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General Relativity, Black Holes, and Cosmology
Andrew J. S. Hamilton
Contents
Preface page 1
Notation 6
PART ONE SPECIAL RELATIVITY 9
Concept Questions 11
What’s important in Special Relativity 13
1 Special Relativity 14
1.1 The postulates of special relativity 14
1.2 The paradox of the constancy of the speed of light 16
1.3 Paradoxes and simultaneity 17
1.4 Time dilation 18
1.5 Lorentz transformation 20
1.6 Paradoxes: Time dilation, Lorentz contraction, and the Twin paradox 22
1.7 The spacetime wheel 23
1.8 Scalar spacetime distance 26
1.9 4-vectors 27
1.10 Energy-momentum 4-vector 28
1.11 Photon energy-momentum 30
1.12 Abstract 4-vectors 31
1.13 What things look like at relativistic speeds 32
1.14 How to programme Lorentz transformations on a computer 33
PART TWO COORDINATE APPROACH TO GENERAL RELATIVITY 35
Concept Questions 37
What’s important? 39
iv Contents
2 Fundamentals of General Relativity 40
2.1 The postulates of General Relativity 40
2.2 Existence of locally inertial frames 41
2.3 Metric 41
2.4 Basis gµ of tangent vectors 42
2.5 4-vectors and tensors 43
2.6 Covariant derivatives 45
2.7 Coordinate 4-velocity 51
2.8 Geodesic equation 51
2.9 Coordinate 4-momentum 52
2.10 Affine parameter 52
2.11 Affine distance 53
2.12 Riemann curvature tensor 53
2.13 Symmetries of the Riemann tensor 54
2.14 Ricci tensor, Ricci scalar 55
2.15 Einstein tensor 55
2.16 Bianchi identities 55
2.17 Covariant conservation of the Einstein tensor 56
2.18 Einstein equations 56
2.19 Summary of the path from metric to the energy-momentum tensor 57
2.20 Energy-momentum tensor of an ideal fluid 57
2.21 Newtonian limit 58
3 ∗More on the coordinate approach 59
3.1 Weyl tensor 59
3.2 Evolution equations for the Weyl tensor 59
3.3 Geodesic deviation 61
3.4 Commutator of the covariant derivative revisited 62
4 ∗Action principle 65
4.1 Principle of least action for point particles 66
4.2 Action for a test particle 67
4.3 Action for a charged test particle in an electromagnetic field 68
4.4 Generalized momentum 69
4.5 Hamiltonian 69
4.6 Derivatives of the action 70
PART THREE IDEAL BLACK HOLES 71
Concept Questions 73
What’s important? 75
Contents v
5 Observational Evidence for Black Holes 76
6 Ideal Black Holes 78
6.1 Definition of a black hole 78
6.2 Ideal black hole 78
6.3 No-hair theorem 79
7 Schwarzschild Black Hole 80
7.1 Schwarzschild metric 80
7.2 Birkhoff’s theorem 81
7.3 Stationary, static 81
7.4 Spherically symmetric 82
7.5 Horizon 83
7.6 Proper time 84
7.7 Redshift 84
7.8 Proper distance 85
7.9 “Schwarzschild singularity” 85
7.10 Embedding diagram 85
7.11 Energy-momentum tensor 86
7.12 Weyl tensor 86
7.13 Gullstrand-Painleve coordinates 86
7.14 Eddington-Finkelstein coordinates 87
7.15 Kruskal-Szekeres coordinates 88
7.16 Penrose diagrams 89
7.17 Schwarzschild white hole, wormhole 90
7.18 Collapse to a black hole 91
7.19 Killing vectors 92
7.20 Time translation symmetry 92
7.21 Spherical symmetry 92
7.22 Killing equation 93
8 Reissner-Nordstrom Black Hole 95
8.1 Reissner-Nordstrom metric 95
8.2 Energy-momentum tensor 96
8.3 Weyl tensor 96
8.4 Horizons 97
8.5 Gullstrand-Painleve metric 97
8.6 Complete Reissner-Nordstrom geometry 98
8.7 Antiverse: Reissner-Nordstrom geometry with negative mass 100
8.8 Ingoing, outgoing 100
8.9 Mass inflation instability 101
8.10 Inevitability of mass inflation 103
vi Contents
8.11 The black hole particle accelerator 104
8.12 The X point 104
8.13 Extremal Reissner-Nordstrom geometry 105
8.14 Reissner-Nordstrom geometry with charge exceeding mass 106
8.15 Reissner-Nordstrom geometry with imaginary charge 106
9 Kerr-Newman Black Hole 109
9.1 Boyer-Lindquist metric 109
9.2 Oblate spheroidal coordinates 110
9.3 Time and rotation symmetries 110
9.4 Ring singularity 111
9.5 Horizons 111
9.6 Angular velocity of the horizon 113
9.7 Ergospheres 113
9.8 Antiverse 114
9.9 Closed timelike curves 114
9.10 Energy-momentum tensor 116
9.11 Weyl tensor 116
9.12 Electromagnetic field 116
9.13 Doran coordinates 117
9.14 Extremal Kerr-Newman geometry 117
9.15 Trajectories of test particles in the Kerr-Newman geometry 118
9.16 Penrose process 122
9.17 Constant latitude trajectories in the Kerr-Newman geometry 123
9.18 Principal null congruence 123
9.19 Circular orbits in the Kerr-Newman geometry 124
PART FOUR HOMOGENEOUS, ISOTROPIC COSMOLOGY 133
Concept Questions 135
What’s important? 137
10 Homogeneous, Isotropic Cosmology 138
10.1 Observational basis 138
10.2 Cosmological Principle 139
10.3 Friedmann-Robertson-Walker metric 140
10.4 Spatial part of the FRW metric: informal approach 140
10.5 Comoving coordinates 142
10.6 Spatial part of the FRW metric: more formal approach 143
10.7 FRW metric 144
10.8 Einstein equations for FRW metric 144
Contents vii
10.9 Newtonian “derivation” of Friedmann equations 145
10.10 Hubble parameter 146
10.11 Critical density 147
10.12 Omega 147
10.13 Redshifting 148
10.14 Types of mass-energy 148
10.15 Evolution of the cosmic scale factor 149
10.16 Conformal time 151
10.17 Looking back along the lightcone 151
10.18 Horizon 152
PART FIVE TETRAD APPROACH TO GENERAL RELATIVITY 157
Concept Questions 159
What’s important? 161
11 The tetrad formalism 162
11.1 Tetrad 162
11.2 Vierbein 162
11.3 The metric encodes the vierbein 163
11.4 Tetrad transformations 164
11.5 Tetrad Tensor 165
11.6 Raising and lowering indices 165
11.7 Gauge transformations 165
11.8 Directed derivatives 166
11.9 Tetrad covariant derivative 166
11.10 Relation between tetrad and coordinate connections 168
11.11 Torsion tensor 168
11.12 No-torsion condition 168
11.13 Antisymmetry of the connection coefficients 169
11.14 Connection coefficients in terms of the vierbein 169
11.15 Riemann curvature tensor 170
11.16 Ricci, Einstein, Bianchi 171
11.17 Electromagnetism 171
12 ∗More on the tetrad formalism 174
12.1 Spinor tetrad formalism 174
12.2 Newman-Penrose tetrad formalism 177
12.3 Electromagnetic field tensor 179
12.4 Weyl tensor 183
12.5 Petrov classification of the Weyl tensor 186
viii Contents
12.6 Raychaudhuri equations and the Sachs optical scalars 187
12.7 Focussing theorem 189
13 ∗The 3+1 (ADM) formalism 191
13.1 ADM tetrad 192
13.2 Traditional ADM approach 193
13.3 Spatial tetrad vectors and tensors 194
13.4 ADM connections, gravity, and extrinsic curvature 194
13.5 ADM Riemann, Ricci, and Einstein tensors 195
13.6 ADM action 196
13.7 ADM equations of motion 199
13.8 Constraints and energy-momentum conservation 200
14 ∗The geometric algebra 201
14.1 Products of vectors 202
14.2 Geometric product 203
14.3 Reverse 204
14.4 The pseudoscalar and the Hodge dual 205
14.5 Reflection 206
14.6 Rotation 207
14.7 A rotor is a spin- 12 object 209
14.8 A multivector rotation is an active rotation 210
14.9 2D rotations and complex numbers 210
14.10 Quaternions 212
14.11 3D rotations and quaternions 213
14.12 Pauli matrices 215
14.13 Pauli spinors 216
14.14 Pauli spinors as scaled 3D rotors, or quaternions 218
14.15 Spacetime algebra 219
14.16 Complex quaternions 221
14.17 Lorentz transformations and complex quaternions 223
14.18 Spatial Inversion (P ) and Time Inversion (T ) 224
14.19 Electromagnetic field bivector 225
14.20 How to implement Lorentz transformations on a computer 225
14.21 Dirac matrices 229
14.22 Dirac spinors 231
14.23 Dirac spinors as complex quaternions 232
14.24 Non-null Dirac spinor — particle and antiparticle 235
14.25 Null Dirac Spinor 236
14.26 Chiral decomposition of a Dirac spinor 237
14.27 Dirac equation 238
Contents ix
14.28 Antiparticles are negative mass particles moving backwards in time 239
14.29 Dirac equation with electromagnetism 240
14.30 CPT 240
14.31 Charge conjugation C 241
14.32 Parity reversal P 242
14.33 Time reversal T 243
14.34 Majorana spinor 243
14.35 Covariant derivatives revisited 243
14.36 General relativistic Dirac equation 244
14.37 3D Vectors as rank-2 spinors 244
PART SIX BLACK HOLE INTERIORS 247
Concept Questions 249
What’s important? 250
15 Black hole waterfalls 251
15.1 Tetrads move through coordinates 251
15.2 Gullstrand-Painleve waterfall 252
15.3 Boyer-Lindquist tetrad 258
15.4 Doran waterfall 259
16 General spherically symmetric spacetime 265
16.1 Spherical spacetime 265
16.2 Spherical electromagnetic field 276
16.3 General relativistic stellar structure 277
16.4 Self-similar spherically symmetric spacetime 278
17 The interiors of spherical black holes 290
17.1 The mechanism of mass inflation 290
17.2 The far future? 293
17.3 Self-similar models of the interior structure of black holes 295
17.4 Instability at outer horizon? 310
PART SEVEN GENERAL RELATIVISTIC PERTURBATION THEORY 311
Concept Questions 313
What’s important? 315
18 Perturbations and gauge transformations 316
18.1 Notation for perturbations 316
18.2 Vierbein perturbation 316
18.3 Gauge transformations 317
x Contents
18.4 Tetrad metric assumed constant 317
18.5 Perturbed coordinate metric 317
18.6 Tetrad gauge transformations 318
18.7 Coordinate gauge transformations 319
18.8 Coordinate gauge transformation of a coordinate scalar 319
18.9 Coordinate gauge transformation of a coordinate vector or tensor 320
18.10 Coordinate gauge transformation of a tetrad vector 320
18.11 Coordinate gauge transformation of the vierbein 321
18.12 Coordinate gauge transformation of the metric 321
18.13 Lie derivative 322
19 Scalar, vector, tensor decomposition 323
19.1 Decomposition of a vector in flat 3D space 323
19.2 Fourier version of the decomposition of a vector in flat 3D space 324
19.3 Decomposition of a tensor in flat 3D space 325
20 Flat space background 326
20.1 Classification of vierbein perturbations 326
20.2 Metric, tetrad connections, and Einstein and Weyl tensors 328
20.3 Spinor components of the Einstein tensor 330
20.4 Too many Einstein equations? 331
20.5 Action at a distance? 332
20.6 Comparison to electromagnetism 333
20.7 Harmonic gauge 337
20.8 What is the gravitational field? 338
20.9 Newtonian (Copernican) gauge 338
20.10 Synchronous gauge 339
20.11 Newtonian potential 341
20.12 Dragging of inertial frames 342
20.13 Quadrupole pressure 343
20.14 Gravitational waves 344
20.15 Energy-momentum carried by gravitational waves 347
PART EIGHT COSMOLOGICAL PERTURBATIONS 349
Concept Questions 351
21 An overview of cosmological perturbations 352
22 ∗Cosmological perturbations in a flat Friedmann-Robertson-Walker background 357
22.1 Unperturbed line-element 357
22.2 Comoving Fourier modes 358
22.3 Classification of vierbein perturbations 358
Contents xi
22.4 Metric, tetrad connections, and Einstein tensor 360
22.5 ADM gauge choices 362
22.6 Conformal Newtonian gauge 362
22.7 Synchronous gauge 363
23 Cosmological perturbations: a simplest set of assumptions 364
23.1 Perturbed FRW line-element 364
23.2 Energy-momenta of ideal fluids 364
23.3 Diffusive damping 367
23.4 Equations for the simplest set of assumptions 368
23.5 Unperturbed background 370
23.6 Generic behaviour of non-baryonic cold dark matter 371
23.7 Generic behaviour of radiation 372
23.8 Regimes 373
23.9 Superhorizon scales 373
23.10 Radiation-dominated, adiabatic initial conditions 376
23.11 Radiation-dominated, isocurvature initial conditions 379
23.12 Subhorizon scales 380
23.13 Matter-dominated 381
23.14 Recombination 382
23.15 Post-recombination 383
23.16 Matter with dark energy 384
23.17 Matter with dark energy and curvature 385
24 ∗Cosmological perturbations: a more careful treatment of photons and baryons 387
24.1 Lorentz-invariant spatial and momentum volume elements 388
24.2 Occupation numbers 388
24.3 Occupation numbers in thermodynamic equilibrium 389
24.4 Boltzmann equation 390
24.5 Non-baryonic cold dark matter 392
24.6 The left hand side of the Boltzmann equation for photons 394
24.7 Spherical harmonics of the photon distribution 395
24.8 Energy-momentum tensor for photons 396
24.9 Collisions 396
24.10 Electron-photon scattering 398
24.11 The photon collision term for electron-photon scattering 399
24.12 Boltzmann equation for photons 402
24.13 Diffusive (Silk) damping 403
24.14 Baryons 403
24.15 Viscous baryon drag damping 405
24.16 Photon-baryon wave equation 405
xii Contents
24.17 Damping of photon-baryon sound waves 407
24.18 Ionization and recombination 409
24.19 Neutrinos 409
24.20 Summary of equations 409
24.21 Legendre polynomials 410
25 Fluctuations in the Cosmic Microwave Background 411
25.1 Primordial power spectrum 411
25.2 Normalization of the power spectrum 412
25.3 CMB power spectrum 412
25.4 Matter power spectrum 413
25.5 Radiative transfer of CMB photons 413
25.6 Integrals over spherical Bessel functions 415
25.7 Large-scale CMB fluctuations 417
25.8 Monopole, dipole, and quadrupole contributions to Cℓ 418
25.9 Integrated Sachs-Wolfe (ISW) effect 418
Preface
Illusory preface
As of writing (May 2010), this book is incomplete. If you happen to discover this draft on the internet, you
are welcome to it (and especially welcome to send me helpful advice and criticism). This book has been
written during two semesters of teaching graduate general relativity at the University of Colorado, Boulder,
during Spring 2008 and 2010. I hope to complete the book the next time I teach the course, which could
possibly be Spring 2012.
Meanwhile, I am vividly aware of the book’s shortcomings. The book is incomplete in many parts, and
needs pruning in others. If the early chapters read more like notes than a book, that is true; I was some
way into writing before I realised that a book was taking shape. Especially, the book is missing many
planned figures. Many of the anticipated figures can be found at three websites: “Special Relativity”
(http://casa.colorado.edu/~ajsh/sr/sr.shtml), “Falling into a Black Hole” (http://casa.colorado.
edu/~ajsh/schw.shtml), and “Inside Black Holes” (http://jila.colorado.edu/~ajsh/insidebh/index.
html).
Although the book is incomplete, I have tried hard to keep mathematical errors from creeping in. If you
find an error — especially in a minus sign or a factor — please let me know.
True preface
This book is driven by one overriding question: “What do students want?”
A fundamental premise of this book is that, in the field of general relativity, the number one thing, by
far, that students want to learn about is black holes. And the second thing they want to learn about is the
cosmic microwave background.
This book is born in part out of frustration with the relentlessly left-brained character of so many texts
on general relativity. I’m probably one of those left-brained characters myself. However, my experience in
general relativistic visualization has convinced me that those interested in black holes from a mathematical
perspective are vastly outnumbered by those fascinated by black holes for other reasons. Among those
2 Preface
reasons are that black holes are, like dinosaurs, awesomely powerful, and supremely mysterious. I worry that
general relativity is (as I hear from students) often taught as if it were little more than tensor calculus. I
fear that the abstract approach repels and culls our right-brained students until only the most left-leaning
of our students remain to trasmit left-brained relativity to the next generation.
Although I was fascinated by general relativity already as a graduate student, my active involvement with
relativity was stimulated by students who insisted that I teach it. From the beginning it seemed obvious
that the way to teach relativity was through visualization. Thus, through teaching, I began to do general
relativistic visualizations of black holes. Initially the visualizations were simple animations, which I put
together into a website “Falling into a Black Hole” in 1997 and 1998. The visualizations touched a chord
with the outside world. In 2001/2 I had the privelege of spending a year’s sabbatical with the Denver
Museum of Nature and Science, where I began developing the Black Hole Flight Simulator (BHFS). That
sabbatical eventually led to a large-format immersive digital dome show “Black Holes: The Other Side of
Infinity,” produced at the DMNS and directed by Tom Lucas. Premiering in 2006, that dome show has
been distributed to some 40 digital domes worldwide. Since that time visualizations with the BHFS have
appeared in several TV documentaries and in a number of exhibits. The experience of working with non-
science professionals has been, and continues to be, intensely enjoyable, and has sensitized me to the insidious
cultural chasms that divide us in an increasingly specialized society.
These experiences have left a prominent dent in my thinking. For example, I think that the highlight of
special relativity is the question of what you see and experience when you pass through a scene at near the
speed of light. Yet most texts scarcely mention the subject, if at all. Similarly, I think that the highlight
of black holes is what (general relativity predicts) you see and experience when you fall into a black hole.
Again, most texts hardly address the issue. Texts often do mention Penrose-Hawking singularity theorems.
Yet few texts mention the all-important mass inflation instability discovered by E. Poisson & W. Israel
(1990). The inflationary instability probably plays the central role in determining the interior structure of
astronomically realistic black holes, and in particular in cutting off the wormhole and white hole connections
to other universes that exist in the ideal Kerr geometry of a rotating black hole. Even E. Poisson (1994) “A
Relativist’s Toolkit: The Mathematics of Black-Hole Mechanics” mentions the inflationary instability only
in the problems at the end of the last chapter. An important goal of this book is to redress this hole in the
teaching of general relativity.
The second focus of this book, after black holes, is the Cosmic Microwave Background (CMB). The CMB
offers a profound window on the genesis of our Universe. Observations of fluctuations in the CMB are in
astonishing agreement with the predictions of general relativistic perturbation theory coupled with some
well-understood physics and some less well-understood but neverthess successful ideas about inflation. For
this book, the goal I set myself was to attempt the simplest possible treatment of CMB fluctuations that
would yield a result that could be compared to observation. This is not an easy goal, since calculation of
CMB fluctuations presents many technical challenges. I applaud recent texts such as M. P. Hobson, G.
P. Efstathiou, & A. N. Lasenby (2006) “General Relativity: An Introduction for Physicists,” and T. Pad-
manabhan (2010) “Gravitation: Foundations and Frontiers,” which include chapters on the CMB power
spectrum. General texts like Hobson et al., Padmanabhan, and the present book by no means replace spe-
cialized books on the Cosmology and the CMB, such as S. Dodelson (2002) “Modern Cosmology,” R. Durrer
Preface 3
(2008) “The Cosmic Microwave Background,” or D. H. Lyth & A. R. Liddle (2009) “The Primordial Density
Perturbation.” However, for many students a course on general relativity may be the only opportunity they
get to learn about the CMB, and I think that a modern course on general relativity should include a basic
introduction to the CMB.
Notwithstanding its intended focus on applications rather than mathematics, this is not an easy book.
It is a serious graduate-level text. R. M. Wald (2006 “Teaching the mathematics of general relativity”,
Am. J. Phys. 74, 471–477 http://arxiv.org/abs/gr-qc/0511073) describes the challenges of teaching the
necessary mathematics in a course on general relativity. This book will not make climbing the mathematical
mountain of general relativity any easier. But the intention is that this book will help you get a clear view
from the top, and not abandon you in fog. While the book does not shirk mathematics, I have tried hard to
make the logic and derivations as clear and tight as possible.
So much for the overall goals of this book. What about the strategy to achieve those goals? Firstly, this
book is intended as a book from which a student can learn, and a lecturer can teach. It is not intended as
a reference book. In a learning/teaching book, one must choose carefully not only what to include, but also
what not to include, because the latter distracts and dilutes.
The grand strategy is to go through general relativity in two passes. In the first pass, the aim is to
run through the foundations of general relativity, and to get to ideal black holes as quickly as possible.
In the second pass, the book essentially starts all over again, using a tetrad-based approach rather than a
coordinate-based approach. The tetrad approach provides the basis for the subsequent treatment of non-ideal
black holes, and of the cosmic microwave background.
The emphasis of (the second half of) this book on tetrads is unusual for a textbook, but consistent with
its right-brained emphasis. If you want to see what’s happening in a spacetime, then you need to look
at it with respect to the frame of an observer, which means working in a locally inertial (orthonormal)
frame. The problem with coordinate frames is that they prescribe that the axes of the spacetime are the
tangent vectors to the coordinates. These tangent vectors are skewed, not orthonormal. Looking at things
in a coordinate frame is like looking at a scene with eyes crossed. Even with geometries as simple as the
Friedmann-Robertson-Walker geometry of homogeneous, isotropic cosmology, it is necessary to play games
to see plainly what its energy-momentum is (for FRW, raise one of the indices on the energy-momentum
tensor — but that trick fails in more complicated spacetimes). Tetrads obviate the need to waste time
attempting to conceptualize the distinction between vectors and covectors (one-forms). My own suspicion
is that the locally flat structure of general relativity may be more fundamental than its globally geometric
character, which could be an emergent phenomenon.
A virtue of the two-pass approach is that the student gets to revisit the fundamentals of general relativity
from two similar but not identical perspectives. This reaffirmation of fundamentals is especially important
given the fast pace and stripped-down coverage.
The course that I teach to senior physics undergraduates and beginning graduate students at the University
of Colorado covers the following 8 topics, each topic taking about 2 weeks during the 16-week semester:
1. Pass 1.
a. Chapter 1: Special relativity.
4 Preface
b. Chapter 2: Coordinate approach to general relativity.
c. Chapters 5–9: Ideal black holes, namely Schwarzschild, Reissner-Nordstrom, and Kerr-Newman.
d. Chapter 10: Homogeneous, isotropic cosmology.
2. Pass 2.
a. Chapter 11: Tetrad approach to general relativity.
b. Chapters 15–17: Black hole interiors.
c. Chapters 18–20: General relativistic perturbation theory.
d. Chapters 21, 23, and 25: Cosmological perturbations.
The first of the eight topics is special relativity. Special relativity is an essential precursor to general
relativity, since a fundamental postulate of general relativity is the Principle of Equivalence, which asserts
that at any point there exist frames, called locally inertial, or free-fall, with respect to which special relativity
operates locally. The strategy of Chapter 1 on Special Relativity is first to confront the paradox of the
constancy of the speed of light, and from there to proceed rapidly to the highlight of special relativity, the
question of what you see and experience when you pass through a scene at near the speed of light. I choose
not to pause to discuss electromagnetism, actions, or other important topics in special relativity, since that
would get in the way of the driving goal, to head to a black hole at the fastest possible pace1.
The second of the eight topics is what I call the coordinate approach to general relativity. This is a lightning
introduction to the fundamental ingredients of general relativity, from the metric through to the Einstein
tensor, using the traditional coordinate-based approach, where components of tensors are expressed relative
to a basis of coordinate tangent vectors. To make the material more accessible, and to lay the groundwork
for tetrads, the book builds on concepts of vectors familiar from high school, and avoids unncecessary
mathematical distractions, such as emphasizing the distinction between vectors and 1-forms. Typically,
texts go through this material at a more leisurely pace, taking time to convey challenging conceptual issues.
Here however I choose not to linger, for two reasons. The first is the obvious one: the goal is to get to black
holes post-haste. The second reason is that, as mentioned earlier in this preface, looking at tensors in a
coordinate basis is like looking at the world with eyes crossed. As my mother used to say, “If you do that
and the wind changes, you’ll be stuck like that forever.”
The third topic is ideal black holes, and here the pace slows. An ideal black hole is one that is stationary
(time translation invariant), and empty outside its singularity, except for the contribution of a static electric
field. In the 4 dimensions of the spacetime we live in, ideal black holes come in just a few varieties: the
Schwarzschild geometry for a spherical, uncharged black hole, the Reissner-Nordstrom for a spherical, charged
black hole, and the Kerr-Newman geometry for a rotating, charged black hole.
The fourth topic is homogeneous, isotropic cosmology, the Friedmann-Robertson-Walker (FRW) geometry.
The FRW geometry forms the essential background spacetime for the cosmological perturbation theory to
be encountered later.
The book now enters the second pass. Tetrads — systems of locally inertial (or other) frames attached to
1 A strategy that might fail for students like myself. I learned special relativity from L. D. Landau & E. M. Lifshitz’sincomparable “The Classical Theory of Fields.” I recall vividly the extraordinary delight in discovering a text that, incontrast to those dreadful books that conveyed the idea that electromagnetism was something to do with resistors andcapacitors, put relativity, Maxwell’s equations, and actions up front.
Preface 5
each point of spacetime — are well known to, and widely used by, general relativists. The tetrad approach
to general relatiivity is more complicated than the coordinate approach in that it requires an additional
superstructure. However, the advantage of being able to see straight, because you are working in an or-
thonormal frame, outweighs the disadvantage of the additional overhead. While the coordinate approach is
adequate for simple spacetimes — ideal black holes, and the FRW geometry — its defects are a barrier to
understanding more complicated spacetimes. Most texts do not cover tetrads, or cover them as an aside. In
this book, tetrads are developed systematically, in one self-contained chapter.
The sixth topic is black hole interiors. Cool, but needs more work.
The seventh topic is general relativistic perturbation theory, some understanding of which is prerequisite
for dealing with cosmological perturbations and the CMB. The approach starts — in a thankfully short
chapter — with one of the most difficult aspects of general relativistic perturbation theory, namely the
problem of coordinate and tetrad gauge ambiguities. This might seem a peculiar starting point. A more
typical starting point is to vary the metric, pick a gauge, and lo there are waves. However, I think that it
is important to show how, at least in flat or FRW background spacetimes, the real physical perturbations
emerge naturally from the formalism, without having to pick a gauge. In flat spacetime, the formalism
picks out one particular gauge, the Newtonian gauge (though I think it should be called the Copernican
gauge, because in the solar system it would pick out a Sun-centered almost-Cartesian frame), in which the
perturbations retained are precisely the physical perturbations and no others. The Newtonian/Copernican
gauge provides the natural arena for elucidating general relativistic phenomena such as the dragging of
inertial frames, and gravitational waves.
The eighth and final topic is cosmological perturbation theory, emphasizing the calculation of fluctuations
in the CMB. This was one of the most challenging parts of the book to write, because of the shear volume of
physics that goes into the calculation. I did my best to condense the core of the calculation into “a simplest
set of assumptions,” Chapter 23. However, if you want to calculate a CMB power spectrum that you can
actually compare to observations, then you’ll have to go beyond the “simple” chapter. Subsequent chapters
will help you do that.
Beyond the eight topics described above, there are several chapters, all starred2, that contain relevant,
but lower priority, material. The material contains some fun stuff, my favourite being “How to implement
Lorentz transformations on a computer,” §14.20.
2 The idea of starred chapters comes from Steven L. Weinberg’s classic 1972 text “Gravitation and Cosmology,” from which Ilearned general relativity. Curiously, I found his starred chapters often more interesting than the unstarred ones.
Notation
Except where actual units are needed, units are such that the speed of light is one, c = 1, and Newton’s
gravitational constant is one, G = 1.
The metric signature is −+++.
Greek (brown) letters α, β, ..., denote dummy 4D coordinate indices. Latin (black) letters a, b, ..., denote
dummy 4D tetrad indices. Mid-alphabet Latin letters i, j, ... denote 3D indices, either coordinate (brown)
or tetrad (black). To avoid distraction, colouring is applied only to coordinate indices, not to the coordinates
themselves.
Specific (non-dummy) components of a vector are labelled by the corresponding coordinate (brown) or
tetrad (black) direction, for example Aµ = At, Ax, Ay, Az or Am = At, Ax, Ay, Az. Allowing the same
label to denote either a coordinate or a tetrad index risks ambiguity, but it should apparent from the context
what is meant. Some texts distinguish coordinate and tetrad indices for example by a caret on the latter,
but this produces notational overload.
Boldface denotes abstract vectors, in either 3D or 4D. In 4D, A = Aµgµ = Amγγm, where gµ denote
coordinate tangent axes, and γγm denote tetrad axes.
Repeated paired dummy indices are summed over, the implicit summation convention. In special and
general relativity, one index of a pair must be up (contravariant), while the other must be down (covariant).
If the space being considered is Euclidean, then both indices may be down.
∂/∂xµ denotes coordinate partial derivatives, which commute. ∂m denotes tetrad directed derivatives,
which do not commute. Dµ and Dm denote respectively coordinate-frame and tetrad-frame covariant deriva-
tives.
Choice of metric signature
There is a tendency, by no means unanimous, for general relativists to prefer the −+++ metric signature,
while particle physicists prefer +−−−.
For someone like me who does general relativistic visualization, there is no contest: the choice has to be
−+++, so that signs remain consistent between 3D spatial vectors and 4D spacetime vectors. For example,
Notation 7
the 3D industry knows well that quaternions provide the most efficient and powerful way to implement
spatial rotations. As shown in Chapter 14, complex quaternions provide the best way to implement Lorentz
transformations, with the subgroup of real quaternions continuing to provide spatial rotations. Compatibility
requires −+++. Actually, OpenGL and other graphics languages put spatial coordinates in the first three
indices, leaving time to occupy the fourth index; but in these notes I stick to the physics convention of
putting time in the zeroth index.
In practical calculations it is convenient to be able to switch transparently between boldface and index
notation in both 3D and 4D contexts. This is where the +−−− signature poses greater potential for
misinterpretation in 3D. For example, with this signature, what is the sign of the 3D scalar product
a · b ? (0.1)
Is it a · b =∑3
i=1 aibi or a · b =
∑3i=1 a
ibi? To be consistent with common 3D usage, it must be the
latter. With the +−−− signature, it must be that a · b = −aibi, where the repeated indices signify implicit
summation over spatial indices. So you have to remember to introduce a minus sign in switching between
boldface and index notation.
As another example, what is the sign of the 3D vector product
a× b ? (0.2)
Is it a×b =∑3
jk=1 εijkajbk or a×b =
∑3jk=1 ε
ijka
jbk or a×b =∑3
jk=1 εijkajbk? Well, if you want to switch
transparently between boldface and index notation, and you decide that you want boldface consistently to
signify a vector with a raised index, then maybe you’d choose the middle option. To be consistent with
standard 3D convention for the sign of the vector product, maybe you’d choose εijk to have positive sign for
ijk an even permutation of xyz.
Finally, what is the sign of the 3D spatial gradient operator
∇ ≡ ∂
∂x? (0.3)
Is it ∇ = ∂/∂xi or ∇ = ∂/∂xi? Convention dictates the former, in which case it must be that some boldface
3D vectors must signify a vector with a raised index, and others a vector with a lowered index. Oh dear.
PART ONE
SPECIAL RELATIVITY
Concept Questions
1. What does c = universal constant mean? What is speed? What is distance? What is time?
2. c+ c = c. How can that be possible?
3. The first postulate of special relativity asserts that spacetime forms a 4-dimensional continuum. The fourth
postulate of special relativity asserts that spacetime has no absolute existence. Isn’t that a contradiction?
4. The principle of special relativity says that there is no absolute spacetime, no absolute frame of reference
with respect to which position and velocity are defined. Yet does not the cosmic microwave background
define such a frame of reference?
5. How can two people moving relative to each other at near c both think each other’s clock runs slow?
6. How can two people moving relative to each other at near c both think the other is Lorentz-contracted?
7. All paradoxes in special relativity have the same solution. In one word, what is that solution?
8. All conceptual paradoxes in special relativity can be understood by drawing what kind of diagram?
9. Your twin takes a trip to α Cen at near c, then returns to Earth at near c. Meeting your twin, you see
that the twin has aged less than you. But from your twin’s perspective, it was you that receded at near
c, then returned at near c, so your twin thinks you aged less. Is it true?
10. Blobs in the jet of the galaxy M87 have been tracked by the Hubble Space Telescope to be moving at
about 6c. Does this violate special relativity?
11. If you watch an object move at near c, does it actually appear Lorentz-contracted? Explain.
12. You speed towards the center of our Galaxy, the Milky Way, at near c. Does the center appear to you
closer or farther away?
13. You go on a trip to the center of the Milky Way, 30,000 lightyears distant, at near c. How long does the
trip take you?
14. You surf a light ray from a distant quasar to Earth. How much time does the trip take, from your
perspective?
15. If light is a wave, what is waving?
16. As you surf the light ray, how fast does it appear to vibrate?
17. How does the phase of a light ray vary along the light ray? Draw surfaces of constant phase on a spacetime
diagram.
12 Concept Questions
18. You see a distant galaxy at a redshift of z = 1. If you could see a clock on the galaxy, how fast would the
clock appear to tick? Could this be tested observationally?
19. You take a trip to α Cen at near c, then instantaneously accelerate to return at near c. If you are looking
through a telescope at a clock on the Earth while you instantaneously accelerate, what do you see happen
to the clock?
20. In what sense is time an imaginary spatial dimension?
21. In what sense is a Lorentz boost a rotation by an imaginary angle?
22. You know what it means for an object to be rotating at constant angular velocity. What does it mean for
an object to be boosting at a constant rate?
23. A wheel is spinning so that its rim is moving at near c. The rim is Lorentz-contracted, but the spokes are
not. How can that be?
24. You watch a wheel rotate at near the speed of light. The spokes appear bent. How can that be?
25. Does a sunbeam appear straight or bent when you pass by it at near the speed of light?
26. Energy and momentum are unified in special relativity. Explain.
27. In what sense is mass equivalent to energy in special relativity? In what sense is mass different from
energy?
28. Why is the Minkowski metric unchanged by a Lorentz transformation?
29. What is the best way to program Lorentz transformations on a computer?
What’s important in Special Relativity
See
http://casa.colorado.edu/~ajsh/sr/
1. Postulates of special relativity.
2. Understanding conceptually the unification of space and time implied by special relativity.
a. Spacetime diagrams.
b. Simultaneity.
c. Understanding the paradoxes of relativity — time dilation, Lorentz contraction, the twin paradox.
3. The mathematics of spacetime transformations
a. Lorentz transformations.
b. Invariant spacetime distance.
c. Minkowski metric.
d. 4-vectors.
e. Energy-momentum 4-vector. E = mc2.
f. The energy-momentum 4-vector of massless particles, such as photons.
4. What things look like at relativistic speeds.
1
Special Relativity
1.1 The postulates of special relativity
The theory of special relativity can be derived formally from a small number of postulates:
1. Space and time form a 4-dimensional continuum:
2. The existence of locally inertial frames;
3. The speed of light is constant;
4. The principle of special relativity.
The first two postulates are assertions about the structure of spacetime, while the last two postulates form
the heart of special relativity. Most books mention just the last two postulates, but I think it is important
to know that special (and general) relativity simply postulate the 4-dimensional character of spacetime, and
that special relativity postulates moreover that spacetime is flat.
1. Space and time form a 4-dimensional continuum. The correct mathematical word for continuum
is manifold. A 4-dimensional manifold is defined mathematically to be a topological space that is locally
homeomorphic to Euclidean 4-space R4.
The postulate that spacetime forms a 4-dimensional continuum is a generalization of the classical Galilean
concept that space and time form separate 3 and 1 dimensional continua. The postulate of a 4-dimensional
spacetime continuum is retained in general relativity.
Physicists widely believe that this postulate must ultimately break down, that space and time are quantized
over extremely small intervals of space and time, the Planck length√
G~/c3 ≈ 10−35 m, and the Planck time√
G~/c5 ≈ 10−43 s, where G is Newton’s gravitational constant, ~ ≡ h/(2π) is Planck’s constant divided by
2π, and c is the speed of light.
2. The existence of globally inertial frames. Statement: “There exist global spacetime frames with
respect to which unaccelerated objects move in straight lines at constant velocity.”
A spacetime frame is a system of coordinates for labelling space and time. Four coordinates are needed,
because spacetime is 4-dimensional. A frame in which unaccelerated objects move in straight lines at constant
1.1 The postulates of special relativity 15
velocity is called an inertial frame. One can easily think of non-inertial frames: a rotating frame, an
accelerating frame, or simply a frame with some bizarre Dahlian labelling of coordinates.
A globally inertial frame is an inertial frame that covers all of space and time. The postulate that
globally inertial frames exist is carried over from classical mechanics (Newton’s first law of motion).
Notice the subtle shift from the Newtonian perspective. The postulate is not that particles move in straight
lines, but rather that there exist spacetime frames with respect to which particles move in straight lines.
Implicit in the assumption of the existence of globally inertial frames is the assumption that the geometry
of spacetime is flat, the geometry of Euclid, where parallel lines remain parallel to infinity. In general
relativity, this postulate is replaced by the weaker postulate that local (not global) inertial frames exist. A
locally inertial frame is one which is inertial in a “small neighbourhood” of a spacetime point. In general
relativity, spacetime can be curved.
3. The speed of light is constant. Statement: “The speed of light c is a universal constant, the same in
any inertial frame.”
This postulate is the nub of special relativity. The immediate challenge of this chapter, §1.2, is to confront
its paradoxical implications, and to resolve them.
Measuring speed requires being able to measure intervals of both space and time: speed is distance travelled
divided by time elapsed. Inertial frames constitute a special class of spacetime coordinate systems; it is with
respect to distance and time intervals in these special frames that the speed of light is asserted to be constant.
In general relativity, arbitrarily weird coordinate systems are allowed, and light need move neither in
straight lines nor at constant velocity with respect to bizarre coordinates (why should it, if the labelling
of space and time is totally arbitrary?). However, general relativity asserts the existence of locally inertial
frames, and the speed of light is a universal constant in those frames.
In 1983, the General Conference on Weights and Measures officially defined the speed of light to be
c ≡ 299,792,458 m s−1, (1.1)
and the meter, instead of being a primary measure, became a secondary quantity, defined in terms of the
second and the speed of light.
4. The principle of special relativity. Statement: “The laws of physics are the same in any inertial
frame, regardless of position or velocity.”
Physically, this means that there is no absolute spacetime, no absolute frame of reference with respect to
which position and velocity are defined. Only relative positions and velocities between objects are meaningful.
It is to be noted that the principle of special relativity does not imply the constancy of the speed of light,
although the postulates are consistent with each other. Moreover the constancy of the speed of light does
not imply the Principle of Special Relativity, although for Einstein the former appears to have been the
inspiration for the latter.
An example of the application of the principle of special relativity is the construction of the energy-
momentum 4-vector of a particle, which should have the same form in any inertial frame (§1.10).
16 Special Relativity
1.2 The paradox of the constancy of the speed of light
The postulate that the speed of light is the same in any inertial frame leads immediately to a paradox.
Resolution of this paradox compels a revolution in which space and time are united from separate 3 and
1-dimensional continua into a single 4-dimensional continuum.
Here, Figure ??, is Vermilion. She emits a flash of light. Vermilion thinks that the light moves outward
at the same speed in all directions. So Vermilion thinks that she is at the centre of the expanding sphere of
light.
But here also, Figure ??, is Cerulean, moving away from Vermilion, at about 12 the speed of light. Vermilion
thinks that she is at the centre of the expanding sphere of light, as before. But, says special relativity,
Cerulean also thinks that the light moves outward at the same speed in all directions from him. So Cerulean
should be at the centre of the expanding light sphere too. But he’s not, is he. Paradox!
Concept question 1.1 Would the light have expanded differently if Cerulean had emitted the light?
1.2.1 Challenge
Can you figure out Einstein’s solution to the paradox? Somehow you have to arrange that both Vermilion
and Cerulean regard themselves as being in the centre of the expanding sphere of light.
1.2.2 Spacetime diagram
A spacetime diagram suggests a way of thinking which leads to the solution of the paradox of the constancy
of the speed of light. Indeed, spacetime diagrams provide the way to resolve all conceptual paradoxes in
special relativity, so it is thoroughly worthwhile to understand them.
A spacetime diagram, Figure ??, is a diagram in which the vertical axis represents time, while the
horizontal axis represents space. Really there are three dimensions of space, which can be thought of as
filling additional horizontal dimensions. But for simplicity a spacetime diagram usually shows just one
spatial dimension.
In a spacetime diagram, the units of space and time are chosen so that light goes one unit of distance
in one unit of time, i.e. the units are such that the speed of light is one, c = 1. Thus light always moves
upward at 45 from vertical in a spacetime diagram. Each point in 4-dimensional spacetime is called an
event. Light signals converging to or expanding from an event follow a 3-dimensional hypersurface called
the lightcone. Light converging on to an event in on the past lightcone, while light emerging from an
event is on the future lightcone.
Here is a spacetime diagram of Vermilion emitting a flash of light, and Cerulean moving relative to
Vermilion at about 12 the speed of light. This is a spacetime diagram version of the situation illustrated
in Figure ??. The lines along which Vermilion and Cerulean move through spacetime are called their
worldlines.
Consider again the challenge problem. The problem is to arrange that both Vermilion and Cerulean are
at the centre of the lightcone, from their own points of view.
1.3 Paradoxes and simultaneity 17
Here’s a clue. Cerulean’s concept of space and time may not be the same as Vermilion’s.
1.2.3 Centre of the lightcone
Einstein’s solution to the paradox is that Cerulean’s spacetime is skewed compared to Vermilion’s, as illus-
trated by Figure ??. The thing to notice in the diagram is that Cerulean is in the centre of the lightcone,
according to the way Cerulean perceives space and time. Vermilion remains at the centre of the lightcone
according to the way Vermilion perceives space and time. In the diagram Vermilion and her space are drawn
at one “tick” of her clock past the point of emission, and likewise Cerulean and his space are drawn at one
“tick” of his identical clock past the point of emission. Of course, from Cerulean’s point of view his spacetime
is quite normal, and it’s Vermilion’s spacetime that is skewed.
In special relativity, the transformation between the spacetime frames of two inertial observers is called a
Lorentz transformation. In general, a Lorentz transformation consists of a spatial rotation about some
spatial axis, combined with a Lorentz boost by some velocity in some direction.
Only space along the direction of motion gets skewed with time. Distances perpendicular to the direction
of motion remain unchanged. Why must this be so? Consider two hoops which have the same size when at
rest relative to each other. Now set the hoops moving towards each other. Which hoop passes inside the
other? Neither! For suppose Vermilion thinks Cerulean’s hoop passed inside hers; by symmetry, Cerulean
must think Vermilion’s hoop passed inside his; but both cannot be true; the only possibility is that the hoops
remain the same size in directions perpendicular to the direction of motion.
Cottoned on? Then you have understood the crux of special relativity, and you can now go away and
figure out all the mathematics of Lorentz transformations. Just like Einstein. The mathematical problem is:
what is the relation between the spacetime coordinates t, x, y, z and t′, x′, y′, z′ of a spacetime interval,
a 4-vector, in Vermilion’s versus Cerulean’s frames, if Cerulean is moving relative to Vermilion at velocity v
in, say, the x direction? The solution follows from requiring
1. that both observers consider themselves to be at the centre of the lightcone, and
2. that distances perpendicular to the direction of motion remain unchanged,
as illustrated by Figure ??. (An alternative version of the second condition is that a Lorentz transformation
at velocity v followed by a Lorentz transformation at velocity −v should yield the unit transformation.)
Note that the postulate of the existence of globally inertial frames implies that Lorentz transformations
are linear, that straight lines (4-vectors) in one inertial spacetime frame transform into straight lines in other
inertial frames.
You will solve this problem in the next section but two, §1.5. As a prelude, the next two sections, §§1.3
and 1.4 discuss simultaneity and time dilation.
1.3 Paradoxes and simultaneity
Most (all?) of the apparent paradoxes of special relativity arise because observers moving at different veloc-
ities relative to each other have different notions of simultaneity.
18 Special Relativity
1.3.1 Operational definition of simultaneity
How can simultaneity, the notion of events ocurring at the same time at different places, be defined opera-
tionally?
One way is illustrated in Figure ??. Vermilion surrounds herself with a set of mirrors, equidistant from
Vermilion. She sends out a flash of light, which reflects off the mirrors back to Vermilion. How does Vermilion
know that the mirrors are all the same distance from her? Because the relected flash returns from the mirrors
to Vermilion all at the same instant.
Vermilion asserts that the light flash must have hit all the mirrors simultaneously. Vermilion also asserts
that the instant when the light hit the mirrors must have been the instant, as registered by her wristwatch,
precisely half way between the moment she emitted the flash and the moment she received it back again. If it
takes, say, 2 seconds between flash and receipt, then Vermilion concludes that the mirrors are 1 lightsecond
away from her.
Figure ?? shows a spacetime diagram of Vermilion’s mirror experiment above. According to Vermilion, the
light hits the mirrors everywhere at the same instant, and the spatial hyperplane passing through these events
is a hypersurface of simultaneity. More generally, from Vermilion’s perspective, each horizontal hyperplane
in the spacetime diagram is a hypersurface of simultaneity.
Cerulean defines surfaces of simultaneity using the same operational setup: he encompasses himself with
mirrors, arranging them so that a flash of light returns from them to him all at the same instant. But
whereas Cerulean concludes that his mirrors are all equidistant from him and that the light bounces off
them all at the same instant, Vermilion thinks otherwise. From Vermilion’s point of view, the light bounces
off Cerulean’s mirrors at different times and moreover at different distances from Cerulean. Only so can the
speed of light be constant, as Vermilion sees it, and yet the light return to Cerulean all at the same instant.
Of course from Cerulean’s point of view all is fine: he thinks his mirrors are equidistant from him, and
that the light bounces off them all at the same instant. The inevitable conclusion is that Cerulean must
measure space and time along axes that are skewed relative to Vermilion’s. Events that happen at the same
time according to Cerulean happen at different times according to Vermilion; and vice versa. Cerulean’s
hypersurfaces of simultaneity are not the same as Vermilion’s.
From Cerulean’s point of view, Cerulean remains always at the centre of the lightcone. Thus for Cerulean,
as for Vermilion, the speed of light is constant, the same in all directions.
1.4 Time dilation
Vermilion and Cerulean construct identical clocks, consisting of a light beam which bounces off a mirror.
Tick, the light beam hits the mirror, tock, the beam returns to its owner. As long as Vermilion and Cerulean
remain at rest relative to each other, both agree that each other’s clock tick-tocks at the same rate as their
own.
But now suppose Cerulean goes off at velocity v relative to Vermilion, in a direction perpendicular to the
direction of the mirror. A far as Cerulean is concerned, his clock tick-tocks at the same rate as before, a tick
1.4 Time dilation 19
at the mirror, a tock on return. But from Vermilion’s point of view, although the distance between Cerulean
and his mirror at any instant remains the same as before, the light has further to go. And since the speed
of light is constant, Vermilion thinks it takes longer for Cerulean’s clock to tick-tock than her own. Thus
Vermilion thinks Cerulean’s clock runs slow relative to her own.
1.4.1 Lorentz gamma factor
How much slower does Cerulean’s clock run, from Vermilion’s point of view? In special relativity the factor
is called the Lorentz gamma factor γ, introduced by the Dutch physicist Hendrik A. Lorentz in 1904, one
year before Einstein proposed his theory of special relativity. Let us see how the Lorentz gamma factor is
related to Cerulean’s velocity v.
In units where the speed of light is one, c = 1, Vermilion’s mirror is one tick away from her, and from her
point of view the vertical distance between Cerulean and his mirror is the same, one tick. But Vermilion
thinks that the distance travelled by the light beam between Cerulean and his mirror is γ ticks. Cerulean is
moving at speed v, so Vermilion thinks he moves a distance of γv ticks during the γ ticks of time taken by
the light to travel from Cerulean to his mirror. Thus, from Vermilion’s point of view, the vertical line from
Cerulean to his mirror, Cerulean’s light beam, and Cerulean’s path form a triangle with sides 1, γ, and γv,
as illustrated. Pythogoras’ theorem implies that
12 + (γv)2 = γ2 . (1.2)
From this it follows that the Lorentz gamma factor γ is related to Cerulean’s velocity v by
γ =1√
1− v2, (1.3)
which is Lorentz’s famous formula.
1.4.2 Paradox
Vermilion thinks Cerulean’s clock runs slow, by the Lorentz factor γ. But of course from Cerulean’s perspec-
tive it is Vermilion who is moving, and Vermilion whose clock runs slow. How can both think the other’s
clock runs slow? Paradox!
The resolution of the paradox, as usual in special relativity, involves simultaneity, and as usual it helps to
draw a spacetime diagram, such as the one in Figure ??.
While Vermilion thinks events happen simultaneously along horizontal planes in this diagram, Cerulean
thinks events occur simultaneously along skewed planes. Thus Vermilion thinks her clock ticks when Cerulean
is at the point before Cerulean’s clock ticks. Conversely, Cerulean thinks his clock ticks when Vermilion is
at the point before Vermilion’s clock ticks.
Where do the two light beams in Vermilion’s and Cerulean’s clocks go in this spacetime diagram? Figure ??
shows a 3D spacetime diagram.
20 Special Relativity
Concept question 1.2 Figure ?? shows a picture of a 3D cube. Is one edge shorter than the other?
Projected on to the page, it appears so, but in reality all the edges have equal length. In what ways is this
situation similar or disimilar to time dilation in 4D relativity?
1.5 Lorentz transformation
A Lorentz transformation is a rotation of space and time. Lorentz transformations form a 6-dimensional
group, with 3 dimensions from spatial rotations, and 3 dimensions from Lorentz boosts.
If you wish to understand special relativity mathematically, then it is essential for you to go through the
exercise of deriving the form of Lorentz transformations for yourself. Indeed, this problem is the challenge
problem posed in §1.2, recast as a mathematical exercise. For simplicity, it is enough to consider the case of
a Lorentz boost by velocity v along the x-axis.
You can derive the form of a Lorentz transformation either pictorially (geometrically), or algebraically.
Ideally you should do both.
Exercise 1.3 Pictorial derivation of the Lorentz transformation. Construct, with ruler and com-
pass, a spacetime diagram that looks like the one in Figure ??. You should recognize that the square
represents the paths of lightrays that Vermilion uses to define a hypersurface of simultaneity, while the
rectangle represents the same thing for Cerulean. Notice that Cerulean’s worldline and line of simultaneity
are diagonals along his light rectangle, so the angles between those lines and the lightcone are equal. Notice
also that the areas of the square and the rectangle are the same, which expresses the fact that the area is
multiplied by the determinant of the Lorentz transformation matrix, which must be one (why?). Use your
geometric construction to derive the mathematical form of the Lorentz transformation.
Exercise 1.4 3D model of the Lorentz transformation. Make a 3D spacetime diagram of the Lorentz
transformation, with not only an x-dimension, as in the previous problem, but also a y-dimension. Resist the
temptation to use a 3D computer modelling program. Believe me, you will learn much more from hands-on
model-making. Make the lightcone from flexible paperboard, the spatial hypersurface of simultaneity from
stiff paperboard, and the worldline from wooden dowel.
Exercise 1.5 Mathematical derivation of the Lorentz transformation. Relative to person A (Ver-
milion, unprimed frame), person B (Cerulean, primed frame) moves at velocity v along the x-axis. Derive
the form of the Lorentz transformation between the coordinates (t, x, y, z) of a 4-vector in A’s frame and the
corresponding coordinates (t′, x′, y′, z′) in B’s frame from the assumptions:
1. that the transformation is linear;
2. that the spatial coordinates in the directions orthogonal to the direction of motion are unchanged;
3. that the speed of light c is the same for both A and B, so that x = t in A’s frame transforms to x′ = t′ in
B’s frame, and likewise x = −t in A’s frame transforms to x′ = −t′ in B’s frame;
4. the definition of speed; if B is moving at speed v relative to A, then x = vt in A’s frame transforms to
x′ = 0 in B’s frame;
1.5 Lorentz transformation 21
5. spatial isotropy; specifically, show that if A thinks B is moving at velocity v, then B must think that A
is moving at velocity −v, and symmetry (spatial isotropy) between these two situations then fixes the
Lorentz γ factor.
Your logic should be precise, and explained in clear, concise English.
You should find that the Lorentz transformation for a Lorentz boost by velocity v along the x-axis is
t′ = γt− γvxx′ = − γvt+ γx
y′ = y
z′ = z
,
t = γt′ + γvx′
x = γvt′ + γx′
y = y′
z = z′
. (1.4)
The transformation can be written elegantly in matrix notation:
t′
x′
y′
z′
=
γ −γv 0 0
−γv γ 0 0
0 0 1 0
0 0 0 1
t
x
y
z
, (1.5)
with inverse
t
x
y
z
=
γ γv 0 0
γv γ 0 0
0 0 1 0
0 0 0 1
t′
x′
y′
z′
. (1.6)
A Lorentz transformation at velocity v followed by a Lorentz transformation at velocity v in the opposite
direction, i.e. at velocity −v, yields the unit transformation, as it should:
γ γv 0 0
γv γ 0 0
0 0 1 0
0 0 0 1
γ −γv 0 0
−γv γ 0 0
0 0 1 0
0 0 0 1
=
1 0 0 0
0 1 0 0
0 0 1 0
0 0 0 1
. (1.7)
The determinant of the Lorentz transformation is one, as it should be:
∣
∣
∣
∣
∣
∣
∣
∣
γ −γv 0 0
−γv γ 0 0
0 0 1 0
0 0 0 1
∣
∣
∣
∣
∣
∣
∣
∣
= γ2(1− v2) = 1 . (1.8)
Indeed, requiring that the determinant be one provides another derivation of the formula (1.3) for the Lorentz
gamma factor.
Concept question 1.6 Why must the determinant of a Lorentz transformation be one?
22 Special Relativity
1.6 Paradoxes: Time dilation, Lorentz contraction, and the Twin paradox
There are several classic paradoxes in special relativity. Two of them have already been met above: the
paradox of the constancy of the speed of light in §1.2, and the paradox of time dilation in §1.4. This section
collects three famous paradoxes: time dilation (reiterating §1.4), Lorentz contraction, and the Twin paradox.
If you wish to understand special relativity conceptually, then you should work through all these paradoxes
yourself. As remarked in §1.3, most (all?) paradoxes in special relativity arise because different observers
have different notions of simultaneity, and most (all?) paradoxes can be solved using spacetime diagrams.
The Twin paradox is particularly helpful because it illustrates several different facets of special relativity,
not only time dilation, but also how light travel time modifies what an observer actually sees.
1.6.1 Time dilation
If a timelike interval t, r corresponds to motion at velocity v, then r = vt. The proper time along the
interval is
τ =√
t2 − r2 = t√
1− v2 =t
γ. (1.9)
This is Lorentz time dilation: the proper time interval τ experienced by a moving person is a factor γ less
than the time interval t according to an onlooker.
Exercise 1.7 On a spacetime diagram, show how two observers moving relative to each other can both
consider the other’s clock to run slow compared to their own.
1.6.2 Fitzgerald-Lorentz contraction
Consider a rocket of proper length l, so that in the rocket’s own rest frame (primed) the back and front ends
of the rocket move through time t′ with coordinates
t′, x′ = t′, 0 and t′, l . (1.10)
From the perspective of an observer who sees the rocket move at velocity v in the x-direction, the worldlines
of the back and front ends of the rocket are at
t, x = γt′, γvt′ and γt′ + γvl, γvt′ + γl . (1.11)
However, the observer measures the length of the rocket simultaneously in their own frame, not the rocket
frame. Solving for γt′ = t at the back and γt′ + γvl = t at the front gives
t, x = t, vt and
t, vt+l
γ
(1.12)
which says that the observer measures the front end of the rocket to be a distance l/γ ahead of the back
end. This is Lorentz contraction: an object of proper length l is measured by a moving person to be shorter
by a factor γ.
1.7 The spacetime wheel 23
Exercise 1.8 On a spacetime diagram, show how two observers moving relative to each other can both
consider the other to be contracted along the direction of motion.
1.6.3 Twin paradox
See Exercise 1.11 at the end of the chapter.
1.7 The spacetime wheel
1.7.1 3D wheel
Figure ?? shows an ordinary 3D wheel. As the wheel rotates, a point on the wheel describes an invariant
circle. The coordinates x, y of a point on the wheel relative to its centre change, but the distance r between
the point and the centre remains constant
r2 = x2 + y2 = constant . (1.13)
More generally, the coordinates x, y, z of the interval between any two points in 3-dimensional space (a
vector) change when the coordinate system is rotated in 3 dimensions, but the separation r of the two points
remains constant
r2 = x2 + y2 + z2 = constant . (1.14)
1.7.2 4D spacetime wheel
Figure ?? shows a 4D spacetime wheel. The diagram here is a spacetime diagram, with time t vertical and
space x horizontal. A rotation between time t and space x is a Lorentz boost in the x-direction. As the
spacetime wheel boosts, a point on the wheel describes an invariant hyperbola. The spacetime coordinates
t, x of a point on the wheel relative to its centre change, but the spacetime separation s between the point
and the centre remains constant
s2 = − t2 + x2 = constant . (1.15)
More generally, the coordinates t, x, y, z of the interval between any two events in 4-dimensional spacetime
(a 4-vector) change when the coordinate system is boosted or rotated, but the spacetime separation s of the
two events remains constant
s2 = − t2 + x2 + y2 + z2 = constant . (1.16)
1.7.3 Lorentz boost as a rotation by an imaginary angle
The − sign instead of a + sign in front of the t2 in the spacetime separation formula (1.16) means that time
t can often be treated mathematically as if it were an imaginary spatial dimension. That is, t = iw where
i ≡√−1 and w is a “fourth spatial coordinate”.
24 Special Relativity
A Lorentz boost by a velocity v can likewise be treated as a rotation by an imaginary angle. Consider a
normal spatial rotation in which a primed frame is rotated in the wx-plane clockwise by an angle a about
the origin, relative to the unprimed frame. The relation between the coordinates w′, x′ and w, x of a
point in the two frames is(
w′
x′
)
=
(
cos a − sina
sin a cos a
)(
w
x
)
. (1.17)
Now set t = iw and α = ia with t and α both real. In other words, take the spatial coordinate w to be
imaginary, and the rotation angle a likewise to be imaginary. Then the rotation formula above becomes(
t′
x′
)
=
(
coshα − sinhα
− sinhα coshα
)(
t
x
)
(1.18)
This agrees with the usual Lorentz transformation formula (??) if the boost velocity v and boost angle α
are related by
v = tanhα , (1.19)
so that
γ = coshα , γv = sinhα . (1.20)
This provides a convenient way to add velocities in special relativity: the boost angles simply add (for boosts
along the same direction), just as spatial rotation angles add (for rotations about the same axis). Thus a
boost by velocity v1 = tanhα1 followed by a boost by velocity v2 = tanhα2 in the same direction gives a
net velocity boost of v = tanhα where
α = α1 + α2 . (1.21)
The equivalent formula for the velocities themselves is
v =v1 + v21 + v1v2
, (1.22)
the special relativistic velocity addition formula.
1.7.4 Trip across the Universe at constant acceleration
Suppose that you took a trip across the Universe in a spaceship, accelerating all the time at one Earth
gravity g. How far would you travel in how much time?
The spacetime wheel offers a cute way to solve this problem, since the rotating spacetime wheel can
be regarded as representing spacetime frames undergoing constant acceleration. Specifically, points on the
right quadrant of the rotating spacetime wheel represent worldlines of persons who accelerate with constant
acceleration in their own frame.
If the units of space and time are chosen so that the speed of light and the gravitational acceleration are
both one, c = g = 1, then the proper time experienced by the accelerating person is the boost angle α, and
1.7 The spacetime wheel 25
Table 1.1 Trip across the Universe.
Time elapsed Time elapsedon spaceship on Earth Distance travelled To
in years in years in lightyears
α sinhα coshα− 1
0 0 0 Earth (starting point)1 1.175 .54312 3.627 2.762
2.337 5.127 4.223 Proxima Cen3.962 26.3 25.3 Vega6.60 368 367 Pleiades10.9 2.7 × 104 2.7 × 104 Centre of Milky Way15.4 2.44 × 106 2.44 × 106 Andromeda galaxy18.4 4.9 × 107 4.9 × 107 Virgo cluster19.2 1.1 × 108 1.1 × 108 Coma cluster25.3 5 × 1010 5 × 1010 Edge of observable Universe
the time and space coordinates of the accelerating person, relative to a person who remains at rest, are those
of a point on the spacetime wheel, namely
t, x = sinhα, coshα . (1.23)
In the case where the acceleration is one Earth gravity, g = 9.80665 m s−2, the unit of time is
c
g=
299,792,458 m s−1
9.80665 m s−2= 0.97 yr , (1.24)
just short of one year. For simplicity, Table 1.1, which tabulates some milestones along the way, takes the
unit of time to be exactly one year, which would be the case if one were accelerating at 0.97 g = 9.5 m s−2.
After a slow start, you cover ground at an ever increasing rate, crossing 50 billion lightyears, the distance
to the edge of the currently observable Universe, in just over 25 years of your own time.
Does this mean you go faster than the speed of light? No. From the point of view of a person at rest
on Earth, you never go faster than the speed of light. From your own point of view, distances along your
direction of motion are Lorentz-contracted, so distances that are vast from Earth’s point of view appear
much shorter to you. Fast as the Universe rushes by, it never goes faster than the speed of light.
This rosy picture of being able to flit around the Universe has drawbacks. Firstly, it would take a huge
amount of energy to keep you accelerating at g. Secondly, you would use up a huge amount of Earth time
travelling around at relativistic speeds. If you took a trip to the edge of the Universe, then by the time
you got back not only would all your friends and relations be dead, but the Earth would probably be gone,
swallowed by the Sun in its red giant phase, the Sun would have exhausted its fuel and shrivelled into a
cold white dwarf star, and the Solar System, having orbited the Galaxy a thousand times, would be lost
somewhere in its milky ways.
26 Special Relativity
Technical point. The Universe is expanding, so the distance to the edge of the currently observable Universe
is increasing. Thus it would actually take longer than indicated in the table to reach the edge of the currently
observable Universe. Moreover if the Universe is accelerating, as recent evidence from the Hubble diagram
of Type Ia Supernovae suggests, then you will never be able to reach the edge of the currently observable
Universe, however fast you go.
1.8 Scalar spacetime distance
One of the most fundamental features of a Lorentz transformation is that its leaves invariant a certain
distance. the scalar spacetime distance, between any two events in spacetime. The scalar spacetime distance
∆s between two events separated by ∆t,∆x,∆y,∆z is given by
∆s2 = −∆t2 + ∆r2
= −∆t2 + ∆x2 + ∆y2 + ∆z2 . (1.25)
A quantity such as ∆s2 that remains unchanged under any Lorentz transformation is called a scalar. It is
left to you in exercise 1.9 to show explicitly that ∆s2 is unchanged under Lorentz transformations. Lorentz
transformations can be defined as linear spacetime transformations that leave ∆s2 invariant.
The single scalar spacetime squared interval ∆s2 replaces the two scalar quantities
time interval ∆t
distance interval ∆r =√
∆x2 + ∆y2 + ∆z2(1.26)
of classical Galilean spacetime.
Exercise 1.9 Invariant spacetime interval. Show that the squared spacetime interval ∆s2 defined by
equation (1.25) is unchanged by a Lorentz transformation, that is, show that
−∆t′2 + ∆x′2 + ∆y′2 + ∆z′2 = −∆t2 + ∆x2 + ∆y2 + ∆z2 . (1.27)
You may assume without proof the familiar results that the 3D scalar product ∆r2 = ∆x2 + ∆y2 + ∆z2 is
unchanged by a spatial rotation, so it suffices to consider a Lorentz boost, say in the x direction.
1.8.1 Proper time, proper distance
The scalar spacetime squared interval ∆s2 has a physical meaning.
If an interval ∆t,∆r is timelike, ∆t > ∆r, then the square root of minus the spacetime interval squared
is the proper time ∆τ along it
∆τ =√
−∆s2 =√
∆t2 −∆r2 . (1.28)
This is the time experienced by an observer moving along that interval.
1.9 4-vectors 27
If an interval ∆t,∆r is spacelike, ∆t < ∆r, then the spacetime interval equals the proper distance
∆l along it
∆l =√
∆s2 =√
∆r2 −∆t2 . (1.29)
This is the distance between two events measured by an observer for whom those events are simultaneous.
Concept question 1.10 Justify the assertions (1.28) and (1.29).
1.8.2 Timelike, lightlike, spacelike
A spacetime interval ∆xm is called
timelike if ∆s2 < 0 ,
null or lightlike if ∆s2 = 0 ,
spacelike if ∆s2 > 0 .
(1.30)
1.8.3 Minkowski metric
The scalar spacetime squared interval ∆s2 associated with an interval ∆xm = ∆t,∆r = ∆t,∆x,∆y,∆zcan be written
∆s2 = ∆xm∆xm
= ηmn∆xm∆xn (1.31)
where ηmn is the Minkowski metric
ηmn ≡
−1 0 0 0
0 1 0 0
0 0 1 0
0 0 0 1
. (1.32)
Equation (1.31) uses the implicit summation convention, according to which paired indices are explicitly
summed over. Invariably, one of a pair of repeated indices is raised, the other lowered.
1.9 4-vectors
A 4-vector in special relativity is a quantity am = at, ax, ay, az that transforms under Lorentz transfor-
mations like an interval xm = t, x, y, z of spacetime
a′m = Lmna
n (1.33)
where Lmn denotes a Lorentz transformation.
28 Special Relativity
1.9.1 Index notation
In special and general relativity it is convenient to introduce two versions of the same 4-vector quantity, one
with raised indices, called the contravariant components of the 4-vector,
am ≡ at, ax, ay, az , (1.34)
and one with lowered indices called the covariant components of the 4-vector,
am ≡ −at, ax, ay, az (1.35)
(the naming is crazy, and you do not need to remember it).
The indices run over m = t, x, y, z, or sometimes m = 0, 1, 2, 3.
Why introduce raised and lowered indices? Because
amam ≡
∑
m
amam = ata
t + axax + aya
y + azaz
= − (at)2 + (ax)2 + (ay)2 + (az)2 (1.36)
is a Lorentz scalar.
1.10 Energy-momentum 4-vector
Symmetry argument:
Symmetry Conservation law
Time translation Energy
Space translation Momentum
suggests
energy = time component
momentum = space component
of 4-vector. (1.37)
The Principle of Special Relativity requires that the equation of energy-momentum conservation
energy
momentum= constant (1.38)
should take the same form in any inertial frame. The equation should be Lorentz covariant, that is, the
equation should transform like a Lorentz 4-vector.
1.10 Energy-momentum 4-vector 29
1.10.1 Construction of the energy-momentum 4-vector
Require:
1. it’s a 4-vector
2. goes over to the Newtonian limit as v → 0.
Newtonian limit:
Momentum p is mass m times velocity v
p = mv = mdr
dt. (1.39)
4D version:
Need to do two things to Newtonian momentum:
• replace r by a 4-vector xm = t, r• replace dt by a scalar — the only obvious choice is the proper time interval τ .
Result:
pk = mdxk
dτ
= m
dt
dτ,dr
dτ
= m γ, γv (1.40)
which are special relativistic versions of energy E and momentum p
pk = E,p = mγ,mγv . (1.41)
1.10.2 Special relativistic energy
E = mγ (units c = 1) (1.42)
or, restoring standard units
E = mc2γ . (1.43)
Taylor expand γ for small velocity v:
γ =1
√
1− v2/c2= 1 +
1
2
v2
c2+ ... (1.44)
so
E = mc2(
1 +1
2
v2
c2+ ...
)
= mc2 +1
2mv2 + ... . (1.45)
The first term, mc2, is the rest-mass energy. The second term, 12mv
2, is the non-relativistic kinetic energy.
Higher-order terms give relativistic corrections to the kinetic energy.
30 Special Relativity
1.10.3 Rest mass is a scalar
The scalar quantity constructed from the energy-momentum 4-vector pk = E,p is
pkpk = −E2 + p2
= −m2(γ2 − γ2v2)
= −m2 (1.46)
minus the square of the rest mass.
1.11 Photon energy-momentum
Photons have zero rest mass
m = 0 . (1.47)
Thus
pkpk = −E2 + p2 = −m2 = 0 (1.48)
whence
p ≡ |p| = E . (1.49)
Hence
pk = E,p= E1,n= hν1,n (1.50)
where ν is the photon frequency.
The photon velocity is n, a unit vector. The photon speed is one (the speed of light).
1.11.1 Lorentz transformation of photon energy-momentum 4-vector
Follows the usual rules for 4-vectors.
In the case that the Lorentz transformation is a Lorentz boost along the x-axis, the transformation is
p′t
p′x
p′y
p′z
=
γ −γv 0 0
−γv γ 0 0
0 0 1 0
0 0 0 1
pt
px
py
pz
=
γ(pt − vpx)
γ(px − vpt)
py
pz
. (1.51)
1.12 Abstract 4-vectors 31
Equivalently
hν′
1
n′x
n′y
n′z
=
γ −γv 0 0
−γv γ 0 0
0 0 1 0
0 0 0 1
hν
1
nx
ny
nz
= hν
γ(1− nxv)
γ(nx − v)ny
nz
.
These mathematical relations imply the rules of 4-dimensional perspective, §1.13.1.
1.11.2 Redshift
Astronomers define the redshift z of a photon by
z ≡ λobs − λemit
λemit. (1.52)
In relativity, it is often more convenient to use the redshift factor 1 + z
1 + z ≡ λobs
λemit=νemit
νobs. (1.53)
1.11.3 Special relativistic Doppler shift
If the emitter frame (primed) is moving with velocity v in the x-direction relative to the observer frame
(unprimed) then
hνemit = hνobsγ(1− nxv) (1.54)
so
1 + z =νemit
νobs
= γ(1− nxv)
= γ(1− n.v) . (1.55)
This is the general formula for the special relativistic Doppler shift.
1.12 Abstract 4-vectors
A = Amγγm (1.56)
32 Special Relativity
1.13 What things look like at relativistic speeds
1.13.1 The rules of 4-dimensional perspective
The diagram below illustrates the rules of 4-dimensional perspective, also called “special relativistic beam-
ing,” which describe how a scene appears when you move through it at near light speed.
1
1γv
γ
On the left, you are at rest relative to the scene. Imagine painting the scene on a celestial sphere around
you. The arrows represent the directions of light rays (photons) from the scene on the celestial sphere to
you at the center.
On the right, you are moving to the right through the scene, at some fraction of the speed of light. The
celestial sphere is stretched along the direction of your motion into a celestial ellipsoid. You, the observer,
are not at the center of the ellipsoid, but rather at one of its foci (the left one, if you are moving to the
right). The scene appears relativistically aberrated, which is to say concentrated ahead of you, and expanded
behind you.
The lengths of the arrows are proportional to the energies, or frequencies, of the photons that you see.
When you are moving through the scene at near light speed, the arrows ahead of you, in your direction
of motion, are longer than at rest, so you see the photons blue-shifted, increased in energy, increased in
frequency. Conversely, the arrows behind you are shorter than at rest, so you see the photons red-shifted,
decreased in energy, decreased in frequency. Since photons are good clocks, the change in photon frequency
also tells you how fast or slow clocks attached to the scene appear to you to run.
Numbers? On the right, you are moving through the scene at v = 0.6 c. The celestial ellipsoid is stretched
along the direction of your motion by the Lorentz gamma factor, which here is γ = 1/√
1− 0.62 = 1.25. The
focus of the celestial ellipsoid, where you the observer are, is displaced from center by γv = 1.25×0.6 = 0.75.
1.14 How to programme Lorentz transformations on a computer 33
1.14 How to programme Lorentz transformations on a computer
3D gaming programmers are familiar with the fact that the best way to program spatial rotations on a
computer is with quaternions. Compared to standard rotation matrices, quaternions offer increased speed
and require less storage, and their algebraic properties simplify interpolation and splining.
Section 1.7 showed that a Lorentz boost is mathematically equivalent to a rotation by an imaginary
angle. Thus suggests that Lorentz transformations might be treated as complexified spatial rotations, which
proves to be true. Indeed, the best way to program Lorentz transformation on a computer is with complex
quaternions, as will be demonstrated in Chapter 14.
Exercises
Exercise 1.11 Twin paradox. Your twin leaves you on Earth and travels to the spacestation Alpha, ℓ
= 3 lyr away, at a good fraction of the speed of light, then immediately returns to Earth at the same speed.
The accompanying spacetime diagram shows the corresponding worldlines of both you and your twin. Aside
from part (a) and the first part of (b), I want you to derive your answers mathematically, using logic and
Lorentz transformations. However, the diagram is accurately drawn, and you should be able to check your
answers by measuring.
1. Label the worldlines of you and your twin. Draw the worldline of a light signal which travels from you on
Earth, hits Alpha just when your twin arrives, and immediately returns to Earth. Draw the twin’s “now”
when just arriving at Alpha, and the twin’s “now” just departing from Alpha (in the first case the twin is
moving toward Alpha, while in the second case the twin is moving back toward Earth).
2. From the diagram, measure the twin’s speed v relative to you, in units where the speed of light is unity,
c = 1. Deduce the Lorentz gamma factor γ, and the redshift factor 1 + z = [(1 + v)/(1 − v)]1/2, in the
cases (i) where the twin is receding, and (ii) where the twin is approaching.
3. Choose the spacetime origin to be the event where the twin leaves Earth. Argue that the position 4-vector
of the twin on arrival at Alpha is
t, x, y, z = ℓ/v, ℓ, 0, 0 . (1.57)
Lorentz transform this 4-vector to determine the position 4-vector of the twin on arrival at Alpha, in the
twin’s frame. Express your answer first in terms of ℓ, v, and γ, and then in (light)years. State in words
what this position 4-vector means.
4. How much do you and your twin age respectively during the round trip to Alpha and back? What is the
ratio of these ages? Express your answers first in terms of ℓ, v, and γ, and then in years.
5. What is the distance between the Earth and Alpha from the twin’s point of view? What is the ratio of
this distance to the distance between Earth and Alpha from your point of view? Explain how your arrived
at your result. Express your answer first in terms of ℓ, v, and γ, and then in lightyears.
6. You watch your twin through a telescope. How much time do you see (through the telescope) elapse
34 Special Relativity
on your twin’s wristwatch between launch and arrival on Alpha? How much time passes on your own
wristwatch during this time? What is the ratio of these two times? Express your answers first in terms of
ℓ, v, and γ, and then in years.
7. On arrival at Alpha, your twin looks back through a telescope at your wristwatch. How much time does
your twin see (through the telescope) has elapsed since launch on your watch? How much time has elapsed
on the twin’s own wristwatch during this time? What is the ratio of these two times? Express your answers
first in terms of ℓ, v, and γ, and then in years.
8. You continue to watch your twin through a telescope. How much time elapses on your twin’s wristwatch,
as seen by you through the telescope, during the twin’s journey back from Alpha to Earth? How much
time passes on your own watch as you watch (through the telescope) the twin journey back from Alpha
to Earth? What is the ratio of these two times? Express your answers first in terms of ℓ, v, and γ, and
then in years.
9. During the journey back from Alpha to Earth, your twin likewise continues to look through a telescope
at the time registered on your watch. How much time passes on your wristwatch, as seen by your twin
through the telescope, during the journey back? How much time passes on the twin’s wristwatch from the
twin’s point of view during the journey back? What is the ratio of these two times? Express your answers
first in terms of ℓ, v, and γ, and then in years.
Exercise 1.12 Lines intersecting at right angles. Prove that if two lines appear to intersect at right-
angles projected on the sky in one frame, then they appear to intersect at right-angles in another frame
Lorentz-transformed with respect to the first.
PART TWO
COORDINATE APPROACH TO GENERAL RELATIVITY
Concept Questions
1. What assumption of general relativity makes it possible to introduce a coordinate system?
2. Is the speed of light a universal constant in general relativity? If so, in what sense?
3. What does “locally inertial” mean? How local is local?
4. Why is spacetime locally inertial?
5. What assumption of general relativity makes it possible to introduce clocks and rulers?
6. Consider two observers at the same point and with the same instantaneous velocity, but one is accelerating
and the other is in free-fall. What is the relation between the proper time or proper distance along an
infinitesimal interval measured by the two observers? What assumption of general relativity implies this?
7. Does the (Strong) Principle of Equivalence imply that two unequal masses will fall at the same rate in a
gravitational field? Explain.
8. In what respects is the Strong Principle of Equivalence (gravity is equivalent to acceleration) stronger
than the Weak Principle of Equivalence (gravitating mass equals inertial mass)?
9. Standing on the surface of the Earth, you hold an object of negative mass in your hand, and drop it.
According to the Principle of Equivalence, does the negative mass fall up or down?
10. Same as the previous question, but what does Newtonian gravity predict?
11. You have a box of negative mass particles, and you remove energy from it. Do the particles move faster
or slower? Does the entropy of the box increase or decrease? Does the pressure exerted by the particles
on the walls of the box increase or decrease?
12. You shine two light beams along identical directions in a gravitational field. The two light beams are
identical in every way except that they have two different frequencies. Does the Equivalence Principle
imply that the interference pattern produced by each of the beams individually is the same?
13. What is a “straight line”, according to the Principle of Equivalence?
14. If all objects move on straight lines, how is it that when, standing on the surface of the Earth, you throw
two objects in the same direction but with different velocities, they follow two different trajectories?
15. In relativity, what is the generalization of the “shortest distance between two points”?
16. What kinds of general coordinate transformations are allowed in general relativity?
38 Concept Questions
17. In general relativity, what is a scalar? A 4-vector? A tensor? Which of the following is a scalar/vector/
tensor/none-of-the-above? (a) a set of coordinates xµ; (b) a coordinate interval dxµ; (c) proper time
τ?
18. What does general covariance mean?
19. What does parallel transport mean?
20. Why is it important to define covariant derivatives that behave like tensors?
21. Is covariant differentiation a derivation? That is, is covariant differentiation a linear operation, and does
it obey the Leibniz rule for the derivative of a product?
22. What is the covariant derivative of the metric tensor? Explain.
23. What does a connection coefficient Γκµν mean physically? Is it a tensor? Why, or why not?
24. An astronaut is in free-fall in orbit around the Earth. Can the astronaut detect that there is a gravitational
field?
25. Can a gravitational field exist in flat space?
26. How can you tell whether a given metric is equivalent to the Minkowski metric of flat space?
27. How many degrees of freedom does the metric have? How many of these degrees of freedom can be removed
by arbitrary transformations of the spacetime coordinates, and therefore how many physical degrees of
freedom are there in spacetime?
28. If you insist that the spacetime is spherical, how many physical degrees of freedom are there in the
spacetime?
29. If you insist that the spacetime is spatially homogeneous and isotropic (the cosmological principle), how
many physical degrees of freedom are there in the spacetime?
30. In general relativity, you are free to prescribe any spacetime (any metric) you like, including metrics with
wormholes and metrics that connect the future to the past so as to violate causality. True or false?
31. If it is true that in general relativity you can prescribe any metric you like, then why aren’t you bumping
into wormholes and causality violations all the time?
32. How much mass does it take to curve space significantly (significantly meaning by of order unity)?
33. What is the relation between the energy-momentum 4-vector of a particle and the energy-momentum
tensor?
34. It is straightforward to go from a prescribed metric to the energy-momentum tensor. True or false?
35. It is straightforward to go from a prescribed energy-momentum tensor to the metric. True or false?
36. Does the Principle of Equivalence imply Einstein’s equations?
37. What do Einstein’s equations mean physically?
38. What does the Riemann curvature tensor Rκλµν mean physically? Is it a tensor?
39. The Riemann tensor splits into compressive (Ricci) and tidal (Weyl) parts. What do these parts mean,
physically?
40. Einstein’s equations imply conservation of energy-momentum, but what does that mean?
41. Do Einstein’s equations describe gravitational waves?
42. Do photons (massless particles) gravitate?
43. How do different forms of mass-energy gravitate?
44. How does negative mass gravitate?
What’s important?
This part of the notes adopts the traditional coordinate-based approach to general relativity. The approach
is neither the most insightful nor the most powerful, but it is the fastest route to connecting the metric to
the energy-momentum content of spacetime.
1. Postulates of general relativity. How do the various postulates imply the mathematical structure of general
relativity?
2. The road from spacetime curvature to energy-momentum:
metric gµν
→ connection coefficients Γκµν
→ Riemann curvature tensor Rκλµν
→ Ricci tensor Rκµ and scalar R
→ Einstein tensor Gκµ = Rκµ − 12gκµR
→ energy-momentum tensor Tκµ
3. 4-velocity and 4-momentum. Geodesic equation.
4. Bianchi identities guarantee conservation of energy-momentum.
2
Fundamentals of General Relativity
2.1 The postulates of General Relativity
General relativity follows from three postulates:
1. Spacetime is a 4-dimensional manifold;
2. The (Strong) Principle of Equivalence;
3. Einstein’s Equations.
2.1.1 Spacetime is a 4-dimensional manifold
A 4-dimensional manifold is defined mathematically to be a topological space that is locally homeomorphic
to Euclidean 4-space R4.
This postulate implies that it is possible to set up a coordinate system (possibly in patches)
xµ ≡ x0, x1, x2, x3 (2.1)
such that each point of (the patch of) spacetime has a unique coordinate.
Andrew’s convention:
Greek (brown) dummy indices label curved spacetime coordinates.
Latin (black) dummy indices label locally inertial (more generally, tetrad) coordinates.
2.1.2 (Strong) Principle of Equivalence (PE)
“The laws of physics in a gravitating frame are equivalent to those in an accelerating frame”.
The Weak Principle of Equivalence is “Gravitating mass = inertial mass”.
PE ⇒ spacetime is locally inertial (see §2.2).
2.2 Existence of locally inertial frames 41
2.1.3 Einstein’s equations
Einstein’s equations comprise a 4× 4 symmetric matrix of equations
Gµν = 8πGTµν . (2.2)
Here G is the Newtonian gravitational constant, Gµν is the Einstein tensor, and Tµν is the energy-
momentum tensor.
Physically, Einstein’s equations signify
(compressive part of) curvature = energy-momentum content . (2.3)
Einstein’s equations generalize Poisson’s equation
∇2Φ = 4πGρ (2.4)
where Φ is the Newtonian gravitational potential, and ρ the mass-energy density. Poisson’s equation is the
time-time component of Einstein’s equations in the limit of a weak gravitational field and slowly moving
matter.
2.2 Existence of locally inertial frames
The Principle of Equivalence implies that at each point of spacetime it is possible to choose a locally inertial,
or free-fall, frame, such that the laws of special relativity apply within an infinitesimal neighbourhood of
that point. By this is meant that at each point of spacetime it is possible to choose coordinates such that
(a) the metric at that point is Minkowski, and (b) the first derivatives of the metric are all zero.
It is built into the Principle of Equivalence that general relativity is, like special relativity, a metric
theory. Notably, the proper times and distances measured by an accelerating observer are the same as
those measured by a freely-falling observer at the same point and with the same instantaneous velocity.
2.3 Metric
The metric is the essential mathematical object that converts an infinitesimal coordinate interval
dxµ ≡ dx0, dx1, dx2, dx3 (2.5)
to a proper measurement of an interval of time or space.
Postulate (1) of general relativity means that it is possible to choose coordinates
xµ ≡ x0, x1, x2, x3 (2.6)
covering (a patch of) spacetime.
42 Fundamentals of General Relativity
Postulate (2) of general relativity implies that at each point of spacetime it is possible to choose locally
inertial coordinates
ξm ≡ ξ0, ξ1, ξ2, ξ3 (2.7)
such that the metric is Minkowski,
ds2 = ηmn dξmdξn , (2.8)
in an infinitesimal neighborhood of the point. The spacetime distance squared ds2 is a scalar, a quantity
that is unchanged by the choice of coordinates.
Since
dξm =∂ξm
∂xµdxµ (2.9)
it follows that
ds2 = ηmn∂ξm
∂xµ
∂ξn
∂xνdxµdxν (2.10)
so the scalar spacetime distance squared is
ds2 = gµν dxµdxν (2.11)
where gµν is the metric, a 4× 4 symmetric matrix
gµν = ηmn∂ξm
∂xµ
∂ξn
∂xν. (2.12)
2.4 Basis gµ of tangent vectors
You are familiar with the idea that in ordinary 3D Euclidean geometry it is often convenient to treat vectors
in an abstract coordinate-independent formalism. Thus for example a 3-vector is commonly written as an
abstract quantity r. The coordinates of the vector r may be x, y, z in some particular coordinate system,
but one recognizes that the vector r has a meaning, a magnitude and a direction, that is independent of the
coordinate system adopted. In an arbitrary Cartesian coordinate system, the Euclidean 3-vector r can be
expressed
r =∑
i
xi xi = x x+ y y + z z (2.13)
where xi ≡ x, y, z are unit vectors along each of the coordinate axes.
The same kind of abstract notation is useful in general relativity. Define gµ
gµ ≡ g0, g1, g2, g3 (2.14)
to be the basis of axes tangent to the coordinates xµ. Each axis gµ is a 4D vector object, with both magnitude
and direction in spacetime. (Some texts represent the tangent vectors gµ with the notation ∂µ, but this
notation is not used here, to avoid the potential confusion between ∂µ as a derivative and ∂µ as a vector.)
2.5 4-vectors and tensors 43
An interval dxµ of spacetime can be expressed in coordinate-independent fashion as the abstract vector
interval dx
dx ≡ gµ dxµ = g0 dx
0 + g1 dx1 + g2 dx
2 + g3 dx3 . (2.15)
The scalar length squared of the abstract vector interval dx is
ds2 = dx · dx = gµ · gν dxµdxν (2.16)
whence
gµν = gµ · gν (2.17)
the metric is the (4D) scalar product of tangent vectors.
The tangent vectors gµ form a basis for a 4D tangent space that has three important mathematical
properties. First, the tangent space is a vector space, that is, it has the properties of linearity that define
a vector space. Second, the tangent space has an inner (or scalar) product, defined by the metric (2.17).
Third, vectors in the tangent space can be differentiated with respect to coordinates, as will be elucidated
in §2.6.3.
2.5 4-vectors and tensors
2.5.1 Contravariant coordinate 4-vector
Under a general coordinate transformation
xµ → x′µ (2.18)
a coordinate interval dxµ transforms as
dx′µ =∂x′µ
∂xνdxν . (2.19)
In general relativity, a coordinate 4-vector is defined to be a quantity Aµ = A0, A1, A2, A3 that trans-
forms under a coordinate transformation (2.18) like a coordinate interval
A′µ =∂x′µ
∂xνAν . (2.20)
2.5.2 Abstract 4-vector
A 4-vector may be written in coordinate-independent fashion as
A = gµAµ . (2.21)
The quantity A is an abstract 4-vector. Although A is a 4-vector, it is by construction unchanged by a
coordinate transformation, and is therefore a coordinate scalar. See §2.5.6 for commentary on the distinction
between abstract and coordinate vectors.
44 Fundamentals of General Relativity
2.5.3 Lowering and raising indices
Define gµν to be the inverse metric, satisfying
gλµ gµν = δν
λ =
1 0 0 0
0 1 0 0
0 0 1 0
0 0 0 1
. (2.22)
The metric gµν and its inverse gµν provide the means of lowering and raising coordinate indices. The
components of a coordinate 4-vector Aµ with raised index are called its contravariant components, while
those Aµ with lowered indices are called its covariant components,
Aµ = gµν Aν , (2.23)
Aµ = gµν Aν . (2.24)
2.5.4 Covariant coordinate 4-vector
Under a general coordinate transformation (2.18), the covariant components Aµ of a coordinate 4-vector
transform as
A′µ =
∂xν
∂x′µAν . (2.25)
You can check that the transformation law (2.25) for the covariant components Aµ is consisistent with the
transformation law (2.20) for the contravariant components Aµ.
You can check that the tangent vectors gµ transform as a covariant coordinate 4-vector.
2.5.5 Scalar product
If Aµ and Bµ are coordinate 4-vectors, then their scalar product is
AµBµ = AµBµ = gµνA
µBν . (2.26)
This is a coordinate scalar, a quantity that remains invariant under general coordinate transformations.
In abstract vector formalism, the scalar product of two 4-vectors A = gµAµ and B = gµB
µ is
A ·B = gµ · gν AµBν = gµνA
µBν . (2.27)
2.5.6 Comment on vector naming and notation
Different texts follow different conventions for naming and notating vectors and tensors.
In this book I follow the convention of calling both Aµ (with a dummy index µ) and A ≡ Aµgµ vectors.
Although Aµ and A are both vectors, they are mathematically different objects.
2.6 Covariant derivatives 45
If the index on a vector indicates a specific coordinate, then the indexed vector is the component of the
vector; for example A0 (or At) is the x0 (or time t) component of the coordinate 4-vector Aµ.
In this book, the different species of vector are distinguished by an adjective:
1. A coordinate vector Aµ, identified by Greek (brown) indices µ, is one that changes in a prescribed way
under coordinate transformations. A coordinate transformation is one that changes the coordinates of the
spacetime without actually changing the spacetime or whatever lies in it.
2. An abstract vector A, identified by boldface, is the thing itself, and is unchanged by the choice of
coordinates. Since the abstract vector is unchanged by a coordinate transformation, it is a coordinate
scalar.
All the types of vector have the properties of linearity (additivity, multiplication by scalars) that identify
them mathematically as belonging to vector spaces. The important distinction between the types of vector
is how they behave under transformations.
In referring to both Aµ and A as vectors, I am following the standard physics practice of mentally regarding
Aµ and A as equivalent objects. You are familiar with the advantages of treating a vector in 3D Euclidean
space either as an abstract vector A, or as a coordinate vector Ai. Depending on the problem, sometimes
the abstract notation A is more convenient, and sometimes the coordinate notation Ai is more convenient.
Sometimes it’s convenient to switch between the two in the middle of a calculation. Likewise in general
relativity it is convenient to have the flexibility to work in either coordinate or abstract notation, whatever
suits the problem of the moment.
2.5.7 Coordinate tensor
In general, a coordinate tensor Aκλ...µν... is an object that transforms under general coordinate transforma-
tions (2.18) as
A′κλ...µν... =
∂x′κ
∂xα
∂x′λ
∂xβ...∂xγ
∂x′µ∂xδ
∂x′ν... Aαβ...
γδ... . (2.28)
You can check that the metric tensor gµν and its inverse gµν are indeed coordinate tensors, transforming
like (2.28).
The rank of a tensor is the number of indices. A scalar is a tensor of rank 0. A 4-vector is a tensor of
rank 1.
2.6 Covariant derivatives
2.6.1 Derivative of a coordinate scalar
Suppose that Φ is a coordinate scalar. Then the coordinate derivative of Φ is a coordinate 4-vector
∂Φ
∂xµis a coordinate tensor (2.29)
46 Fundamentals of General Relativity
transforming like equation (2.25).
As a shorthand, the ordinary partial derivative is often denoted in the literature with a comma
∂Φ
∂xµ= Φ,µ . (2.30)
For the most part this book does not use the comma notation.
2.6.2 Derivative of a coordinate 4-vector
The ordinary partial derivative of a covariant coordinate 4-vector Aµ is not a tensor
∂Aµ
∂xνis not a coordinate tensor (2.31)
because it does not transform like a coordinate tensor.
However, the 4-vector A = gµAµ, being by construction invariant under coordinate transformations, is a
coordinate scalar, and its partial derivative is a coordinate 4-vector
∂A
∂xν=∂gµA
µ
∂xν
= gµ∂Aµ
∂xν+∂gµ
∂xνAµ is a coordinate tensor . (2.32)
The last line of equation (2.32) assumes that it is legitimate to differentiate the tangent vectors gµ, but
what does this mean? The partial derivatives of basis vectors gµ are defined in the usual way by
∂gµ
∂xν≡ lim
δxν→0
gµ(x0, ..., xν+δxν , ..., x3)− gµ(x0, ..., xν , ..., x3)
δxν. (2.33)
This definition relies on being able to compare the vectors gµ(x) at some point x with the vectors gµ(x+δx)
at another point x+δx a small distance away. The comparison between two vectors a small distance apart
is made possible by the existence of locally inertial frames. In a locally inertial frame, two vectors a small
distance apart can be compared by parallel-transporting one vector to the location of the other along
the small interval between them, that is, by transporting the vector without accelerating or precessing with
respect to the locally inertial frame. Thus gµ(x+δx) in the definition (2.33) should be interpreted as its
value parallel-transported from position x+δx to position x along the small interval δx between them.
2.6.3 Coordinate connection coefficients (Christoffel symbols)
The partial derivatives of the basis vectors gµ that appear on the right hand side of equation (2.32) define
the coordinate connection coefficients Γκµν , also known as Christoffel symbols,
∂gµ
∂xν≡ Γκ
µν gκ is not a coordinate tensor . (2.34)
The definition (2.34) shows that the connection coefficients express how each tangent vector gµ changes,
relative to parallel-transport, when shifted along an interval δxν .
2.6 Covariant derivatives 47
2.6.4 Covariant derivative of a contravariant 4-vector
Expression (2.32) along with the definition (2.34) of the connection coefficients implies that
∂A
∂xν= gµ
∂Aµ
∂xν+ Γκ
µνgκAµ
= gκ
(
∂Aκ
∂xν+ Γκ
µνAµ
)
is a coordinate tensor . (2.35)
The expression in parentheses is a coordinate tensor, and defines the covariant derivative DνAκ of the
contravariant coordinate 4-vector Aκ
DνAκ ≡ ∂Aκ
∂xν+ Γκ
µνAµ is a coordinate tensor . (2.36)
As a shorthand, the covariant derivative is often denoted in the literature with a semi-colon
DνAκ = Aκ
;ν . (2.37)
For the most part this book does not use the semi-colon notation.
2.6.5 Covariant derivative of a covariant coordinate 4-vector
Similarly,
∂A
∂xν= gκDνAκ is a coordinate tensor (2.38)
where DνAκ is the covariant derivative of the covariant coordinate 4-vector Aκ
DνAκ ≡∂Aκ
∂xν− Γµ
κνAµ is a coordinate tensor . (2.39)
2.6.6 Covariant derivative of a coordinate tensor
In general, the covariant derivative of a coordinate tensor is
DαAκλ...µν... =
∂Aκλ...µν...
∂xα+ Γκ
βαAβλ...µν... + Γλ
βαAκβ...µν... + ...− Γβ
µαAκλ...βν... − Γβ
ναAκλ...µβ... − ... (2.40)
with a positive Γ term for each contravariant index, and a negative Γ term for each covariant index.
2.6.7 No-torsion condition
The existence of locally inertial frames requires that it must be possible to arrange not only that the tangent
axes gµ are orthonormal at a point, but also that they remain orthonormal to first order in a Taylor expansion
48 Fundamentals of General Relativity
about the point. That is, it must be possible to choose the coordinates such that the tangent axes gµ are
orthonormal, and unchanged to linear order:
gµ · gν = ηµν , (2.41)
∂gµ
∂xν= 0 . (2.42)
In view of the definition (2.34) of the connection coefficients, the second condition (2.42) is equivalent to the
vanishing of all the connection coefficients:
Γκµν = 0 . (2.43)
Under a general coordinate transformation xµ → x′µ, the tangent axes transform as gµ → g′µ = ∂xν/∂x′µ gν .
The 4×4 matrix ∂xν/∂x′µ of partial derivatives provides 16 degrees of freedom in choosing the tangent axes
at a point. The 16 degrees of freedom are enough — more than enough — to accomplish the orthonormality
condition (2.41), which is a symmetric 4 × 4 matrix equation with 10 degrees of freedom. The additional
16− 10 = 6 degrees of freedom are Lorentz transformations, which rotate the tangent axes gµ, but leave the
metric ηµν unchanged.
Just as it is possible to reorient the tangent axes gµ at a point by adjusting the matrix ∂x′ν/∂xµ of first
partial derivatives of the coordinate transformation xµ → x′µ, so also it is possible to reorient the derivatives
∂gµ/∂xν of the tangent axes by adjusting the matrix ∂2x′ν/∂xλ∂xµ of second partial derivatives. The second
partial derivatives comprise a set of 4 symmetric 4×4 matrices, for a total of 4×10 = 40 degrees of freedom.
However, there are 4 × 4 × 4 = 64 connection coefficients Γκµν , all of which the condition (2.43) requires to
vanish. The matrix of second derivatives is thus 64 − 40 = 24 degrees of freedom short of being able to
make all the connections vanish. The resolution of the problem is that, as shown below, equation (2.51),
there are 24 combinations of the connections that form a tensor, the torsion tensor. If a tensor is zero in one
frame, then it is automatically zero in any other frame. Thus the requirement that all the connections vanish
requires that the torsion tensor vanish. This requires, from the expression (2.51) for the torsion tensor, the
no-torsion condition that the connection coefficients are symmetric in their last two indices
Γκµν = Γκ
νµ . (2.44)
It should be emphasized that the condition of vanishing torsion is an assumption of general relativity, not
a mathematical necessity. It has been shown in this section that torsion vanishes if and only if spacetime is
locally flat, meaning that at any point coordinates can be found such that conditions (2.41) are true. The
assumption of local flatness is central to the idea of the principle of equivalence. But it is an assumption,
not a consequence, of the theory.
Concept question 2.1 If torsion does not vanish, then there is no locally inertial frame. What does
parallel-transport mean in such a case?
2.6 Covariant derivatives 49
2.6.8 Aside: torsion and the integrability of the position vector
From the definitions (2.34) of the connection coefficients, the no-torsion condition (2.44) is equivalent to
∂gµ
∂xν− ∂gν
∂xµ= 0 . (2.45)
According to Frobenius’ theorem, this condition (2.45) is precisely the condition for the system gµ to be
integrable, that is, there exists a position vector X whose partial derivatives are
gµ =∂X
∂xµ. (2.46)
Equivalently, the total differential of the position vector X is
dX = gµdxµ . (2.47)
The abstract vector interval dx ≡ gµ dxµ was defined by equation (2.15) as the coordinate-independent
version of a spacetime interval dxµ. The notation dx was merely symbolic: dx was not necessarily a total
differential of something. However, the no-torsion condition (2.47) implies that dx is in fact the total
differential dX of the position vector X
dx = dX . (2.48)
The no-torsion condition (2.44) is equivalent to the commutation of partial derivatives of X:
Γκµν gκ ≡
∂gµ
∂xν=
∂2X
∂xν∂xµ=∂gν
∂xµ≡ Γκ
νµ gκ . (2.49)
The physical meaning of torsion is discussed further in §3.4.
2.6.9 Torsion tensor
General relativity assumes no torsion, but it is possible to consider generalizations to theories with torsion.
The torsion tensor Sµκλ is defined by the commutator of the covariant derivative acting on a scalar Φ
[Dκ, Dλ] Φ = Sµκλ
∂Φ
∂xµis a coordinate tensor . (2.50)
Note that the covariant derivative of a scalar is just the ordinary derivative, DλΦ = ∂Φ/∂xλ. The expres-
sion (2.39) for the covariant derivatives shows that the torsion tensor is
Sµκλ = Γµ
κλ − Γµλκ is a coordinate tensor (2.51)
which is evidently antisymmetric in the indices κλ.
In Einstein-Cartan theory, the torsion tensor is related to the spin content of spacetime. Since this vanishes
in empty space, Einstein-Cartan theory is indistinguishable from general relativity in experiments carried
out in vacuum.
50 Fundamentals of General Relativity
Exercise 2.2 Show that
DκAλ −DλAκ =∂Aλ
∂xκ− ∂Aκ
∂xλ+ Sµ
κλAµ . (2.52)
Conclude that, if torsion vanishes as general relativity assumes, Sµκλ = 0, then
DκAλ −DλAκ =∂Aλ
∂xκ− ∂Aκ
∂xλ. (2.53)
⋄
2.6.10 Connection coefficients in terms of the metric
The connection coefficients have been defined, equation (2.34), as derivatives of the tangent basis vectors gµ.
However, the connection coefficients can be expressed purely in terms of the (first derivatives of the) metric,
without reference to the individual basis vectors. The partial derivatives of the metric are
∂gλµ
∂xν=∂gλ · gµ
∂xν
= gλ ·∂gµ
∂xν+ gµ ·
∂gλ
∂xν
= gλ · gκ Γκµν + gµ · gκ Γκ
λν
= gλκ Γκµν + gµκ Γκ
λν
= Γλµν + Γµλν , (2.54)
which is a sum of two connection coefficients. Here Γλµν with all indices lowered is defined to be Γκµν with
the first index lowered by the metric,
Γλµν ≡ gλκΓκµν . (2.55)
Combining the metric derivatives in the following fashion yields an expression for a single connection:
∂gλµ
∂xν+∂gλν
∂xµ− ∂gµν
∂xλ= Γλµν + Γµλν + Γλνµ + Γνλµ − Γµνλ − Γνµλ
= 2 Γλµν − Sλµν − Sµνλ − Sνµλ
= 2 Γλµν , (2.56)
the last line of which follows from the no-torsion condition Sλµν = 0. Thus
Γλµν =1
2
(
∂gλµ
∂xν+∂gλν
∂xµ− ∂gµν
∂xλ
)
is not a coordinate tensor . (2.57)
This is the formula that allows connection coefficients to be calculated from the metric.
2.7 Coordinate 4-velocity 51
2.6.11 Mathematical aside
General relativity is a metric theory. Many of the structures introduced above can be defined mathematically
without a metric. For example, it is possible to define the tangent space of vectors with basis gµ, and to
define a dual vector space with basis gµ such that gµ · gν = δνµ. Elements of the dual vector space are
commonly called one-forms. Similarly it is possible to define connections and covariant derivatives without
a metric. However, this book follows general relativity in assuming that spacetime has a metric.
2.7 Coordinate 4-velocity
Consider a particle following a worldline
xµ(τ) , (2.58)
where τ is the particle’s proper time. The proper time along any interval of the worldline is dτ ≡√−ds2.
Define the coordinate 4-velocity uµ by
uµ ≡ dxµ
dτis a coordinate 4-vector . (2.59)
The magnitude squared of the 4-velocity is constant
uµuµ = gµν
dxµ
dτ
dxν
dτ=ds2
dτ2= −1 . (2.60)
The negative sign arises from the choice of metric signature: with the signature −+++ adopted here, there
is a − sign between ds2 and dτ2. Equation (2.60) can be regarded as an integral of motion associated with
conservation of particle rest mass.
2.8 Geodesic equation
Let u ≡ gµuµ be the 4-velocity in coordinate-independent notation. The principle of equivalence implies
that the geodesic equation, the equation of motion of a freely-falling particle, is
du
dτ= 0 . (2.61)
Why? Because du/dτ = 0 in the particle’s own free-fall frame, and the equation is coordinate-independent.
In the particle’s own free-fall frame, the particle’s 4-velocity is uµ = 1, 0, 0, 0, and the particle’s locally
inertial axes gµ = g0, g1, g2, g3 are constant.
52 Fundamentals of General Relativity
What does the equation of motion look like in coordinate notation? The acceleration is
du
dτ=dxν
dτ
∂u
∂xν
= uνgκDνuκ
= uνgκ
(
∂uκ
∂xν+ Γκ
µνuµ
)
= gκ
(
duκ
dτ+ Γκ
µνuµuν
)
. (2.62)
The geodesic equation is then
duκ
dτ+ Γκ
µνuµuν = 0 . (2.63)
Another way of writing the geodesic equation is
Duκ
Dτ= 0 , (2.64)
where D/Dτ is the covariant proper time derivative
D
Dτ≡ uνDν . (2.65)
2.9 Coordinate 4-momentum
The coordinate 4-momentum of a particle of rest mass m is defined to be
pµ ≡ muµ = mdxµ
dτis a coordinate 4-vector . (2.66)
The momentum squared is
pµpµ = m2uµu
µ = −m2 (2.67)
minus the square of the rest mass. Again, the minus sign arises from the choice −+++ of metric signature.
2.10 Affine parameter
For photons, the rest mass is zero, m = 0, but the 4-momentum pµ remains finite. Define the affine
parameter λ by
λ ≡ τ
mis a coordinate scalar (2.68)
2.11 Affine distance 53
which remains finite in the limit m → 0. The affine parameter λ is unique up to an overall linear transfor-
mation (that is, αλ + β is also an affine parameter, for constant α and β), because of the freedom in the
choice of mass m and the zero point of proper time τ . In terms of the affine parameter, the 4-momentum is
pµ =dxµ
dλ. (2.69)
The geodesic equation is then in coordinate-independent notation
dp
dλ= 0 , (2.70)
or in component form
dpκ
dλ+ Γκ
µνpµpν = 0 , (2.71)
which works for massless as well as massive particles.
Another way of writing this is
Dpκ
Dλ= 0 , (2.72)
where D/Dλ is the covariant affine derivative
D
Dλ≡ pνDν . (2.73)
2.11 Affine distance
The affine parameter is also called the affine distance, because it provides a measure of distance along null
geodesics. When you look at a scene with your eyes, you are looking along null geodesics, and the natural
measure of distance to objects that you see is the affine distance. The freedom in the overall scaling of the
affine distance is fixed by setting it equal to the proper distance near the observer in the observer’s locally
inertial rest frame.
In special relativity, the affine distance coincides with the perceived (e.g. binocular) distance to objects.
2.12 Riemann curvature tensor
The Riemann curvature tensor Rκλµν is defined by the commutator of the covariant derivative acting
on a 4-vector
[Dκ, Dλ]Aµ = RκλµνAν is a coordinate tensor . (2.74)
The expression (2.74) assumes vanishing torsion; the more general expression with non-zero torsion is (3.20).
54 Fundamentals of General Relativity
The expression (2.39) for the covariant derivative yields the following formula for the Riemann tensor in
terms of connection coefficients
Rκλµν =∂Γµνλ
∂xκ− ∂Γµνκ
∂xλ+ Γα
µλΓανκ − ΓαµκΓανλ is a coordinate tensor . (2.75)
This is the formula that allows the Riemann tensor to be calculated from the connection coefficients.
In flat (Minkowski) space, covariant derivatives reduce to partial derivatives, Dκ → ∂/∂xκ, and
[Dκ, Dλ]→[
∂
∂xκ,∂
∂xλ
]
= 0 in flat space (2.76)
so that Rκλµν = 0 in flat space.
Comment: In quantum field theories (QED, QCD), the commutator of the gauge-covariant derivative is
taken to be the field. In conventional general relativity, by contrast, the metric is taken to be the fundamental
field, rather than the curvature. Another difference between quantum field theories and general relativity is
that the Lagrangian of quantum field theories is taken to be quadratic in the field, whereas the Lagrangian
of general relativity is taken to be linear in the curvature (specifically, the general relativity Lagrangian is
the Ricci scalar R).
2.13 Symmetries of the Riemann tensor
In a locally inertial frame, the connection coefficients all vanish, Γλµν = 0, but their partial derivatives,
which are proportional to second derivatives of the metric tensor, do not vanish. Thus in a locally inertial
frame the Riemann tensor is
Rκλµν =∂Γµνλ
∂xκ− ∂Γµνκ
∂xλ
=1
2
(
∂2gµν
∂xκ∂xλ+
∂2gµλ
∂xκ∂xν− ∂2gνλ
∂xκ∂xµ− ∂2gµν
∂xλ∂xκ− ∂2gµκ
∂xλ∂xν+
∂2gνκ
∂xλ∂xµ
)
=1
2
(
∂2gµλ
∂xκ∂xν− ∂2gνλ
∂xκ∂xµ− ∂2gµκ
∂xλ∂xν+
∂2gνκ
∂xλ∂xµ
)
. (2.77)
You can check that the bottom line of equation (2.77):
1. is antisymmetric in κ↔ λ,
2. is antisymmetric in µ↔ ν,
3. is symmetric in κλ↔ µν,
4. has the property that the sum of the cyclic permutations of the last three indices vanishes
Rκλµν +Rκνλµ +Rκµνλ = 0 . (2.78)
The first three of these four symmetries can be summarized by the shorthand notation
R([κλ][µν]) (2.79)
2.14 Ricci tensor, Ricci scalar 55
in which [ ] denotes anti-symmetrization and ( ) symmetrization. These symmetries imply that the Riemann
tensor is a symmetric matrix of antisymmetric matrices. An antisymmetric matrix has 6 degrees of freedom.
A symmetric matrix of these things is a 6×6 symmetric matrix, which has 21 degrees of freedom. The final,
cyclic symmetry of the Riemann tensor, equation (2.78), removes 1 degree of freedom. Thus the Riemann
tensor has a net 20 degrees of freedom.
Although the above symmetries were derived in a locally inertial frame, the fact that the Riemann tensor
is a tensor means that the symmetries hold in any frame. If you prefer, you can add back the products of
connection coefficients in equation (2.75), and check that the claimed symmetries remain.
2.14 Ricci tensor, Ricci scalar
The Ricci tensor Rκµ and Ricci scalar R are the essentially unique contractions of the Riemann curvature
tensor. The Ricci tensor, the compressive part of the Riemann tensor, is
Rκµ ≡ gλνRκλµν is a coordinate tensor . (2.80)
The symmetries of the Riemann tensor imply that the Ricci tensor is symmetric
Rκµ = Rµκ (2.81)
and therefore has 10 independent components.
The Ricci scalar is
R ≡ gκµRκµ is a coordinate tensor (a scalar) . (2.82)
2.15 Einstein tensor
The Einstein tensor Gκµ is defined by
Gκµ ≡ Rκµ − 12 gκµR is a coordinate tensor . (2.83)
The symmetry of the Ricci and metric tensors imply that the Einstein tensor is likewise symmetric
Gκµ = Gµκ . (2.84)
The Einstein tensor has 10 independent components.
2.16 Bianchi identities
The Jacobi identity
[Dκ, [Dλ, Dµ]] + [Dλ, [Dµ, Dκ]] + [Dµ, [Dκ, Dλ]] = 0 (2.85)
56 Fundamentals of General Relativity
implies the Bianchi identities
DκRλµνπ +DλRµκνπ +DµRκλνπ = 0 (2.86)
which can be written in shorthand
D[κRλµ]νπ = 0 . (2.87)
The Bianchi identities constitute a set of differential relations between the components of the Riemann
tensor, which are distinct from the algebraic symmetries of the Riemann tensor.
There are 20 independent Bianchi identities. If just the symmetries (2.79) of the Riemann tensor are
taken into account, then there are 24 identities; but the cyclic symmetry (2.78) eliminates 4, leaving 20
independent identities.
2.17 Covariant conservation of the Einstein tensor
The most important consequence of the Bianchi identities (2.87) is obtained from the double contraction
gκνgλπ (DκRλµνπ +DλRµκνπ +DµRκλνπ) = −DκRκµ −DλRλµ +DµR = 0 (2.88)
which implies that
DκGκµ = 0 . (2.89)
This equation is a primary motivation for the form of the Einstein equations, since it implies energy-
momentum conservation, equation (2.91).
2.18 Einstein equations
Einstein’s equations are
Gκµ = 8πGTκµ is a coordinate tensor equation . (2.90)
What motivates the form of Einstein’s equations?
1. The equation is generally covariant;
2. The Bianchi identities guarantee conservation of energy-momentum;
3. The Einstein tensor depends on the lowest (second) order derivatives of the metric tensor that do not
vanish in a locally inertial frame;
The covariant conservation of the Einstein tensor, equation (2.89), implies the conservation of energy-
momentum
DκTκµ = 0 . (2.91)
Einstein’s equations (2.90) constitute a complete set of gravitational equations, generalizing Poisson’s
2.19 Summary of the path from metric to the energy-momentum tensor 57
equation of Newtonian gravity. However, Einstein’s equations by themselves do not constitute a closed set
of equations: in general, other equations, such as Maxwell’s equations of electromagnetism, and equations
describing the microphysics of the energy-momentum, must be adjoined to form a closed set.
2.19 Summary of the path from metric to the energy-momentum tensor
1. Start by defining the metric gµν .
2. Compute the connection coefficients Γλµν from equation (2.57).
3. Compute the Riemann tensor Rκλµν from equation (2.75).
4. Compute the Ricci tensor Rκµ from equation (2.80), the Ricci scalar R from equation (2.82), and the
Einstein tensor Gκµ from equation (2.83).
5. The Einstein equations (2.90) then imply the energy-momentum tensor Tκµ.
The path from metric to energy-momentum tensor is straightforward to program on a computer, but the
results are typically messy and complicated, even for fairly simple spacetimes. Inverting the path to recover
the metric from a given energy-momentum content is typically highly non-trivial, the subject of a huge
literature.
The great majority of metrics gµν yield an energy-momentum tensor Tκµ that cannot be achieved with
normal matter.
2.20 Energy-momentum tensor of an ideal fluid
The simplest non-trival energy-momentum tensor is that of an ideal fluid. In this case T µν is isotropic in
the locally inertial rest frame of the fluid, taking the form
T µν =
ρ 0 0 0
0 p 0 0
0 0 p 0
0 0 0 p
(2.92)
where
ρ is the proper mass-energy density ,
p is the proper pressure .(2.93)
The expression (2.92) is valid only in the locally inertial rest frame of the fluid. An expression that is valid
in any frame is
T µν = (ρ+ p)uµuν + p gµν , (2.94)
where uµ is the 4-velocity of the fluid. Equation (2.94) is valid because it is a tensor equation, and it is true
in the locally inertial rest frame, where uµ = 1, 0, 0, 0.
58 Fundamentals of General Relativity
2.21 Newtonian limit
The Newtonian limit is obtained in the limit of a weak gravitational field and non-relativistic (pressureless)
matter. In Cartesian coordinates, the metric in the Newtonian limit is
ds2 = − (1 + 2Φ)dt2 + (1− 2Φ)(dx2 + dy2 + dz2) , (2.95)
in which
Φ(x, y, z) = Newtonian potential (2.96)
is a function only of the spatial coordinates x, y, z, not of time t.
For this metric, to first order in the potential Φ the only non-vanishing component of the Einstein tensor
is the time-time component
Gtt = 2∇2Φ , (2.97)
where ∇2 = ∂2/∂x2 + ∂2/∂y2 + ∂2/∂z2 is the usual 3D Laplacian operator. This component (2.97) of the
Einstein tensor plugged into Einstein’s equations (2.90) implies Poisson’s equation (2.4).
3
∗More on the coordinate approach
3.1 Weyl tensor
The trace-free, or tidal, part of the Riemann curvature tensor defines the Weyl tensor Cκλµν
Cκλµν ≡ Rκλµν − 12 (gκµRλν − gκνRλµ + gλνRκµ − gλµRκν) + 1
6 (gκµgλν − gκνgλµ)R is a coordinate tensor .
(3.1)
The Weyl tensor has 10 independent components. These 10 components together with the 10 components
of the Ricci tensor account for the 20 components of the Riemann tensor. The Weyl tensor inherits the
symmetries (2.79) of the Riemann tensor
C([κλ][µν]) . (3.2)
Whereas the Einstein tensor Gκµ, necessarily vanishes in a region of spacetime where there is no energy-
momentum, Tκµ = 0, the Weyl tensor does not. The Weyl tensor expresses the presence of tidal gravitational
forces, and of gravitational waves.
3.2 Evolution equations for the Weyl tensor
This section is included because (a) the comparison to Maxwell’s equations is neat and insightful, (b) it helps
to account for the degrees of freedom of the gravitational field, (c) it shows how the Weyl tensor encodes
gravitational waves.
Contracted on one index, the Bianchi identities (2.86) are
D[κRλµ]νκ = 0 . (3.3)
There are 20 such independent contracted identities. Since this is the same as the number of independent
Bianchi identities, it follows that the contracted Bianchi identities (3.3) are equivalent to the full set of
Bianchi identities (2.87). If the Riemann tensor is separated into its trace (Ricci) and traceless (Weyl) parts,
60 ∗More on the coordinate approach
equation (3.1), then the contracted Bianchi identities (3.3) become the Weyl evolution equations
DκCκλµν = Jλµν , (3.4)
where Jλµν is the Weyl current
Jλµν ≡ 12 (DµGλν −DνGλµ)− 1
6 (gλνDµG− gλµDνG) . (3.5)
The Weyl evolution equations (3.4) can be regarded as the gravitational analogue of Maxwell’s equations of
electromagnetism.
The Weyl current Jλµν is a vector of bivectors, which would suggest that it has 4 × 6 = 24 components,
but it loses 4 of those components because of the cyclic identity (2.78), which implies the cyclic symmetry
J[λµν] = 0 . (3.6)
Thus Jλµν has 20 independent components, in agreement with the above assertion that there are 20 inde-
pendent contracted Bianchi identities. Since the Weyl tensor is traceless, contracting the Weyl evolution
equations (3.4) on λν yields zero on the left hand side, so that the contracted Weyl current satisfies
Jλλµ = 0 . (3.7)
This doubly-contracted Bianchi identity, which is the same as equation (2.89), enforces conservation of
energy-momentum. Unlike the cyclic symmetries (3.6), which are automatically satisfied, equations (3.7)
constitute a non-trivial set of 4 conditions on the Einstein tensor. Besides the algebraic relations (3.6) and
(3.7), the Weyl current satisfies 6 differential identities comprising the conservation law
DλJλµν = 0 (3.8)
in view of equation (3.4) and the antisymmetry of Cκλµν with respect to the indices κλ. The Weyl current
conservation law (3.8) follows automatically from the form (3.5) of the Weyl current, coupled with energy-
momentum conservation (2.89), so does not impose any additional non-trivial conditions on the Riemann
tensor.
The 4 relations (3.7) and the 6 identities (3.8) account for 10 of the 20 contracted Bianchi identities (3.3).
The remaining 10 equations comprise Maxwell-like equations (3.4) for the evolution of the 10 components
of the Weyl tensor.
Whereas the Einstein equations relating the Einstein tensor to the energy-momentum tensor are postulated
equations of general relativity, the 10 evolution equations for the Weyl tensor, and the 4 equations enforcing
covariant conservation of the Einstein tensor, follow mathematically from the Bianchi identities, and do not
represent additional assumptions of the theory.
Exercise 3.1 Confirm the counting of degrees of freedom. ⋄
Exercise 3.2 From the Bianchi identities, show that the Riemann tensor satisfies the covariant wave
equation
Rκλµν = DκDµRλν −DκDνRλµ +DλDνRκµ −DλDµRκν , (3.9)
3.3 Geodesic deviation 61
where is the D’Alembertian operator, the 4-dimensional wave operator
≡ DπDπ . (3.10)
Show that contracting equation (3.9) with gλν yields the identity Rκµ = Rκµ. Conclude that the wave
equation (3.9) is non-trivial only for the trace-free part of the Riemann tensor, the Weyl tensor Cκλµν . Show
that the wave equation for the Weyl tensor is
Cκλµν = (DκDµ − 12 gκµ )Rλν − (DκDν − 1
2 gκν )Rλµ
+ (DλDν − 12 gλν )Rκµ − (DλDµ − 1
2 gλµ )Rκν
+ 16 (gκµgλν − gκνgλµ)R . (3.11)
Conclude that in a vacuum, where Rκµ = 0,
Cκλµν = 0 . (3.12)
⋄
3.3 Geodesic deviation
This section on geodesic deviation is included not because the equation of geodesic deviation is crucial to
everyday calculations in general relativity, but rather for two reasons. First, the equation offers insight into
the physical meaning of the Riemann tensor. Second, the derivation of the equation offers a fine illustration
of the fact that in general relativity, whenever you take differences at infinitesimally separated points in
space or time, you should always take covariant differences.
Consider two objects that are free-falling along two infinitesimally separated geodesics. In flat space the
acceleration between the two objects would be zero, but in curved space the curvature induces a finite
acceleration between the two objects. This is how an observer can measure curvature, at least in principle:
set up an ensemble of objects initially at rest a small distance away from the observer in the observer’s
locally inertial frame, and watch how the objects begin to move. The equation (3.18) that describes this
acceleration between objects an infinitesimal distance apart is called the equation of geodesic deviation.
The covariant difference in the velocities of two objects an infinitesimal distance δxµ apart is
Dδxµ
Dτ= δuµ . (3.13)
In general relativity, the ordinary difference between vectors at two points a small interval apart is not
a physically meaningful thing, because the frames of reference at the two points are different. The only
physically meaningful difference is the covariant difference, which is the difference in the two vectors parallel-
transported across the gap between them. It is only this covariant difference that is independent of the frame
of reference. On the left hand side of equation (3.13), the proper time derivative must be the covariant proper
time derivative, D/Dτ = uλDλ. On the right hand side of equation (3.13), the difference in the 4-velocity
62 ∗More on the coordinate approach
at two points δxκ apart must be the covariant difference δ = δxκDκ. Thus equation (3.13) means explicitly
the covariant equation
uλDλδxµ = δxκDκu
µ . (3.14)
To derive the equation of geodesic deviation, first vary the geodesic equation Duµ/Dτ = 0 (I’ve put the
index µ downstairs so that the final equation (3.18) looks cosmetically better, but of course since everything
is covariant the µ index could just as well be put upstairs everywhere):
0 = δDuµ
Dτ
= δxκDκ
(
uλDλuµ
)
= δuλDλuµ + δxκuλDκDλuµ . (3.15)
On the second line, the covariant diffence δ between quantities a small distance δxκ apart has been set equal
to δxκDκ, while D/Dτ has been set equal to the covariant time derivative uλDλ along the geodesic. On the
last line, δxκDκuλ has been replaced by δuµ. Next, consider the covariant acceleration of the interval δxµ,
which is the covariant proper time derivative of the covariant velocity difference δuµ:
D2δxµ
Dτ2=Dδuµ
Dτ
= uλDλ (δxκDκuµ)
= δuκDκuµ + δxκuλDλDκuµ . (3.16)
As in the previous equation (3.15), on the second line D/Dτ has been set equal to uλDλ, while δ has been
set equal to δxκDκ. On the last line, uλDλδxκ has been set equal to δuµ, equation (3.14). Subtracting (3.15)
from (3.16) gives
D2δxµ
Dτ2= δxκuλ[Dλ, Dκ]uµ , (3.17)
or equivalently
D2δxµ
Dτ2+Rκλµνδx
κuλuν = 0 , (3.18)
which is the desired equation of geodesic deviation.
3.4 Commutator of the covariant derivative revisited
The commutator of the covariant derivative is of fundamental importance because it defines what is meant
by the field in gauge theories. It was seen above that the commutator of the covariant derivative acting
on a scalar defined the torsion tensor, equation (2.50), which general relativity assumes vanishes, while the
commutator of the covariant derivative acting on a vector defined the Riemann tensor, equation (2.74). Does
the commutator of the covariant derivative acting on a general tensor introduce any other distinct tensor? No:
3.4 Commutator of the covariant derivative revisited 63
the torsion and Riemann tensors completely define the action of the commutator of the covariant derivative
on any tensor, equation (3.22).
The general expression (3.22) for the commutator of the covariant derivative reveals the meaning of the
torsion and Riemann tensors. The torsion and Riemann tensors describe respectively the displacement and
the Lorentz transformation experienced by an object when parallel-transported around a curve. Displacement
and Lorentz transformations together constitute the Poincare group, the complete group of symmetries of
flat spacetime.
How can an object detect a displacement when parallel-transported around a curve? If you go around
a curve back to the same coordinate in spacetime where you began, won’t you necessarily be at the same
position? This is a question that goes to heart of the meaning of spacetime. To answer the question, you
have to consider how fundamental particles are able to detect position, orientation, and velocity. Classically,
particles may be structureless points, but quantum mechanically, particles possess frequency, wavelength,
spin, and (in the relativistic theory) boost, and presumably it is these properties that allow particles to
“measure” the properties of the spacetime in which they live. Specifically, a Dirac spinor (relativistic spin- 12
particle) has 8 degrees of freedom, of which 6 define a Lorentz transformation (a Lorentz rotor, a member of
the group of spin- 12 Lorentz transformations), and the remaining 2 comprise a complex number REALLY?
THE COMPLEX NUMBER IS WITH RESPECT TO THE PSEUDOSCALAR I ∼ e−ip·x whose phase
encodes the displacement. Thus a Dirac spinor could potentially detect a displacement through a change in
its phase when parallel-transported around a curve back to the same point in spacetime. General relativity,
which assumes that torsion vanishes, asserts that there is no such change of phase.
In the presence of torsion, the expression for the connection coefficients Γλµν is, from equation (2.57),
Γλµν =1
2
(
∂gλµ
∂xν+∂gλν
∂xµ− ∂gµν
∂xλ+ Sλµν + Sµνλ + Sνµλ
)
. (3.19)
The first part 12 (gλµ,ν + gλν,µ − gµν,λ) of this expression is called the Christoffel symbol of the first kind [the
same thing with the first index raised, 12g
κλ (gλµ,ν + gλν,µ − gµν,λ), is called the Christoffel symbol of the
second kind], while the second part 12 (Sλµν + Sµνλ + Sνµλ) is called the contortion (not contorsion!) tensor.
There’s no need to remember the crazy jargon, but in case you meet it, that’s what it means.
If torsion does not vanish, then the commutator of the covariant derivative acting on a contravariant
4-vector is
[Dκ, Dλ]Aµ = SνκλDνAµ +RκλµνA
ν is a coordinate tensor (3.20)
where the Riemann tensor Rκλµν is given in terms of the connection coefficients by the same formula (2.75)
as before, but the connection coefficients Γλµν themselves are given by (3.19). The Riemann tensor is still
antisymmetric in each of κ↔ λ and µ↔ ν, but with torsion it is no longer symmetric in κλ↔ µν. In other
words, the symmetries of the Riemann tensor with torsion are
R[κλ][µν] . (3.21)
As a matrix of antisymmetric tensors, the Riemann tensor with torsion has 6 × 6 = 36 degrees of freedom.
64 ∗More on the coordinate approach
Because the Riemann tensor Rκλµν is no longer symmetric in κλ↔ µν, the Ricci tensor Rκµ ≡ gλνRκλµν is
no longer symmetric, and likewise the Einstein tensor Gκµ ≡ Rκµ− 12Rgκµ is no longer symmetric. Evidently
the antisymmetric part of the Einstein tensor depends on torsion.
Acting on a general tensor, the commutator of the covariant derivative is
[Dκ, Dλ]Aπρ...µν... = Sα
κλDαAπρ...µν... +Rκλµ
αAπρ...αν... +Rκλν
αAπρ...µα... −Rκλα
πAαρ...µν... −Rκλα
ρAπα...µν... . (3.22)
In more abstract notation, the commutator of the covariant derivative is the operator
[Dκ, Dλ] = Sκλ ·D + Rκλ (3.23)
where Sκλ ≡ gµSµκλ and D ≡ gµDµ, and the Riemann curvature operator Rκλ is an operator whose action
on any tensor is specified by equation (3.22). The action of the operator Rκλ is analogous to that of the
covariant derivative (2.40): there’s a positive R term for each covariant index, and a negative R term for each
contravariant index. The action of Rκλ on a scalar is zero, which reflects the fact that a scalar is unchanged
by a Lorentz transformation.
4
∗Action principle
Hamilton’s principle of least action postulates that any dynamical system is characterized by a scalar action
S, which has the property that when the system evolves from one specified state to another, the path by
which it gets between the two states is such as to minimize the action. The action need not be a global
minimum, just a local minimum with respect to small variations in the path between fixed initial and final
states.
That nature appears to respect the principle of least action is of the profoundest significance.
x0
x1
λ
Figure 4.1 The action principle considers various paths through spacetime between fixed initial and finalconditions, and chooses that path that minimizes the action.
66 ∗Action principle
4.1 Principle of least action for point particles
The path of a point particle through spacetime is specified by its coordinates xµ(λ) as a function of some
arbitrary parameter λ. In non-relativistic mechanics it is usual to take the parameter λ to be the time t,
and the path of a particle through space is then specified by three spatial coordinates xi(t). In relativity
however it is more natural to treat the time and space coordinates on an equal footing, and regard the path
of a particle as being specified by four spacetime coordinates xµ(λ) as a function of an arbitrary parameter
λ. The parameter λ is simply a continuous parameter that labels points along the path, and has no physical
significance (for example, it is not necessarily an affine parameter).
The path of a system of N point particles through spacetime is specified by 4N coordinates xµ(λ). The
action principle postulates that, for a system of N point particles, the action S is an integral of a Lagrangian
L(xµ, dxµ/dλ) which is a function of the 4N coordinates xµ(λ) together with the 4N velocities dxµ/dλ with
respect to the arbitrary parameter λ. The action from an initial state at λi to a final state at λf is thus
S =
∫ λf
λi
L
(
xµ,dxµ
dλ
)
dλ . (4.1)
The principle of least action demands that the actual path taken by the system between given initial and
final coordinates xµi and xµ
f is such as to minimize the action. Thus the variation δS of the action must be
zero under any change δxµ in the path, subject to the constraint that the coordinates at the endpoints are
fixed, δxµi = 0 and δxµ
f = 0,
δS =
∫ λf
λi
(
∂L
∂xµδxµ +
∂L
∂(dxµ/dλ)δ(dxµ/dλ)
)
dλ = 0 . (4.2)
The change in the velocity along the path is just the velocity of the change, δ(dxµ/dλ) = d(δxµ)/dλ.
Integrating the second term in the integrand of equation (4.2) by parts yields
δS =
[
∂L
∂(dxµ/dλ)δxµ
]λf
λi
+
∫ λf
λi
(
∂L
∂xµ− d
dλ
∂L
∂(dxµ/dλ)
)
δxµ dλ = 0 . (4.3)
The surface term in equation (4.3) vanishes, since by hypothesis the coordinates are held fixed at the end
points, so δxµ = 0 at the end points. Therefore the integral in equation (4.3) must vanish. Indeed least
action requires the integral to vanish for all possible variations δxµ in the path. The only way this can
happen is that the integrand must be identically zero. The result is the Euler-Lagrange equations of
motion
d
dλ
∂L
∂(dxµ/dλ)− ∂L
∂xµ= 0 . (4.4)
It might seem that the Euler-Lagrange equations (4.4) are inadequately specified, since they depend on
some arbitrary unknown parameter λ. But in fact the Euler-Lagrange equations are the same regardless of
the choice of λ. An example of the irrelevance of λ will be seen in the next section, §4.2. Since λ can be
chosen arbitrarily, it is usual to choose it in some convenient fashion. For a massive particle, λ can be taken
4.2 Action for a test particle 67
equal to the proper time τ of the particle. For a massless particle, whose proper time never progresses, λ
can be taken equal to an affine parameter.
4.2 Action for a test particle
According to the principle of equivalence, a test particle in a gravitating system moves along a geodesic, a
straight line relative to local free-falling frames. A geodesic is the shortest distance between two points. In
relativity this translates, for a massive particle, into the longest proper time between two points. The proper
time along any path is dτ =√−ds2 =
√
−gµνdxµdxν . Thus the action Sm of a test particle of rest mass m
in a gravitating system is
Sm = −m∫ λf
λi
dτ = −m∫ λf
λi
√
−gµνdxµ
dλ
dxν
dλdλ . (4.5)
The factor of rest massm brings the action, which has units of angular momentum, to standard normalization.
The overall minus sign comes from the fact that the action is a minimum whereas the proper time is a
maximum along the path. The action principle requires that the Lagrangian be written as a function
of the coordinates xµ and velocities dxν/dλ, and it is seen that the integrand in the last expression of
equation (4.5) has the desired form, the metric gµν being considered a given function of the coordinates.
Thus the Lagrangian Lm of a test particle of mass m is
Lm = −m√
−gµνdxµ
dλ
dxν
dλ. (4.6)
The partial derivatives that go in the Euler-Lagrange equations (4.4) are then
∂Lm
∂(dxκ/dλ)= −m
−gκνdxν
dλ√
−gπρ(dxπ/dλ)(dxρ/dλ), (4.7a)
∂Lm
∂xκ= −m
−1
2
∂gµν
dxκ
dxµ
dλ
dxν
dλ√
−gπρ(dxπ/dλ)(dxρ/dλ). (4.7b)
The denominators in the expressions (4.7) for the partial derivatives of the Lagrangian are√
−gπρ(dxπ/dλ)(dxρ/dλ) = dτ/dλ. It was not legitimate to make this substitution before taking the partial
derivatives, since the Euler-Lagrange equations require that the Lagrangian be expressed in terms of xµ and
dxµ/dλ, but it is fine to make the substitution now that the partial derivatives have been obtained. The
partial derivatives (4.7) thus simplify to
∂Lm
∂(dxκ/dλ)= mgκν
dxν
dλ
dλ
dτ= muκ , (4.8)
∂Lm
∂xκ=
1
2m∂gµν
dxκ
dxµ
dλ
dxν
dλ
dλ
dτ= mΓµνκu
µuν dτ
dλ, (4.9)
68 ∗Action principle
in which uκ ≡ dxκ/dτ is the usual 4-velocity, and the derivative of the metric has replaced by connections
in accordance with equation (2.54). The resulting Euler-Lagrange equations of motion (4.4) are
dmuκ
dλ= mΓµνκu
µuν dτ
dλ. (4.10)
As remarked in §4.1, the choice of the arbitrary parameter λ has no effect on the equations of motion.
With a factor of mdτ/dλ cancelled, and with the derivative converted to a covariant derivative by (2.39),
equation (4.10) becomes
Duκ
Dτ= Sµνκu
µuν , (4.11)
where Sµνκ is the torsion, equation (2.51). If torsion vanishes, as general relativity assumes, then the result
is the usual equation of geodesic motion
Duκ
Dτ= 0 . (4.12)
The fact that motion is geodesic only if torsion vanishes is to be expected, since, as argued in §2.6.7, space
is locally inertial only if torsion vanishes.
4.3 Action for a charged test particle in an electromagnetic field
Aim is to reproduce the Lorentz force law.
S = Sm + Sq (4.13)
where
Sq = q
∫ λf
λi
Aµ dxµ = q
∫ λf
λi
Aµdxµ
dλdλ . (4.14)
The Lagrangian is therefore
Lq = qAµdxµ
dλ. (4.15)
Partial derivatives are
∂Lq
∂(dxκ/dλ)= qAκ ,
∂Lq
∂xκ= q
∂Aµ
∂xκ
dxµ
dλ= q
∂Aµ
∂xκuµ dτ
dλ. (4.16a)
Applied to the Lagrangian L = Lm + Lq, the Euler-Lagrange equations (4.4) are
d
dλ(muκ + qAκ) =
(
mΓµνκuµuν + q
∂Aµ
∂xκuµ
)
dτ
dλ. (4.17)
If torsion vanishes, as general relativity assumes, then the result is the Lorentz force law for a test particle
of mass m and charge q moving in a prescribed gravitational and electromagnetic field
Dmuκ
Dτ= qFµκu
µ . (4.18)
4.4 Generalized momentum 69
Exercise 4.1 Show that if torsion does not vanish, then the Lorentz force law becomes
Dmuκ
Dτ= [qFµκ + Sνµκ (muν + qAν)]uµ . (4.19)
[Hint: Recall the relations (2.51) and (2.52).] ⋄
4.4 Generalized momentum
The generalized momentum
πκ ≡∂L
∂(dxκ/dλ). (4.20)
The generalized momentum πκ of a test particle coincides with the ordinary momentum pκ:
πκ = pκ = muκ . (4.21)
The generalized momentum of a test particle of charge q in an electromagnetic field of potential Aκ
πκ = pκ + qAκ . (4.22)
4.5 Hamiltonian
Work with coordinates and generalized momenta instead of coordinates and velocities. Define Hamiltonian
H by
H ≡ πµdxµ
dλ− L (4.23)
S =
∫(
πµdxµ
dλ−H
)
dλ (4.24)
δS = [πµδxµ] +
∫
−(
dπµ
dλ+∂H
∂xµ
)
δxµ +
(
dxµ
dλ− ∂H
∂πµ
)
δπµ
dλ (4.25)
dπµ
dλ= − ∂H
∂xµ,
dxµ
dλ=∂H
∂πµ. (4.26)
What is this useful for?
70 ∗Action principle
4.6 Derivatives of the action
Besides being a scalar whose minimum value between fixed endpoints defines the path between those points,
the action can also be treated as a function of its endpoints along the actual path. Along the actual path,
the equations of motion are satisfied, so the integral in the variation (4.3) of the action vanishes identically.
The surface term in the variation (4.3) then implies that δS = πµδxµ, so that the partial derivatives of the
action with respect to the coordinates are equal to the generalized momenta,
∂S
∂xµ= πµ . (4.27)
This is the basis of the Hamilton-Jacobi method.
PART THREE
IDEAL BLACK HOLES
Concept Questions
1. What evidence do astronomers currently accept as indicating the presence of a black hole in a system?
2. Why can astronomers measure the masses of supermassive black holes only in relatively nearby galaxies?
3. To what extent (with what accuracy) are real black holes in our Universe described by the no-hair theorem?
4. Does the no-hair theorem apply inside a black hole?
5. Black holes lose their hair on a light-crossing time. How long is a light-crossing time for a typical stellar-
sized or supermassive astronomical black hole?
6. Relativists say that the metric is gµν , but they also say that the metric is ds2 = gµν dxµdxν . How can
both statements be correct?
7. The Schwarzschild geometry is said to describe the geometry of spacetime outside the surface of the Sun
or Earth. But the Schwarzschild geometry supposedly describes non-rotating masses, whereas the Sun
and Earth are rotating. If the Sun or Earth collapsed to a black hole conserving their mass M and angular
momentum L, roughly what would the spin a/M = L/M2 of the black hole be relative to the maximal
spin a/M = 1 of a Kerr black hole?
8. What happens at the horizon of a black hole?
9. As cold matter becomes denser, it goes through the stages of being solid/liquid like a planet, then electron
degenerate like a white dwarf, then neutron degenerate like a neutron star, then finally it collapses to a
black hole. Why could there not be a denser state of matter, denser than a neutron star, that brings a
star to rest inside its horizon?
10. How can an observer determine whether they are “at rest” in the Schwarzschild geometry?
11. An observer outside the horizon of a black hole never sees anything pass through the horizon, even to the
end of the Universe. Does the black hole then ever actually collapse, if no one ever sees it do so?
12. If nothing can ever get out of a black hole, how does its gravity get out?
13. Why did Einstein believe that black holes could not exist in nature?
14. In what sense is a rotating black hole “stationary”, but not “static”?
15. What is a white hole? Do they exist?
16. Could the expanding Universe be a white hole?
17. Could the Universe be the interior of a black hole?
18. You know the Schwarzschild metric for a black hole. What is the corresponding metric for a white hole?
74 Concept Questions
19. What is the best kind of black hole to fall into if you want to avoid being tidally torn apart?
20. Why do astronomers often assume that the inner edge of an accretion disk around a black hole occurs at
the innermost stable orbit?
21. A collapsing star of uniform density has the geometry of a collapsing Friedmann-Robertson-Walker cos-
mology. If a spatially flat FRW cosmology corresponds to a star that starts from zero velocity at infinity,
then to what do open or closed FRW cosmologies correspond?
22. Is the singularity of a Reissner-Nordstrom black hole gravitationally attractive or repulsive?
23. If you are a charged particle, which dominates near the singularity of the Reissner-Nordstrom geometry,
the electrical attraction/repulsion or the gravitational attraction/repulsion?
24. Is a white hole gravitationally attractive or repulsive?
25. What happens if you fall into a white hole?
26. Which way does time go in Parallel Universes in the Reissner-Nordstrom geometry?
27. What does it mean that geodesics inside a black hole can have negative energy?
28. Can geodesics have negative energy outside a black hole? How about inside the ergosphere?
29. Physically, what causes mass inflation?
30. Is mass inflation likely to occur inside real astronomical black holes?
31. What happens at the X point, where the ingoing and outgoing inner horizons of the Reissner-Nordstrom
geometry intersect?
32. Can a particle like an electron or proton, whose charge far exceeds its mass (in geometric units), be
modeled as Reissner-Nordstrom black hole?
33. Does it makes sense that a person might be at rest in the Kerr-Newman geometry? How would the
Boyer-Linquist coordinates of such a person vary along their worldline?
34. In identifying M as the mass and a the angular momentum per unit mass of the black hole in the Boyer-
Linquist metric, why is it sufficient to consider the behaviour of the metric at r →∞?
35. Does space move faster than light inside the ergosphere?
36. If space moves faster than light inside the ergosphere, why is the outer boundary of the ergosphere not a
horizon?
37. Do closed timelike curves make sense?
38. What does Carter’s fourth integral of motion Q signify physically?
39. What is special about a principal null congruence?
40. Evaluated in the locally inertial frame of a principal null congruence, the spin-0 component of the Weyl
scalar of the Kerr geometry is C = −M/(r− ia cosθ)3, which looks like the Weyl scalar C = −M/r3 of the
Schwarzschild geometry but with radius r replaced by the complex radius r− ia cos θ. Is there something
deep here? Can the Kerr geometry be constructed from the Schwarzschild geometry by complexifying the
radial coordinate r?
What’s important?
1. Astronomical evidence suggests that stellar-sized and supermassive black holes exist ubiquitously in nature.
2. The no-hair theorem, and when and why it applies.
3. The physical picture of black holes as regions of spacetime where space is falling faster than light.
4. A physical understanding of how the metric of a black hole relates to its physical properties.
5. Penrose (conformal) diagrams. In particular, the Penrose diagrams of the various kinds of vacuum black
hole: Schwarzschild, Reissner-Nordstrom, Kerr-Newman.
6. What really happens inside black holes. Collapse of a star. Mass inflation instability.
5
Observational Evidence for Black Holes
It is beyond the scope of this course to discuss the observational evidence for black holes in any detail.
However, it is useful to summarize a few facts.
1. Observational evidence supports the idea that black holes occur ubiquitously in nature. They are not
observed directly, but reveal themselves through their effects on their surroundings. Two kinds of black
hole are observed: stellar-sized black holes in x-ray binary systems, mostly in our own Milky Way galaxy,
and supermassive black holes in Active Galactic Nuclei (AGN) found at the centers of our own and other
galaxies.
2. The primary evidence that astronomers accept as indicating the presence of a black hole is a lot of mass
compacted into a tiny space.
a. In an x-ray binary system, if the mass of the compact object exceeds 3 M⊙, the maximum theoretical
mass of a neutron star, then the object is considered to be a black hole. Many hundreds of x-ray binary
systems are known in our Milky Way galaxy, but only 10s of these have measured masses, and in about
20 the measured mass indicates a black hole.
b. Several tens of thousands of AGN have been cataloged, identified either in the radio, optical, or x-rays.
But only in nearby galaxies can the mass of a supermassive black hole be measured directly. This is
because it is only in nearby galaxies that the velocities of gas or stars can be measured sufficiently close
to the nuclear center to distinguish a regime where the velocity becomes constant, so that the mass
can be attribute to an unresolved central point as opposed to a continuous distribution of stars. The
masses of about 40 supermassive black holes have been measured in this way. The masses range from
the 4 × 106 M⊙ mass of the black hole at the center of the Milky Way to the 3 × 109 M⊙ mass of the
black hole at the center of the M87 galaxy at the center of the Virgo cluster at the center of the Local
Supercluster of galaxies.
3. Secondary evidences for the presence of a black hole are:
a. high luminosity;
b. non-stellar spectrum, extending from radio to gamma-rays;
c. rapid variability.
d. relativistic jets.
Observational Evidence for Black Holes 77
Jets in AGN are often one-sided, and a few that are bright enough to be resolved at high angular resolution
show superluminal motion. Both evidences indicate that jets are commonly relativistic, moving at close
to the speed of light. There are a few cases of jets in x-ray binary systems.
4. Stellar-sized black holes are thought to be created in supernovae as the result of the core-collapse of
stars more massive than about 25 M⊙ (this number depends in part on uncertain computer simulations).
Supermassive black holes are probably created initially in the same way, but they then grow by accretion of
gas funnelled to the center of the galaxy. The growth rates inferred from AGN luminosities are consistent
with this picture.
5. Long gamma-ray bursts (lasting more than about 2 seconds) are associated observationally with super-
novae. It is thought that in such bursts we are seeing the formation of a black hole. As the black hole gulps
down the huge quantity of material needed to make it, it regurgitates a relativistic jet that punches through
the envelope of the star. If the jet happens to be pointed in our direction, then we see it relativistically
beamed as a gamma-ray burst.
6. Astronomical black holes present the only realistic prospect for testing general relativity in the strong field
regime, since such fields cannot be reproduced in the laboratory. At the present time the observational
tests of general relativity from astronomical black holes are at best tentative. One test is the redshifting
of 7 keV iron lines in a small number of AGN, notably MCG-6-30-15, which can be interpreted as being
emitted by matter falling on to a rotating (Kerr) black hole.
7. At present, no gravitational waves have been definitely detected from anything. In the future, gravitational
wave astronomy should eventually detect the merger of two black holes. If the waveforms of merging black
holes are consistent with the predictions of general relativity, it will provide a far more stringent test of
strong field general relativity than has been possible to date.
8. Although gravitational waves have yet to be detected directly, their existence has been inferred from the
gradual speeding up of the orbit of the Hulse-Taylor binary, which consists of two neutron stars, one of
which, PSR1913+16, is a pulsar. The parameters of the orbit have been measured with exquisite precision,
and the rate of orbital speed-up is in good agreement with the energy loss by quadrupole gravitational
wave emission predicted by general relativity.
6
Ideal Black Holes
6.1 Definition of a black hole
What is a black hole? Doubtless you have heard the standard definition many times: It is a region whose
gravity is so strong that not even light can escape.
But why can light not escape from a black hole? A standard answer, which John Michell (1784, Phil.
Trans. Roy. Soc. London 74, 35) would have found familiar, is that the escape velocity exceeds the speed of
light. But that answer brings to mind a Newtonian picture of light going up, turning around, and coming
back down, that is altogether different from what general relativity actually predicts.
A better definition of a black hole is that it is a
region where space is falling faster than light.
Inside the horizon, light emitted outwards is carried inward by the faster-than-light inflow of space, like a
fish trying but failing to swim up a waterfall.
The definition may seem jarring. If space has no substance, how can it fall faster than light? It means
that inside the horizon any locally inertial frame is compelled to fall to smaller radius as its proper time goes
by. This fundamental fact is true regardless of the choice of coordinates.
A similar concept of space moving arises in cosmology. Astronomers observe that the Universe is expand-
ing. Cosmologists find it convenient to conceptualize the expansion by saying that space itself is expanding.
For example, the picture that space expands makes it more straightforward, both conceptually and math-
ematically, to deal with regions of spacetime beyond the horizon, the surface of infinite redshift, of an
observer.
6.2 Ideal black hole
The simplest kind of black hole, an ideal black hole, is one that is stationary, electrovac outside its singularity,
and extends to asymptotically flat empty space at infinity. Electrovac means that the energy-momentum
tensor Tµν is zero except for the contribution from a stationary electromagnetic field.
6.3 No-hair theorem 79
The next several chapters deal with ideal black holes. The importance of ideal black holes stems from
the no-hair theorem, discussed in the next section. The no-hair theorem has the consequence that, except
during their initial collapse, or during a merger, real astronomical black holes are accurately described as
ideal outside their horizons.
6.3 No-hair theorem
I will state and justify the no-hair theorem, but I will not prove it mathematically, since the proof is technical.
The no-hair theorem states that a stationary black hole in asymptotically flat space is characterized by
just three quantities:
1. Mass M ;
2. Electric charge Q;
3. Spin, usually parameterized by the angular momentum a per unit mass.
The mechanism by which a black hole loses its hair is gravitational radiation. When initially formed,
whether from the collapse of a massive star or from the merger of two black holes, a black hole will form
a complicated, oscillating region of spacetime. But over the course of several light crossing times, the
oscillations lose energy by gravitational radiation, and damp out, leaving a stationary black hole.
Real astronomical black holes are not isolated, and continue to accrete (cosmic microwave background
photons, if nothing else). However, the timescale (a light crossing time) for oscillations to damp out by
gravitational radiation is usually far shorter than the timescale for accretion, so in practice real black holes
are extremely well described by no-hair solutions almost all of their lives.
The physical reason that the no-hair theorem applies is that space is falling faster than light inside the
horizon. Consequently, unlike a star, no energy can bubble up from below to replace the energy lost by
gravitational radiation, so that the black hole tends to the lowest energy state characterized by conserved
quantities.
As a corollary, the no-hair theorem does not apply from the inner horizon of a black hole inward, because
there space ceases to fall superluminally.
If there exist other absolutely conserved quantities, such as magnetic charge (magnetic monopoles), or
various supersymmetric charges in theories where supersymmetry is not broken, then the black hole will also
be characterized by those quantities.
Black holes are expected not to conserve quantities such as baryon or lepton number that are thought not
to be absolutely conserved, even though they appear to be conserved in low energy physics.
Other stationary solutions exist that describe black holes in spacetimes that are not asymptotically flat,
such as spacetimes with a cosmological constant, or with a uniform electromagnetic field.
It is legitimate to think of the process of reaching a stationary state as analogous to reaching a condition
of thermodynamic equilibrium, in which a macroscopic system is described by a small number of parameters
associated with the conserved quantities of the system.
7
Schwarzschild Black Hole
The Schwarzschild geometry was discovered by Karl Schwarzschild in late 1915 at essentially the same time
that Einstein was arriving at his final version of the General Theory of Relativity.
7.1 Schwarzschild metric
The Schwarzschild metric is, in a polar coordinate system t, r, θ, φ, and in geometric units c = G = 1,
ds2 = −(
1− 2M
r
)
dt2 +
(
1− 2M
r
)−1
dr2 + r2do2 , (7.1)
where do2 (this is the Landau & Lifshitz notation) is the metric of a unit 2-sphere
do2 = dθ2 + sin2θ dφ2 . (7.2)
The Schwarzschild geometry describes the simplest kind of black hole: a black hole with mass M , but no
electric charge, and no spin.
The geometry describes not only a black hole, but also any empty space surrounding a spherically sym-
metric mass. Thus the Schwarzschild geometry describes to a good approximation the spacetime outside the
surfaces of the Sun and the Earth.
Comparison with the spherically symmetric Newtonian metric
ds2 = − (1 + 2Φ)dt2 + (1− 2Φ)(dr2 + r2do2) (7.3)
with Newtonian potential
Φ(r) = −Mr
(7.4)
establishes that the M in the Schwarzschild metric is to be interpreted as the mass of the black hole.
The Schwarzschild geometry is asymptotically flat, because the metric tends to the Minkowski metric in
7.2 Birkhoff’s theorem 81
polar coordinates at large radius
ds2 → − dt2 + dr2 + r2do2 as r →∞ . (7.5)
Exercise 7.1 The Schwarschild metric (7.1) does not have the same form as the spherically symmetric
Newtonian metric (7.3). By a suitable transformation of the radial coordinate r, bring the Schwarschild
metric (7.1) to the isotropic form
ds2 = −(
1−M/2R
1 +M/2R
)2
dt2 + (1 +M/2R)4 (dR2 +R2do2) . (7.6)
What is the relation between R and r? Hence conclude that the identification (7.4) is correct, and therefore
that M is indeed the mass of the black hole. Is the isotropic form (7.6) of the Schwarzschild metric valid
inside the horizon?
7.2 Birkhoff’s theorem
Birkhoff’s theorem states that the geometry of empty space surrounding a spherically symmetric matter
distribution is the Schwarzschild geometry. That is, if the metric is of the form
ds2 = A(t, r) dt2 +B(t, r) dt dr + C(t, r) dr2 +D(t, r) do2 , (7.7)
where the metric coefficients A, B, C, and D are allowed to be arbitrary functions of t and r, and if the
energy momentum tensor vanishes, Tµν = 0, outside some value of the circumferential radius r′ defined by
r′2 = D, then the geometry is necessarily Schwarzschild outside that radius.
This means that if a mass undergoes spherically symmetric pulsations, then those pulsations do not affect
the geometry of the surrounding spacetime. This reflects the fact that there are no spherically symmetric
gravitational waves.
7.3 Stationary, static
The Schwarzschild geometry is stationary. A spacetime is said to be stationary if and only if there exists
a timelike coordinate t such that the metric is independent of t. In other words, the spacetime possesses
time translation symmetry: the metric is unchanged by a time translation t → t + t0 where t0 is some
constant. Evidently the Schwarzschild metric (7.1) is independent of the timelike coordinate t, and is
therefore stationary, time translation symmetric.
The Schwarzschild geometry is also static. A spacetime is static if and only if the coordinates can be
chosen so that, in addition to being stationary with respect to a time coordinate t, the spatial coordinates
82 Schwarzschild Black Hole
do not change along the direction of the tangent vector gt. This requires that the tangent vector gt be
orthogonal to all the spatial tangent vectors
gt · gµ = gtµ = 0 for µ 6= t . (7.8)
The Gullstrand-Painleve metric for the Schwarzschild geometry, discussed in section 7.13, is an example of a
metric that is stationary but not static (although the underlying spacetime, being Schwarzschild, is static).
The Gullstrand-Painleve metric is independent of the free-fall time tff , so is stationary, but observers who
follow the tangent vector gtff fall into the black hole, so the metric is not manifestly static.
The Schwarzschild time coordinate t is thus identified as a special one: it is the unique time coordinate
with respect to which the Schwarzschild geometry is manifestly static.
7.4 Spherically symmetric
The Schwarzschild geometry is also spherically symmetric. This is evident from the fact that the angular
part r2do2 of the metric is the metric of a 2-sphere of radius r. This can be see as follows. Consider the
metric of ordinary flat 3-dimensional Euclidean space in Cartesian coordinates x, y, z:
ds2 = dx2 + dy2 + dz2 . (7.9)
Convert to polar coordinates r, θ, φ, defined so that
x = r sin θ cosφ ,
y = r sin θ sinφ ,
z = r cos θ .
(7.10)
Substituting equations (7.10) into the Euclidean metric (7.9) gives
ds2 = dr2 + r2(dθ2 + sin2θ dφ2) . (7.11)
Restricting to a surface r = constant of constant radius then gives the metric of a 2-sphere of radius r
ds2 = r2(dθ2 + sin2θ dφ2) (7.12)
as claimed.
The radius r in Schwarzschild coordinates is the circumferential radius, defined such that the proper
circumference of the 2-sphere measured by observers at rest in Schwarschild coordinates is 2πr. This is a
coordinate-invariant definition of the meaning of r, which implies that r is a scalar.
7.5 Horizon 83
7.5 Horizon
The horizon of the Schwarzschild geometry lies at the Schwarzschild radius r = rs
rs =2GM
c2. (7.13)
where units of c and G have been restored. Where does this come from? The Schwarzschild metric shows
that the scalar spacetime distance squared ds2 along an interval at rest in Schwarzschild coordinates, dr =
dθ = dφ = 0, is timelike, lightlike, or spacelike depending on whether the radius is greater than, equal to, or
less than r = 2M :
ds2 = −(
1− 2M
r
)
dt2
< 0 if r > 2M ,
= 0 if r = 2M ,
> 0 if r < 2M .
(7.14)
Since the worldline of a massive observer must be timelike, it follows that a massive observer can remain at
rest only outside the horizon, r > 2M . An object at rest at the horizon, r = 2M , follows a null geodesic,
which is to say it is a possible worldline of a massless particle, a photon. Inside the horizon, r < 2M , neither
massive nor massless objects can remain at rest.
A full treatment of what is going on requires solving the geodesic equation in the Schwarzschild geometry,
but the results may be anticipated already at this point. In effect, space is falling into the black hole. Outside
the horizon, space is falling less than the speed of light; at the horizon space is falling at the speed of light;
and inside the horizon, space is falling faster than light, carrying everything with it. This is why light cannot
escape from a black hole: inside the horizon, space falls inward faster than light, carrying light inward even if
that light is pointed radially outward. The statement that space is falling superluminally inside the horizon
of a black hole is a coordinate-invariant statement: massive or massless particles are carried inward whatever
their state of motion and whatever the coordinate system.
Whereas an interval of coordinate time t switches from timelike outside the horizon to spacelike inside the
horizon, an interval of coordinate radius r does the opposite: it switches from spacelike to timelike:
ds2 =
(
1− 2M
r
)−1
dr2
> 0 if r > 2M ,
= 0 if r = 2M ,
< 0 if r < 2M .
(7.15)
It appears then that the Schwarzschild time and radial coordinates swap roles inside the horizon. Inside the
horizon, the radial coordinate becomes timelike, meaning that it becomes a possible worldline of a massive
observer. That is, a trajectory at fixed t and decreasing r is a possible wordline. Again this reflects the fact
that space is falling faster than light inside the horizon. A person inside the horizon is inevitably compelled
as time goes by to move to smaller radial coordinate r.
84 Schwarzschild Black Hole
7.6 Proper time
The proper time experienced by an observer at rest in Schwarzschild coordinates, dr = dθ = dφ = 0, is
dτ =√
−ds2 =
(
1− 2M
r
)1/2
dt . (7.16)
For an observer at rest at infinity, r →∞, the proper time is the same as the coordinate time,
dτ → dt as r →∞ . (7.17)
Among other things, this implies that the Schwarzschild time coordinate t is a scalar: not only is it the
unique coordinate with respect to which the metric is manifestly static, but it coincides with the proper time
of observers at rest at infinity. This coordinate-invariant definition of time t implies that it is a scalar.
At finite radii outside the horizon, r > 2M , the proper time dτ is less than the Schwarzchild time dt, so
the clocks of observers at rest run slower at smaller than at larger radii.
At the horizon, r = 2M , the proper time dτ of an observer at rest goes to zero,
dτ → 0 as r→ 2M . (7.18)
This reflects the fact that an object at rest at the horizon is following a null geodesic, and as such experiences
zero proper time.
7.7 Redshift
An observer at rest at infinity looking through a telesope at an emitter at rest at radius r sees the emitter
redshifted by a factor
1 + z ≡ λobs
λemit=νemit
νobs=
dτobs
dτemit=
(
1− 2M
r
)−1/2
. (7.19)
This is an example of the universally valid statement that photons are good clocks: the redshift factor is
given by the rate at which the emitter’s clock appears to tick relative to the observer’s own clock.
It should be emphasized that the redshift factor (7.19) is valid only for an observer and an emitter at rest
in the Schwarzschild geometry. If the observer and emitter are not at rest, then additional special relativistic
factors will fold into the redshift.
The redshift goes to infinity for an emitter at the horizon
1 + z →∞ as r → 2M . (7.20)
Here the redshift tends to infinity regardless of the motion of the observer or emitter. An observer watching
an emitter fall through the horizon will see the emitter appear to freeze at the horizon, becoming ever slower
and more redshifted. Physically, photons emitted vertically upward at the horizon by an emitter falling
through it remain at the horizon for ever, taking an infinite time to get out to the outside observer.
7.8 Proper distance 85
7.8 Proper distance
The proper radial distance measure by observers at rest in Schwarzschild coordinates, dr = dθ = dφ = 0, is
dl =√ds2 =
(
1− 2M
r
)−1/2
dr . (7.21)
For an observer at rest at infinity, r → ∞, an interval of proper radial distance equals an interval of
circumferential radial distance, as you might expect for asymptotically flat space
dl→ dr as r →∞ . (7.22)
At the horizon, r = 2M , a proper radial interval dl measured by an observer at rest goes to infinity
dl→∞ as r→ 2M . (7.23)
7.9 “Schwarzschild singularity”
The apparent singularity in the Schwarzschild metric at the horizon r = 2M is not a real singularity, because
it can be removed by a change of coordinates, such as to Gullstrand-Painleve coordinates (7.26). Prior to
as late as the 1950s, people, including Einstein, thought that the “Schwarzschild singularity” at r = 2M
marked the physical boundary of the Schwarzschild spacetime. After all, an outside observer watching stuff
fall in never sees anything beyond that boundary.
Schwarzschild’s choice of coordinates was certainly a natural one. It was natural to search for static
solutions, and his time coordinate t is the only one with respect to which the metric is manifestly static.
The problem is that physically there can be no static observers inside the horizon: they must necessarily fall
inward as time passes. The fact that Schwarzschild’s coordinate system shows an apparent singularity at the
horizon reflects the fact that the assumption of a static spacetime necessarily breaks down at the horizon,
where space is falling at the speed of light.
Does stuff “actually” fall in, even though no outside observer ever sees it happen? Classically, the answer
is yes: when a black hole forms, it does actually collapse, and when an observer falls through the horizon,
they really do fall through the horizon. The reason that an outside observer sees everything freeze at the
horizon is simply a light travel time effect: it takes an infinite time for light to lift off the horizon and make
it to the outside world.
7.10 Embedding diagram
An embedding diagram is a visual aid to understanding geometry. It is a depiction of a lower dimensional
geometry in a higher dimension. A classic example is the illustration of the geometry of a 2-sphere embedded
in 3-dimensional space. The 2-sphere has a meaning independent of any embedding in 3 dimensions because
86 Schwarzschild Black Hole
the geometry of the 2-sphere can be measured by 2-dimensional inhabitants of its surface without reference
to any encompassing 3-dimensional space. Nevertheless, the pictorial representation aids imagination.
Textbooks sometimes illustrate the Schwarzschild geometry with an embedding diagram that shows the
spatial geometry at a fixed instant of Schwarzschild time t. The diagram illustrates the stretching of proper
distances in the radial direction. I’ll let you figure out how to construct this embedding diagram.
It should be emphasized that the embedding diagram of the Schwarzshild geometry at fixed Schwarzschild
time t has a limited physical meaning. Fixing the time t means choosing a certain hypersurface through the
geometry. Other choices of hypersurface will yield different diagrams. For example, the Gullstrand-Painleve
metric is spatially flat at fixed free-fall time tff , so in that case the embedding diagram would simply illustrate
flat space, with no funny business at the horizon.
7.11 Energy-momentum tensor
The energy-momentum tensor of the Schwarzschild geometry is zero, by construction.
7.12 Weyl tensor
It turns out that the 10 components of the Weyl tensor, the tidal part of the Riemann tensor, can be decom-
posed in any locally inertial frame into 5 complex components of spin 0, ±1, and ±2. In the Schwarzschild
metric, all components vanish except the real spin 0 component. This component is a coordinate-invariant
scalar, the Weyl scalar C
C = −Mr3
. (7.24)
The Weyl scalar, which expresses the presence of tidal forces, goes to infinity at zero radius,
C →∞ as r → 0 , (7.25)
signalling the presence of a real singularity at zero radius.
7.13 Gullstrand-Painleve coordinates
The Gullstrand-Painleve metric is an alternative metric for the Schwarzschild geometry, discovered indepen-
dently by Allvar Gullstrand and Paul Painleve in (1921). When we have done tetrads, we will recognize
that the standard way in which metrics are written encodes not only metric but also a complete tetrad. The
Gullstrand-Painleve line-element (7.26) encodes a tetrad that represents locally inertial frames free-falling
radially into the black hole at the Newtonian escape velocity. Unlike Schwarzschild coordinates, there is no
singularity at the horizon in Gullstrand-Painleve coordinates. It is striking that the mathematics was known
long before physical understanding emerged.
7.14 Eddington-Finkelstein coordinates 87
The Gullstrand-Painleve metric is
ds2 = − dt2ff + (dr − β dtff)2 + r2do2 . (7.26)
Here β is the Newtonian escape velocity (with a minus sign because space is falling inward)
β = −(
2M
r
)1/2
(7.27)
and tff is the proper time experienced by an object that free falls radially inward from zero velocity at infinity.
The free fall time tff is related to the Schwarzschild time coordinate t by
dtff = dt− β
1− β2dr , (7.28)
which integrates to
tff = t+ 2M
[
2( r
2M
)1/2
+ ln
∣
∣
∣
∣
(r/2M)1/2 − 1
(r/2M)1/2 + 1
∣
∣
∣
∣
]
. (7.29)
The time axis gtff in Gullstrand-Painleve coordinates is not orthogonal to the radial axis gr, but rather is
tilted along the radial axis, gtff · gr = gtffr = −β.
The proper time of a person at rest in Gullstrand-Painleve coordinates, dr = dθ = dφ = 0, is
dτ = dtff√
1− β2 . (7.30)
The horizon occurs where this proper time vanishes, which happens when the infall velocity β is the speed
of light
|β| = 1 . (7.31)
According to equation (7.27), this happens at r = 2M , which is the Schwarzschild radius, as it should be.
7.14 Eddington-Finkelstein coordinates
In Schwarzschild coordinates, radially infalling or outfalling light rays appear never to cross the horizon
of the Schwarzschild black hole. This feature of Schwarzschild coordinates contributed to the historical
misconception that black holes stopped at their horizons. In 1958, David Finkelstein carried out a trivial
transformation of the time coordinate which seeemed to show that infalling light rays could indeed pass
through the horizon. It turned out that Eddington had already discovered the transformation in 1924,
though at that time the physical implications were not grasped. Again, it is striking that the mathematics
was in place long before physical understanding.
In Schwarzschild coordinates, light rays that fall radially (dθ = dφ = 0) inward or outward follow null
geodesics
ds2 = −(
1− 2M
r
)
dt2 +
(
1− 2M
r
)−1
dr2 = 0 . (7.32)
88 Schwarzschild Black Hole
Radial null geodesics thus follow
dr
dt= ±
(
1− 2M
r
)
(7.33)
in which the ± sign is + for outfalling, − for infalling rays. Equation (7.33) shows that dr/dt → 0 as
r → 2M , suggesting that null rays, whether infalling or outfalling, never cross the horizon. The solution to
equation (7.33) is
t = ± (r + 2M ln|r − 2M |) , (7.34)
which shows that Schwarzschild time t approaches ±∞ logarithmically as null rays approach the horizon.
Finkelstein defined his time coordinate tF by
tF ≡ t+ 2M ln |r − 2M | , (7.35)
which has the property that infalling null rays follow
tF + r = 0 . (7.36)
In other words, on a spacetime diagram in Finkelstein coordinates, radially infalling light rays move at 45,the same as in special relativistic spacetime diagrams.
7.15 Kruskal-Szekeres coordinates
Since Finkelstein transformed coordinates so that radially infalling light rays moved at 45 in a spacetime
diagram, it is natural to look for coordinates in which outfalling as well as infalling light rays are at 45.Kruskal and Szekeres independently provided such a transformation, in 1960.
Define the tortoise (or Regge-Wheeler 1959) coordinate r∗ by
r∗ ≡∫
dr
1− 2M/r= r + 2M ln |r − 2M | . (7.37)
Then radially infalling and outfalling null rays follow
r∗ + t = 0 infalling ,
r∗ − t = 0 outfalling .(7.38)
In a spacetime diagram in coordinates t and r∗, infalling and outfalling light rays are indeed at 45. Unfor-
tunately the metric in these coordinates is still singular at the horizon r = 2M :
ds2 =
(
1− 2M
r
)
(
− dt2 + dr∗2)
+ r2do2 . (7.39)
The singularity at the horizon can be eliminated by the following transformation into Kruskal-Szekeres
coordinates tK and rK :
rK + tK = e(r∗+t)/2 ,
rK − tK = ±e(r∗−t)/2 ,(7.40)
7.16 Penrose diagrams 89
r = ∞
r =∞Hor
izon
Antihorizon Universe
Singularity (r = 0)
Black Hole
Tim
e
Space
Light
Light
Figure 7.1 Penrose diagram of the Schwarzschild geometry.
where the ± sign in the last equation is + outside the horizon, − inside the horizon. The Kruskal-Szekeres
metric is
ds2 = r−1e−r(
− dt2K + dr2K)
+ r2do2 , (7.41)
which is non-singular at the horizon. The Schwarzschild radial coordinate r, which appears in the factors
r−1e−r and r2 in the Kruskal metric, is to be understood as an implicit function of the Kruskal coordinates
tK and rK .
7.16 Penrose diagrams
Roger Penrose, as so often, had a novel take on the business of spacetime diagrams. Penrose conceived that
the primary purpose of a spacetime diagram should be to portray the causal structure of the spacetime, and
that the specific choice of coordinates was largely irrelevant. After all, general relativity allows arbitrary
choices of coordinates.
In addition to requiring that light rays be at 45, Penrose wanted to bring regions at infinity (in time or
space) to a finite position on the spacetime diagram, so that the entire spacetime could be seen at once. He
calls these thing conformal diagrams, but the rest of us commonly call them Penrose diagrams.
Penrose diagrams are bona-fide spacetime diagrams. For example, a coordinate transformation from
Kruskal to “Penrose” coordinates (the following transformation is not analytic, but Penrose does not care)
rP + tP =rK + tK
2 + |rK + tK |,
rP − tP =rK − tK
2− |rK − tK |,
(7.42)
90 Schwarzschild Black Hole
brings spatial and temporal infinity to finite values of the coordinates, while keeping infalling and outfalling
light rays at 45 in the spacetime diagram. However, there are many such transformations, and Penrose
would be the last person to advocate any one of them in particular.
r = 0
r = ∞
r =∞
r =∞
r = ∞
Horiz
on
Universe
Black Hole
r = 0
Parallel Horizon
Antihorizon
Para
llel Ant
ihor
izon
Parallel Universe
White Hole
Figure 7.2 Penrose diagram of the complete, analytically extended Schwarzschild geometry.
7.17 Schwarzschild white hole, wormhole
The Kruskal-Szekeres spacetime diagram reveals a new feature that was not apparent in Schwarzschild or
Finkelstein coordinates. Dredged from the depths of t = −∞ appears a null line rK + tK = 0. The null line
is at radius r = 2M , but it does not correspond to the horizon that a person might fall into. The null line is
called the antihorizon. The horizon is sometimes called the true horizon, and the antihorizon the illusory
horizon. In a real black hole, only the true horizon is real. The antihorizon is replaced by an exponentially
dimming and redshifting image of the star that collapsed to form the black hole.
The Kruskal-Szekeres (= Schwarzschild) geometry is analytic, and there is a unique analytic continuation
of the geometry through the antihorizon. The analytic continuation is a time-reversed copy of the original
Schwarzschild geometry, glued at the antihorizon. Whereas the original Schwarzschild geometry showed an
asymptotically flat region and a black hole region separated by a horizon, the complete analytically extended
Schwarzschild geometry shows two asymptotically flat regions, together with a black hole and a white hole.
Relativists label the regions I, II, III, and IV, but I like to call them by name: “Universe”, “Black Hole”,
“Parallel Universe”, and “White Hole”.
The white hole is a time-reversed version of the black hole. Whereas space falls inward faster than light
inside the black hole, space falls outward faster than light inside the white hole. In the Gullstrand-Painleve
metric (7.26), the velocity β = ±(2M/r)1/2 is negative for the black hole, positive for the white hole.
The Kruskal or Penrose diagrams show that the universe and the parallel universe are connected, but
7.18 Collapse to a black hole 91
only by spacelike lines. This spacelike connection is called the Einstein-Rosen bridge, and constitutes a
wormhole connecting the two universes. Because the connection is spacelike, it is impossible for a traveler
to pass through this wormhole.
Although two travelers, one from the universe and one from the parallel universe, cannot travel to each
other’s universe, they can meet, but only inside the black hole. Inside the black hole, they can talk to each
other, and they can see light from each other’s universe. Sadly, the enlightenment is only temporary, because
they are doomed soon to hit the central singularity.
It should be emphasized that the white hole and the wormhole in the Schwarzschild geometry are a
mathematical construction with as far as anyone knows no relevance to reality. Nevertheless it is intriguing
that such bizarre objects emerge already in the simplest general relativistic solution for a black hole.
7.18 Collapse to a black hole
Realistic collapse of a star to a black hole is not expected to produce a white hole or parallel universe.
The simplest model of a collapsing star is a spherical ball of uniform density and zero pressure which free
falls from zero velocity at infinity. In this simple model, the interior of the star is described by a collapsing
Friedmann-Robertson-Walker metric (the canonical cosmological metric), while the exterior is described by
the Schwarzschild solution. The assumption that the star collapses from zero velocity at infinity implies
that the FRW metric is spatially flat, the simplest case. To continue the geometry between Schwarzchild
and FRW metrics, it is neatest to use the Gullstrand-Painleve metric, with the Gullstrand-Painleve infall
velocity β at the edge of the star set equal to minus r times the Hubble parameter −rH ≡ −r d ln a/dt of
the collapsing FRW metric.
The simple model shows that the antihorizon of the complete Schwarzschild geometry is replaced by the
surface of the collapsing star, and that beyond the antihorizon is not a parallel universe and a white hole,
but merely the interior of the star (and the distant Universe glimpsed through the star’s interior).
Since light can escape from the collapsing star system as long as it is even slightly larger than its Schwarz-
schild radius, it is possible to take the view that the horizon comes instantaneously into being at the moment
the star collapses through its Schwarzschild radius. This definition of the horizon is called the apparent hori-
zon.
Hawking has advocated that a better definition of the horizon is to take it to be the boundary between
outgoing null rays that fall into the black hole versus those that go to infinity. In any evolving situation,
this definition of the horizon, which is called the absolute horizon, depends formally on what happens in the
infinite future, though in slowly evolving systems the absolute horizon can be located with some precision
without knowing the future. The absolute horizon of the collapsing star forms before the star has collapsed,
and grows to meet the apparent horizon as the star falls through its Schwarzchild radius.
In this simple model, the central singularity forms slightly before the star has collapsed to zero radius.
The formation of the singularity is marked by the fact that light rays emitted at zero radius cease to be able
to move outward. In other words, the singularity forms when space starts to fall into it faster than light.
92 Schwarzschild Black Hole
7.19 Killing vectors
The Schwarzschild metric presents an opportunity to introduce the concept of Killing vectors (after Wil-
helm Killing, not because the vectors kill things, though the latter is true), which are associated with
symmetries of the spacetime.
7.20 Time translation symmetry
The time translation invariance of the Schwarzschild geometry is evident from the fact that the metric is
independent of the time coordinate t. Equivalently, the partial time derivative ∂/∂t of the Schwarzschild
metric is zero. The associated Killing vector ξµ is then defined by
ξµ ∂
∂xµ=
∂
∂t(7.43)
so that in Schwarzschild coordinates t, r, θ, φ
ξµ = 1, 0, 0, 0 . (7.44)
In coordinate-independent notation, the Killing vector is
ξ = gµξµ = gt . (7.45)
This may seem like overkill – couldn’t we just say that the metric is independent of time t and be done
with it? The answer is that symmetries are not always evident from the metric, as will be seen in the next
section 7.21.
Because the Killing vector gt is the unique timelike Killing vector of the Schwarzschild geometry, it has
a definite meaning independent of the coordinate system. It follows that its scalar product with itself is a
coordinate-independent scalar
ξµξµ = gt · gt = gtt = −
(
1− 2M
r
)
. (7.46)
In curved spacetimes, it is hugely important to be able to identify scalars, which have a physical meaning
independent of the choice of coordinates.
7.21 Spherical symmetry
The rotational symmetry of the Schwarzschild metric about the azimuthal axis is evident from the fact that
the metric is independent of the azimuthal coordinate φ. The associated Killing vector is
gφ (7.47)
with components 0, 0, 0, 1 in Schwarzschild coordinates t, r, θ, φ.
7.22 Killing equation 93
The Schwarzschild metric is fully spherically symmetric, not just azimuthally symmetric. Since the 3D
rotation group O(3) is 3-dimensional, it is to be expected that there are three Killing vectors. You may
recognize from quantum mechanics that ∂/∂φ is (modulo factors of i and ~) the z-component of the angular
momentum operator L = Lx, Ly, Lz in a coordinate system where the azimuthal axis is the z-axis. The 3
components of the angular momentum operator are given by:
iLx = y∂
∂z− z ∂
∂y= − sinφ
∂
∂θ− cot θ cosφ
∂
∂φ,
iLy = z∂
∂x− x ∂
∂z= cosφ
∂
∂θ− cot θ sinφ
∂
∂φ,
iLz = x∂
∂y− y ∂
∂x=
∂
∂φ.
(7.48)
The 3 rotational Killing vectors are correspondingly:
rotation about x-axis: − sinφgθ − cot θ cosφgφ ,
rotation about y-axis: cosφgθ − cot θ sinφgφ ,
rotation about z-axis: gφ .
(7.49)
You can check that the action of the x and y rotational Killing vectors on the metric does not kill the
metric. For example, iLxgφφ = 2r2 cosφ sin θ cos θ does not vanish. This example shows that a more powerful
and general condition, described in the next section 7.22, is needed to establish whether a quantity is or is
not a Killing vector.
Because spherical symmetry does not define a unique azimuthal axis gφ, its scalar product with itself
gφ · gφ = gφφ = −r2 sin2θ is not a coordinate-invariant scalar. However, the sum of the scalar products of
the 3 rotational Killing vectors is rotationally invariant, and is therefore a coordinate-invariant scalar
(− sinφgθ − cot θ cosφgφ)2 + (cosφgθ − cot θ sinφgφ)2 + g2φ = gθθ + (cot2θ + 1)gφφ = −2r2 . (7.50)
This shows that the circumferential radius r is a scalar, as you would expect.
7.22 Killing equation
As seen in the previous section, a Killing vector does not always kill the metric in a given coordinate system.
This is not really surprising given the arbitrariness of coordinates in GR. What is true is that a quantity is
a Killing vector if and only if there exists a coordinate system such that the Killing vector kills the metric
in that system.
Suppose that in some coordinate system the metric is independent of the coordinate φ. In problem set 2
you showed that in such a case the covariant component uφ of the 4-velocity along a geodesic is constant
uφ = constant . (7.51)
Equivalently
ξνuν = constant (7.52)
94 Schwarzschild Black Hole
where ξν is the associated Killing vector, whose only non-zero component is ξφ = 1 in this particular
coordinate system. The converse is also true: if ξνuν = constant along all geodesics, then ξν is a Killing
vector. The constancy of ξνuν along all geodesics is equivalent to the condition that its proper time derivative
vanish along all geodesics
dξνuν
dτ= 0 . (7.53)
But this is equivalent to
0 = uµDµ(ξνuν) = uµuνDµξν =1
2uµuν(Dµξν +Dνξµ) (7.54)
where the second equality follows from the geodesic equation, uµDµuν = 0, and the last equality is true
because of the symmetry of uµuν in µ ↔ ν. A necessary and sufficient condition for equation (7.54) to be
true for all geodesics is that
Dµξν +Dνξµ = 0 (7.55)
which is Killing’s equation. This equation is the desired necessary and sufficient condition for ξν to be a
Killing vector. It is a generally covariant equation, valid in any coordinate system.
8
Reissner-Nordstrom Black Hole
The Reissner-Nordstrom geometry, discovered independently by Hans Reissner in 1916, Hermann Weyl in
1917, and Gunnar Nordstrom in 1918, describes the unique spherically symmetric static solution for a black
hole with mass and electric charge in asymptotically flat spacetime.
8.1 Reissner-Nordstrom metric
The Reissner-Nordstrom metric for a black hole of mass M and electric charge Q is, in geometric units
c = G = 1,
ds2 = −(
1− 2M
r+Q2
r2
)
dt2 +
(
1− 2M
r+Q2
r2
)−1
dr2 + r2do2 (8.1)
which looks like the Schwarzschild metric with the replacement
M →M(r) = M − Q2
2r. (8.2)
In fact equation (8.2) has a coordinate independent interpretation as the mass M(r) interior to radius r,
which here is the mass M at infinity, minus the mass in the electric field E = Q/r2 outside r∫ ∞
r
E2
8π4πr2dr =
∫ ∞
r
Q2
8πr44πr2dr =
Q2
2r. (8.3)
This seems like a Newtonian calculation of the energy in the electric field, but it turns out to be valid also
in general relativity.
Real astronomical black holes probably have very little electric charge, because the Universe as a whole
appears almost electrically neutral (and Maxwell’s equations in fact require that the Universe in its entirety
should be exactly electrically neutral), and a charged black hole would quickly neutralize itself. It would
probably not neutralize itself completely, but have some small residual positive charge, because protons
(positive charge) are more massive than electrons (negative charge), so it is slightly easier for a black hole
to accrete protons than electrons.
96 Reissner-Nordstrom Black Hole
Nevertheless, the Reissner-Nordstrom solution is of more than passing interest because its internal geom-
etry resembles that of the Kerr solution for a rotating black hole.
Concept question 8.1 What is the charge Q in standard (gaussian) units?
8.2 Energy-momentum tensor
The Einstein tensor of the Reissner-Nordstrom metric (8.1) is diagonal, with elements given by
Gνµ =
Gtt 0 0 0
0 Grr 0 0
0 0 Gθθ 0
0 0 0 Gφφ
= 8π
−ρ 0 0 0
0 pr 0 0
0 0 p⊥ 0
0 0 0 p⊥
=Q2
r4
−1 0 0 0
0 −1 0 0
0 0 1 0
0 0 0 1
. (8.4)
The trick of writing one index up and the other down on the Einstein tensor Gνµ partially cancels the
distorting effect of the metric, yielding the proper energy density ρ, the proper radial pressure pr, and
transverse pressure p⊥, up to factors of ±1. A more systematic way to extract proper quantities is to work
in the tetrad formalism, but this will do for now.
The energy-momentum tensor is that of a radial electric field
E =Q
r2. (8.5)
Notice that the radial pressure pr is negative, while the transverse pressure p⊥ is positive. It is no coincidence
that the sum of the energy density and pressures is twice the energy density, ρ+ pr + 2p⊥ = 2ρ.
The negative pressure, or tension, of the radial electric field produces a gravitational repulsion that domi-
nates at small radii, and that is responsible for much of the strange phenomenology of the Reissner-Nordstrom
geometry. The gravitational repulsion mimics the centrifugal repulsion inside a rotating black hole, for which
reason the Reissner-Nordstrom geometry is often used a surrogate for the rotating Kerr-Newman geometry.
At this point, the statements that the energy-momentum tensor is that of a radial electric field, and that
the radial tension produces a gravitational repulsion that dominates at small radii, are true but unjustified
assertions.
8.3 Weyl tensor
As with the Schwarzschild geometry (indeed, any spherically symmetric geometry), only 1 of the 10 inde-
pendent spin components of the Weyl tensor is non-vanishing, the real spin-0 component, the Weyl scalar
8.4 Horizons 97
C. The Weyl scalar for the Reissner-Nordstrom geometry is
C = − M
r3+Q2
r4. (8.6)
The Weyl scalar goes to infinity at zero radius
C →∞ as r → 0 (8.7)
signalling the presence of a real singularity at zero radius.
8.4 Horizons
The Reissner-Nordstrom geometry has not one but two horizons. The horizons occur where an object at
rest in the geometry, dr = dθ = dφ = 0, follows a null geodesic, ds2 = 0, which occurs where
1− 2M
r+Q2
r2= 0 . (8.8)
This is a quadratic equation in r, and it has two solutions, an outer horizon r+ and an inner horizon r−
r± = M ±√
M2 −Q2 . (8.9)
It is straightforward to check that the Reissner-Nordstrom time coordinate t is timelike outside the outer
horizon, r > r+, spacelike between the horizons r− < r < r+, and again timelike inside the inner horizon
r < r−. Conversely, the radial coordinate r is spacelike outside the outer horizon, r > r+, timelike between
the horizons r− < r < r+, and spacelike inside the inner horizon r < r−.
The physical meaning of this strange behaviour is akin to that of the Schwarzschild geometry. As in the
Schwarzschild geometry, outside the outer horizon space is falling at less than the speed of light; at the outer
horizon space hits the speed of light; and inside the outer horizon space is falling faster than light. But
a new ingredient appears. The gravitational repulsion caused by the negative pressure of the electric field
slows down the flow of space, so that it slows back down to the speed of light at the inner horizon. Inside
the inner horizon space is falling at less than the speed of light.
8.5 Gullstrand-Painleve metric
Deeper insight into the Reissner-Nordstrom geometry comes from examining its Gullstrand-Painleve metric.
The Gullstrand-Painleve metric for the Reissner-Nordstrom geometry is the same as that for the Schwarz-
schild geometry
ds2 = − dt2ff + (dr − β dtff)2 + r2do2 . (8.10)
98 Reissner-Nordstrom Black Hole
The velocity β is again the escape velocity, but this is now
β = ∓√
2M(r)
r(8.11)
where M(r) = M − Q2/2r is the interior mass already given as equation (8.2). Horizons occur where the
magnitude of the velocity β equals the speed of light
|β| = 1 (8.12)
which happens at the outer and inner horizons r = r+ and r = r−, equation (8.9).
The Gullstrand-Painleve metric once again paints the picture of space falling into the black hole. Outside
the outer horizon r+ space falls at less than the speed of light, at the horizon space falls at the speed of
light, and inside the horizon space falls faster than light. But the gravitational repulsion produced by the
tension of the radial electric field starts to slow down the inflow of space, so that the infall velocity reaches
a maximum at r = Q2/M . The infall slows back down to the speed of light at the inner horizon r−. Inside
the inner horizon, the flow of space slows all the way to zero velocity, β = 0, at the turnaround radius
r0 =Q2
2M. (8.13)
Space then turns around, the velocity β becoming positive, and accelerates back up to the speed of light.
Space is now accelerating outward, to larger radii r. The outfall velocity reaches the speed of light at the
inner horizon r−, but now the motion is outward, not inward. Passing back out through the inner horizon,
space is falling outward faster than light. This is not the black hole, but an altogether new piece of spacetime,
a white hole. The white hole looks like a time-reversed black hole. As space falls outward, the gravitational
repulsion produced by the tension of the radial electric field declines, and the outflow slows. The outflow
slows back to the speed of light at the outer horizon r+ of the white hole. Outside the outer horizon of the
white hole is a new universe, where once again space is flowing at less than the speed of light.
What happens inward of the turnaround radius r0, equation (8.13)? Inside this radius the interior mass
M(r), equation (8.2), is negative, and the velocity β is imaginary. The interior massM(r) diverges to negative
infinity towards the central singularity at r → 0. The singularity is timelike, and infinitely gravitationally
repulsive, unlike the central singularity of the Schwarzschild geometry. Is it physically realistic to have a
singularity that has infinite negative mass and is infinitely gravitationally repulsive? Undoubtedly not.
8.6 Complete Reissner-Nordstrom geometry
As with the Schwarzschild geometry, it is possible to go through the steps: Reissner-Nordstrom coordinates→Eddington-Finkelstein coordinates → Kruskal-Szekeres coordinates → Penrose coordinates. The conclusion
of these constructions is that the Reissner-Nordstrom geometry can be analytically continued, and the
complete analytic continuation consists of an infinite ladder of universes and parallel universes connected to
each other by black hole → wormhole → white hole tunnels. I like to call the various pieces of spacetime
8.6 Complete Reissner-Nordstrom geometry 99
r = ∞
r =∞
r =∞
r = ∞
r =−∞
r = −∞
r = −∞
r =−∞
Para
llel Ant
ihor
izon Antihorizon
Parallel Universe Universe
Parallel HorizonHor
izon
Black Hole
Inne
r Horiz
on
Parallel Inner Horizon
Wor
mho
le ParallelW
ormhole
AntiverseParallel
Antiverse
White HoleInner Antihorizon
Para
llel In
ner Ant
ihor
izon
New Parallel Universe New Universe
Figure 8.1 Penrose diagram of the complete Reissner-Nordstrom geometry.
100 Reissner-Nordstrom Black Hole
“Universe”, “Parallel Universe”, “Black Hole”, “Wormhole”, “Parallel Wormhole”, and “White Hole”. These
pieces repeat in an infinite ladder.
The Wormhole and Parallel Wormhole contain separate central singularities, the “Singularity” and the
“Parallel Singularity”, which are oppositely charged. If the black hole is positively charged as measured by
observers in the Universe, then it is negatively charged as measured by observers in the Parallel Universe,
and the Wormhole contains a positive charge singularity while the Parallel Wormhole contains a negative
charge singularity.
Where does the electric charge of the Reissner-Nordstrom geometry “actually” reside? This comes down
to the question of how observers detect the presence of charge. Observers detect charge by the electric field
that it produces. Equip all (radially moving) observers with a gyroscope that they orient consistently in
the same radial direction, which can be taken to be towards the black hole as measured by observers in
the Universe. Observers in the Parallel Universe find that their gyroscope is pointed away from the black
hole. Inside the black hole, observers from either Universe agree that the gyroscope is pointed towards the
Wormhole, and away from the Parallel Wormhole. All observes agree that the electric field is pointed in the
same radial direction. Observers who end up inside the Wormhole measure an electric field that appears to
emanate from the Singularity, and which they therefore attribute to charge in the Singularity. Observers
who end up inside the Parallel Wormhole measure an electric field that appears to emanate in the opposite
direction from the Parallel Singularity, and which they therefore attribute to charge of opposite sign in the
Parallel Singularity. Strange, but all consistent.
8.7 Antiverse: Reissner-Nordstrom geometry with negative mass
It is also possible to consider the Reissner-Nordstrom geometry for negative values of the radius r. I call the
extension to negative r the “Antiverse”. There is also a “Parallel Antiverse”.
Changing the sign of r in the Reissner-Nordstrom metric (8.1) is equivalent to changing the sign of the
mass M . Thus the Reissner-Nordstrom metric with negative r describes a charged black hole of negative
mass
M < 0 . (8.14)
The negative mass black hole is gravitationally repulsive at all radii, and it has no horizons.
8.8 Ingoing, outgoing
The black hole in the Reissner-Nordstrom geometry is bounded at its inner edge by not one but two inner
horizons. The two distinct horizons play a crucial role in the mass inflation instability described in §8.9
below.
The inner horizons can be called ingoing and outgoing. Persons freely falling in the Black Hole region are
8.9 Mass inflation instability 101
all moving inward in coordinate radius r, but they may be moving either forward or backward in Reissner-
Nordstrom coordinate time t. In the Black Hole region, the conserved energy along a geodesic is positive if
the time coordinate t is decreasing, negative if the time coordinate t is increasing1. Persons with positive
energy are ingoing, while persons with negative energy are outgoing. Both ingoing and outgoing persons
fall inward, to smaller radii, but ingoing persons think that the inward direction is towards the Wormhole,
while outgoing persons think that the inward direction is in the opposite direction, towards the Parallel
Wormhole. Ingoing persons fall through the ingoing inner horizon, while outgoing persons fall through the
outgoing inner horizon.
Coordinate time t moves forwards in the Universe and Wormhole regions, and geodesics have positive
energy in these regions. Conversely, coordinate time t moves backwards in the Parallel Universe and Parallel
Wormhole regions, and geodesics have negative energy in these regions. Of course, all observers, whereever
they may be, always perceive their own proper time to be moving forward in the usual fashion, at the rate
of one second per second.
8.9 Mass inflation instability
Roger Penrose (1968) first pointed out that a person passing through the outgoing inner horizon (also
called the Cauchy horizon) of the Reissner-Nordstrom geometry would see the outside Universe infinitely
blueshifted, and he suggested that this would destabilize the geometry. Perturbation theory calculations,
starting with Simpson & Penrose (1973) and culminating with Chandrasekhar and Hartle (1982), confirmed
that waves become infinitely blueshifted as they approach the outgoing inner horizon, and that their energy
density diverges. The perturbation theory calculations were widely construed as indicating that the Reissner-
Nordstrom geometry was “unstable”, although the precise nature of this instability remained obscure.
It was not until a seminal paper by Poisson & Israel (1990) that the true nonlinear nature of the instability
at the inner horizon was clarified. Poisson & Israel showed that the Reissner-Nordstrom geometry is subject
to an exponentially growing instability which they dubbed mass inflation. The term refers to the fact that
the interior mass M(r) grows exponentially during mass inflation. The interior mass M(r) has the property
of being a gauge-invariant, scalar quantity, so it has a physical meaning independent of the coordinate system.
What causes mass inflation? Actually it has nothing to do with mass: the inflating mass is just a symptom
of the underlying cause. What causes mass inflation is relativistic counter-streaming between ingoing and
outgoing streams. As the Penrose diagram of the Reissner-Nordstrom geometry shows, ingoing and outgoing
streams must drop through separate ingoing and outgoing inner horizons into separate pieces of spacetime,
the Wormhole and the Parallel Wormhole. The regions of spacetime must be separate because coordinate
time t is timelike in both regions, but going in opposite directions in the two regions, forward in the Wormhole,
backward in the Parallel Wormhole. In other words, ingoing and outgoing streams cannot co-exist in the
1 The apparently wrong direction of time comes from requiring that t → ∞ for ingoing trajectories near the horizon, whetherjust above or just below the horizon, while t → −∞ for outgoing trajectories near the horizon, whether just above or justbelow the horizon. It would be unconventional, but consistent, to flip the sign of time t between the horizons so thatpositive energy increased with time, and negative energy decreased with time.
102 Reissner-Nordstrom Black Hole
Universe
Horiz
on
Black HoleIn
goin
gin
ner ho
rizon
Outgoinginner horizon
IngoingOut
goin
g
t
t
t t
Figure 8.2 Penrose diagram illustrating why the Reissner-Nordstrom geometry is subject to the mass infla-tion instability. Ingoing and outgoing streams just outside the inner horizon must pass through separateingoing and outgoing inner horizons into causally separated pieces of spacetime where the timelike timecoordinate t goes in opposite directions. To accomplish this, the ingoing and outgoing streams must exceedthe speed of light through each other, which physically they cannot do. The mass inflation instability isdriven by the pressure of the relativistic counter-streaming between ingoing and outgoing streams. Theinset shows the direction of coordinate time t in the various regions. Proper time of course always increasesupward in a Penrose diagram.
same subluminal region of spacetime because they would have to be moving in opposite directions in time,
which cannot be.
In the Reissner-Nordstrom geometry, ingoing and outgoing streams resolve their differences by exceeding
the speed of light relative to each other, and passing into causally separated regions. As the ingoing and
outgoing streams drop through their respective inner horizons, they each see the other stream infinitely
blueshifted.
In reality however, this cannot occur: ingoing and outgoing streams cannot exceed the speed of light
relative to each other. Instead, as the ingoing and outgoing streams move ever faster through each other
in their effort to drop through the inner horizon, their counter-streaming generates a radial pressure. The
pressure, which is positive, exerts an inward gravitational force. As the counter-streaming approaches the
8.10 Inevitability of mass inflation 103
speed of light, the gravitational force produced by the counter-streaming pressure eventually exceeds the
gravitational force produced by the background Reissner-Nordstrom geometry. At this point, mass inflation
begins.
The gravitational force produced by the counter-streaming is inwards, but, in the strange way that general
relativity operates, the inward direction is in opposite directions for the ingoing streams, towards the black
hole for the ingoing stream, and away from the black hole for the outgoing stream. Consequently the counter-
streaming pressure simply accelerates the ingoing and outgoing streams ever faster through each other. The
result is an exponential feedback instability. The increasing pressure accelerates the streams faster through
each other, which increases the pressure, which increases the acceleration.
The interior mass is not the only thing that increases exponentially during mass inflation. The proper
density and pressure, and the Weyl scalar (all gauge-invariant scalars) exponentiate together.
Exercise 8.2 Show that, in the Reissner-Nordstrom geometry, the blueshift of a photon with energy
vt = ±1 and angular momentum per unit energy v⊥ = J observed by observer on a geodesic with energy
per unit mass ut = E and angular momentum per unit mass u⊥ = L is
uµvµ =
something
B. (8.15)
Argue that the blueshift diverges at the horizon for ingoing observers observing outgoing photons, and for
outgoing observers observing ingoing photons.
8.10 Inevitability of mass inflation
Mass inflation requires the simultaneous presence of both ingoing and outgoing streams near the inner
horizon. Will that happen in real black holes? Any real black hole will of course accrete matter from its
surroundings, so certainly there will be a stream of one kind or another (ingoing or outgoing) inside the
black hole. But is it guaranteed that there will also be a stream of the other kind? The answer is probably.
One of the remarkable features of the mass inflation instability is that, as long as ingoing and outgoing
streams are both present, the smaller the perturbation the more violent the instability. That is, if say the
outgoing stream is reduced to a tiny trickle compared to the ingoing stream (or vice versa), then the length
scale (and time scale) over which mass inflation occurs gets shorter. During mass inflation, as the counter-
streaming streams drop through an interval ∆r of circumferential radius, the interior mass M(r) increases
exponentially with length scale l
M(r) ∝ e∆r/l . (8.16)
It turns out that the inflationary length scale l is proportional to the accretion rate
l ∝ M , (8.17)
104 Reissner-Nordstrom Black Hole
so that smaller accretion rates produce more violent inflation. Physically, the smaller accretion rate, the
closer the streams must approach the inner horizon before the pressure of their counter-streaming begins to
dominate the gravitational force. The distance between the inner horizon and where mass inflation begins
effectively sets the length scale l of inflation.
Given this feature of mass inflation, that the tinier the perturbation the more rapid the growth, it seems
almost inevitable that mass inflation must occur inside real black holes. Even the tiniest piece of stuff going
the wrong way is apparently enough to trigger the mass inflation instability.
One way to avoid mass inflation inside a real black hole is to have a large level of dissipation inside the
black hole, sufficient to reduce the charge (or spin) to zero near the singularity. In that case the central
singularity reverts to being spacelike, like the Schwarzschild singularity. While the electrical conductivity of
a realistic plasma is more than adequate to neutralize a charged black hole, angular momentum transport
is intrinsically a much weaker process, and it is not clear whether the dissipation of angular momentum
might be large enough to eliminate the spin near the singularity of a rotating black hole. There has been no
research on the latter subject.
8.11 The black hole particle accelerator
A good way to think conceptually about mass inflation is that it acts like a particle accelerator. The counter-
streaming pressure accelerates ingoing and outgoing streams through each other at an exponential rate, so
that a Lagrangian gas element spends equal amounts of proper time accelerating through equal decades of
counter-streaming velocity. The center of mass energy easily exceeds the Planck energy, where quantum
mechanics presumably comes into play.
Mass inflation is expected to occur just above the inner horizon of a black hole. In a realistic rotating
astronomical black hole, the inner horizon is likely to be at a considerable fraction of the radius of the
outer horizon. Thus the black hole accelerator operates not near a central singularity, but rather at a
macroscopically huge scale. This machine is truly monstrous.
Undoubtedly much fascinating physics occurs in the black hole particle accelerator. The situation is far
more extreme than anywhere else in our Universe today. Who knows what Nature does there? To my
knowledge, there has been no research on the subject.
8.12 The X point
The point in the Reissner-Nordstrom geometry where the ingoing and outgoing inner horizons intersect, the
X point, is a special one. This is the point through which geodesics of zero energy must pass. Persons with
zero energy who reach the X point see both ingoing and outgoing streams, coming from opposite directions,
infinitely blueshifted.
8.13 Extremal Reissner-Nordstrom geometry 105
r =∞
r =∞
r =∞
r = ∞
r = ∞
r = −∞
r = −∞
r =−∞
r =−∞
r =−∞
Antihorizon
Universe
Horiz
onWor
mho
leAntiverse
New Universe
Figure 8.3 Penrose diagram of the extremal Reissner-Nordstrom geometry.
8.13 Extremal Reissner-Nordstrom geometry
So far the discussion of the Reissner-Nordstrom geometry has centered on the case Q < M (or more generally,
|Q| < |M |) where there are separate outer and inner horizons. In the special case that the charge and mass
are equal,
Q = M , (8.18)
106 Reissner-Nordstrom Black Hole
the inner and outer horizons merge into one, r+ = r−, equation (8.9). This special case describes the
extremal Reissner-Nordstrom geometry.
The extremal Reissner-Nordstrom geometry proves to be of particular interest in quantum gravity because
its Hawking temperature is zero, and in string theory because extremal black holes arise as solutions under
certain duality transformations.
The Penrose diagram of the extremal Reissner-Nordstrom geometry is different from that of the standard
Reissner-Nordstrom geometry.
8.14 Reissner-Nordstrom geometry with charge exceeding mass
The Reissner-Nordstrom geometry with charge greater than mass,
Q > M , (8.19)
has no horizons. The change in geometry from an extremal black hole, with horizon at finite radius r+ =
r− = M , to one without horizons is discontinuous. This suggests that there is no way to pack a black hole
with more charge than its mass. Indeed, if you try to force additional charge into an extremal black hole,
then the work needed to do so increases its mass so that the charge Q does not exceed its mass M .
Real fundamental particles nevertheless have charge far exceeding their mass. For example, the charge-
to-mass ratio of a proton ise
mp≈ 1018 (8.20)
where e is the square root of the fine-structure constant α ≡ e2/~c ≈ 1/137, and mp ≈ 10−19 is the mass of
the proton in Planck units. However, the Schwarzschild radius of such a fundamental particle is far tinier
than its Compton wavelength ∼ ~/m (or its classical radius e2/m = α~/m), so quantum mechanics, not
general relativity, governs the structure of these fundamental particles.
8.15 Reissner-Nordstrom geometry with imaginary charge
It is possible formally to consider the Reissner-Nordstrom geometry with imaginary charge Q
Q2 < 0 . (8.21)
This is completely unphysical. If charge were imaginary, then electromagnetic energy would be negative.
However, the Reissner-Nordstrom metric with Q2 < 0 is well-defined, and it is possible to calculate
geodesics in that geometry. What makes the geometry interesting is that the singularity, instead of being
gravitationally repulsive, becomes gravitationally attractive. Thus particles, instead of bouncing off the
singularity, are attracted to it, and it turns out to be possible to continue geodesics through the singularity.
Mathematically, the geometry can be considered as the Kerr-Newman geometry in the limit of zero spin. In
8.15 Reissner-Nordstrom geometry with imaginary charge 107
r = ∞
r =∞
r =∞
r = ∞
r =−∞
r = −∞
r = −∞
r =−∞
SingularityPara
llel Ant
ihor
izon Antihorizon
Parallel Universe Universe
Parallel HorizonHor
izon
Black Hole
Singularity (r = 0)Black Hole
Inne
r Horiz
on
Parallel Inner Horizon
AntiverseParallel
Antiverse
White Hole
SingularityWhite Hole
Inner Antihorizon
Para
llel In
ner Ant
ihor
izon
New Parallel Universe New Universe
i
i
i
Figure 8.4 Penrose diagram of the Reissner-Nordstrom geometry with imaginary charge Q. If charge wereimaginary, then electromagnetic energy would be negative, which is completely unphysical. But the metricis well-defined, and the spacetime is fun.
the Kerr-Newman geometry, geodesics can pass from positive to negative radius r, and the passage through
the singularity of the Reissner-Nordstrom geometry can be regarded as this process in the limit of zero spin.
Suffice to say that it is intriguing to see what it looks like to pass through the singularity of a charged
108 Reissner-Nordstrom Black Hole
black hole of imaginary charge, however unrealistic. The Penrose diagram is even more eventful than that
for the usual Reissner-Nordstrom geometry.
9
Kerr-Newman Black Hole
The geometry of a stationary, rotating, uncharged black hole in asymptotically flat empty space was dis-
covered unexpectedly by Roy Kerr in 1963. Kerr’s (2007) own account of the history of the discovery is at
http://arxiv.org/abs/0706.1109. You can read in that paper that the discovery was not mere chance:
Kerr used sophisticated mathematical methods to make it. The extension to a rotating electrically charged
black hole was made shortly thereafter by Ted Newman (Newman et al. 1965). Newman told me (private
communication 2009) that, after seeing Kerr’s work, he quickly realized that the extension to a charged black
hole was straightforward. He set the problem to the graduate students in his relativity class, who became
coauthors of Newman et al. (1965).
The importance of the Kerr-Newman geometry stems in part from the no-hair theorem, which states
that this geometry is the unique end state of spacetime outside the horizon of an undisturbed black hole in
asymptotically flat space.
9.1 Boyer-Lindquist metric
The Boyer-Linquist metric of the Kerr-Newman geometry is
ds2 = − ∆
ρ2
(
dt− a sin2θ dφ)2
+ρ2
∆dr2 + ρ2dθ2 +
R4 sin2θ
ρ2
(
dφ− a
R2dt)2
(9.1)
where R and ρ are defined by
R ≡√
r2 + a2 , ρ ≡√
r2 + a2 cos2θ , (9.2)
and ∆ is the horizon function defined by
∆ ≡ R2 − 2Mr +Q2 . (9.3)
At large radius r, the Boyer-Linquist metric is
ds2 → −(
1− 2M
r
)
dt2 +
(
1 +2M
r
)
dr2 + r2(
dθ2 + sin2θ dφ2)
− 4aM sin2θ
rdtdφ . (9.4)
110 Kerr-Newman Black Hole
Comparison of this metric to the metric of a weak field establishes that M is the mass of the black hole and
a is its angular momentum per unit mass. For positive a, the black hole rotates right-handedly about its
polar axis θ = 0.
The Boyer-Linquist line-element (9.1) defines not only a metric but also a tetrad. The Boyer-Linquist
coordinates and tetrad are carefully chosen to exhibit the symmetries of the geometry. In the locally inertial
frame defined by the Boyer-Linquist tetrad, the energy-momentum tensor (which is non-vanishing for charged
Kerr-Newman) and the Weyl tensor are both diagonal. These assertions becomes apparent only in the tetrad
frame, and are obscure in the coordinate frame.
9.2 Oblate spheroidal coordinates
Boyer-Linquist coordinates r, θ, φ are oblate spheroidal coordinates (not polar coordinates). Corresponding
Cartesian coordinates are
x = R sin θ cosφ ,
y = R sin θ sinφ ,
z = r cos θ .
(9.5)
Surfaces of constant r are confocal oblate spheroids, satisfying
x2 + y2
r2 + a2+z2
r2= 1 . (9.6)
Equation (9.6) implies that the spheroidal coordinate r is given in terms of x, y, z by the quadratic equation
r4 − r2(x2 + y2 + z2 − a2)− a2z2 = 0 . (9.7)
9.3 Time and rotation symmetries
The Boyer-Linquist metric coefficients are independent of the time coordinate t and of the azimuthal angle
φ. This shows that the Kerr-Newman geometry has time translation symmetry, and rotational symmetry
about its azimuthal axis. The time and rotation symmetries means that the tangent vectors gt and gφ in
Boyer-Linquist coordinates are Killing vectors. It follows that their scalar products
gt · gt = gtt = − 1
ρ2
(
∆− a2 sin2θ)
,
gt · gφ = gtφ = − a sin2θ
ρ2
(
R2 −∆)
,
gφ · gφ = gφφ =sin2θ
ρ2
(
R4 − a2 sin2θ∆)
, (9.8)
9.4 Ring singularity 111
are all gauge-invariant scalar quantities. As will be seen below, gtt = 0 defines the boundary of ergospheres,
gtφ = 0 defines the turnaround radius, and gφφ = 0 defines the boundary of the toroidal region containing
closed timelike curves.
The Boyer-Linquist time t and azimuthal angle φ are arranged further to satisfy the condition that gt and
gφ are each orthogonal to both gr and gθ.
9.4 Ring singularity
The Kerr-Newman geometry contains a ring singularity where the Weyl tensor (9.21) diverges, ρ = 0, or
equivalently at
r = 0 and θ = π/2 . (9.9)
The ring singularity is at the focus of the confocal ellipsoids of the Boyer-Linquist metric. Physically, the
singularity is kept open by the centrifugal force.
9.5 Horizons
The horizon of a Kerr-Newmman black hole rotates, as observed by a distant observer, so it is incorrect to
try to solve for the location of the horizon by assuming that the horizon is at rest. The worldline of a photon
that sits on the horizon, battling against the inflow of space, remains at fixed radius r and polar angle θ, but
it moves in time t and azimuthal angle φ. The photon’s 4-velocity is vµ = vt, 0, 0, vφ, and the condition
that it is on a null geodesic is
0 = vµvµ = gµνv
µvν = gtt(vt)2 + 2 gtφ v
tvφ + gφφ(vφ)2 . (9.10)
This equation has solutions provided that the determinant of the 2× 2 matrix of metric coefficients in t and
φ is less than or equal to zero (why?). The determinant is
gttgφφ − g2tφ = − sin2θ∆ (9.11)
where ∆ is the horizon function defined above, equation (9.3). Thus if ∆ ≥ 0, then there exist null geodesics
such that a photon can be instantaneously at rest in r and θ, whereas if ∆ < 0, then no such geodesics exist.
The boundary
∆ = 0 (9.12)
defines the location of horizons. With ∆ given by equation (9.3), equation (9.12) gives outer and inner
horizons at
r± = M ±√
M2 −Q2 − a2 . (9.13)
Between the horizons ∆ is negative, and photons cannot be at rest. This is consistent with the picture that
space is falling faster than light between the horizons.
112 Kerr-Newman Black Hole
Ergosphere
CTCs
Rot
atio
nax
is
Inner horizon
Outer horizon
r = 0Ringsingularity
Ergosphere
CTCs
Rot
atio
nax
is
Inner horizon
Outer horizon
Turnaround
r = 0Ringsingularity
Figure 9.1 Geometry of (upper) a Kerr black hole with spin parameter a = 0.96M , and (lower) a Kerr-Newman black hole with charge Q = 0.8M and spin parameter a = 0.56M . The upper half of each diagramshows r ≥ 0, while the lower half shows r ≤ 0, the Antiverse. The outer and inner horizons are confocaloblate spheroids whose focus is the ring singularity. For the Kerr geometry, the turnaround radius is atr = 0. CTCs are closed timelike curves.
9.6 Angular velocity of the horizon 113
Figure 9.2 Not a mouse’s eye view of a snake coming down its mousehole, uhoh. Contours of constant ρ,and their normals, in Boyer-Linquist coordinates, in a Kerr black hole of spin parameter a = 0.96M . Thethicker contours are the outer and inner horizons, which are confocal spheroids with the ring singularity attheir focus.
9.6 Angular velocity of the horizon
The Boyer-Linquist metric (9.1) has been cunningly written so that you can read off the angular velocity of
the horizon as observed by observers at rest at infinity. The horizon is at dr = dθ = 0 and ∆ = 0, and then
the null condition ds2 = 0 implies that the angular velocity is
dφ
dt=
a
R2. (9.14)
The derivative is with respect to the proper time t of observers at rest at infinity, so this is the angular
velocity observed by such observers.
9.7 Ergospheres
There are finite regions, just outside the outer horizon and just inside the inner horizon, within which the
worldline of an object at rest, dr = dθ = dφ = 0, is spacelike. These regions, called ergospheres, are places
where nothing can remain at rest (the place where little children come from). Objects can escape from within
the outer ergosphere (whereas they cannot escape from within the outer horizon), but they cannot remain
114 Kerr-Newman Black Hole
at rest there. A distant observer will see any object within the outer ergosphere being dragged around by
the rotation of the black hole. The direction of dragging is the same as the rotation direction of the black
hole in both outer and inner ergospheres.
The boundary of the ergosphere is at
gtt = 0 (9.15)
which occurs where
∆ = a2 sin2θ . (9.16)
Equation (9.16) has two solutions, the outer and inner ergospheres. The outer and inner ergospheres touch
respectively the outer and inner horizons at the poles, θ = 0 and π.
9.8 Antiverse
The surface at zero radius, r = 0, forms a disk bounded by the ring singularity. Objects can pass through
this disk into the region at negative radius, r < 0, the Antiverse.
The Boyer-Lindquist metric (9.1) is unchanged by a symmetry transformation that simultaneously flips
the sign both of the radius and mass, r → −r and M → −M . Thus the Boyer-Linquist geometry at
negative r with positive mass is equivalent to the geometry at positive r with negative mass. In effect, the
Boyer-Linquist metric with negative r describes a rotating black hole of negative mass
M < 0 . (9.17)
9.9 Closed timelike curves
Inside the inner horizon there is a toroidal region around the ring singularity within which the light cone in
t-φ coordinates opens up to the point that φ as well as t are timelike coordinates. The direction of increasing
proper time along t is t increasing, and along φ is φ decreasing, which is retrograde. Within the toroidal
region, there exist timelike trajectories that go either forwards or backwards in coordinate time t as they wind
retrograde around the toroidal tunnel. Because the φ coordinate is periodic, these timelike curves connect
not only the past to the future (the usual case), but also the future to the past, which violates causality. In
particular, as first pointed out by Carter (1968), there exist closed timelike curves (CTCs), trajectories
that connect to themselves, connecting their own future to their own past, and repeating interminably, like
Sisyphus pushing his rock up the mountain.
The boundary of this toroidal region is at
gφφ = 0 (9.18)
which occurs whereR4
∆= a2 sin2θ . (9.19)
9.9 Closed timelike curves 115
r = ∞
r =∞
r =∞
r = ∞
r =−∞
r = −∞
r = −∞
r =−∞
White HolePa
ralle
l Antih
oriz
on Antihorizon
Parallel Universe Universe
Parallel HorizonHor
izon
Black Hole
Inne
r Horiz
on
Parallel Inner Horizon
Wor
mho
le ParallelW
ormhole
AntiverseParallel
Antiverse
White HoleInner Antihorizon
Para
llel In
ner Ant
ihor
izon
New Parallel Universe New Universe
Figure 9.3 Penrose diagram of the Kerr-Newman geometry. The diagram is similar to that of the Reissner-Nordstrom geometry, except that it is possible to pass through the disk at r = 0 from the Wormhole regioninto the Antiverse region. This Penrose diagram, which represents a slice at fixed θ and φ, does not capturethe full richness of the geometry, which contains closed timelike curves in a torus around the ring singularity.
116 Kerr-Newman Black Hole
In the uncharged Kerr geometry the CTC torus is entirely at negative radius, r < 0, but in the Kerr-Newman
geometry the CTC torus extends to positive radius.
9.10 Energy-momentum tensor
The Einstein tensor of the Kerr-Newman geometry in Boyer-Linquist coordinates is a bit of a mess, so I won’t
write it down. The trick of raising one index, which for the Reissner-Nordstrom metric brought the Einstein
tensor to diagonal form, equation (8.4), fails for Boyer-Linquist (because the Boyer-Linquist metric is not
diagonal). The problem is endemic to the coordinate approach to general relativity. When we have done
tetrads we will find that, in the Boyler-Linquist tetrad, the Einstein tensor is diagonal, and that the proper
density ρ, the proper radial pressure pr, and the proper transverse pressure p⊥ in that frame are (do not
confuse the notation ρ for proper density with the radial parameter ρ, equation (9.2), of the Boyer-Linquist
metric)
ρ = −pr = p⊥ =Q2
8πρ4. (9.20)
This looks like the energy-momentum tensor (8.4) of the Reissner-Nordstrom geometry with the replacement
r → ρ. The energy-momentum is that of an electric field produced by a charge Q seemingly located at the
singularity.
9.11 Weyl tensor
The Weyl tensor of the Kerr-Newman geometry in Boyer-Linquist coordinates is likewise a mess. After
tetrads, we will find that the 10 components of the Weyl tensor can be decomposed into 5 complex components
of spin 0, ±1, and ±2. In the Boyer-Linquist tetrad, the only non-vanishing component is the spin-0
component, the Weyl scalar C, but in contrast to the Schwarzschild and Reissner-Nordstrom geometries the
spin-0 component is complex, not real:
C = − 1
(r − ia cos θ)3
(
M − Q2
r + ia cos θ
)
. (9.21)
9.12 Electromagnetic field
The expression for the electromagnetic field in Boyer-Linquist coordinates is again a mess. After tetrads, we
will discover that, in the Boyer-Linquist tetrad, the electromagnetic field is purely radial, and the electro-
magnetic potential has only a time component. For subsequent reference, the electromagnetic potential Aµ
in the Boyer-Linquist coordinate (not tetrad) frame is
Aµ =Qr
ρ2
1, 0, 0, −a sin2θ
. (9.22)
9.13 Doran coordinates 117
9.13 Doran coordinates
For the Kerr-Newman geometry, the analog of the Gullstrand-Painleve metric is the Doran (2000) metric
ds2 = − dt2ff +
[
ρ
Rdr − βR
ρ
(
dtff − a sin2θdφff
)
]2
+ ρ2dθ2 +R2 sin2θ dφ2ff (9.23)
where the free-fall time tff and azimuthal angle φff are related to the Boyer-Linquist time t and azimuthal
angle φ by
dtff = dt− β
1− β2dr , dφff = φ− aβ
R2(1− β2)dr . (9.24)
The free-fall time tff is the proper time experienced by persons who free-fall from rest at infinity, with zero
angular momentum. They follow trajectories of fixed θ and φff , with radial velocity dr/dtff = β. In other
words, the 4-velocity uν ≡ dxν/dτ of such free-falling observers is
utff = 1 , ur = β , uθ = 0 , uφff = 0 . (9.25)
For the Kerr-Newman geometry, the velocity β is
β = ∓√
2Mr −Q2
R(9.26)
where the ∓ sign is − (infalling) for black hole solutions, and + (outfalling) for white hole solutions.
Horizons occur where the magnitude of the velocity β equals the speed of light
β = ∓1 . (9.27)
The boundaries of ergospheres occur where the velocity is
β = ∓ ρR. (9.28)
The turnaround radius is where the velocity is zero
β = 0 . (9.29)
The region containing closed timelike curves is bounded by the imaginary velocity
β = iρ
a sin θ. (9.30)
9.14 Extremal Kerr-Newman geometry
The Kerr-Newman geometry is called extremal when the outer and inner horizons coincide, r+ = r−, which
occurs where
M2 = Q2 + a2 . (9.31)
118 Kerr-Newman Black Hole
Rot
atio
nax
is
Inner horizon
Outer horizon
Figure 9.4 Geometry of a Kerr black hole with spin parameter a = 0.96M . The arrows show the velocityβ in the Doran metric. The flow follows lines of constant θ, which form nested hyperboloids orthogonal toand confocal with the nested spheroids of constant r.
Figure 9.5 illustrates the structure of an extremal Kerr (uncharged) black hole, and an extremal Kerr-Newman
(charged) black hole.
9.15 Trajectories of test particles in the Kerr-Newman geometry
Geodesics of test particles in the Kerr-Newman geometry have the expected three constants of motion
associated with conservation of energy, conservation of azimuthal angular momentum, and conservation of
rest mass. Remarkably, Carter (1968) was able to show by separation of variables in the Hamilton-Jacobi
equation that a fourth integral of motion exists, the Carter integral, so that there is a complete set of
four integrals of motion. Moreover, the complete set of integrals exists not only for uncharged particles
(geodesics), but also for charged particles.
The Hamilton-Jacobi method is the most powerful known method for solving equations of motion. For a
test particle it starts with conservation of rest mass m,
gµνpµpν = −m2 , (9.32)
where pµ ≡ mdxµ/dτ is the usual coordinate 4-momentum of the particle. Hamilton-Jacobi replaces the
usual momentum pµ in favour of the generalized momentum πµ = pµ + qAµ, so that the equation (9.32) of
9.15 Trajectories of test particles in the Kerr-Newman geometry 119
Ergosphere
CTCs
Rot
atio
nax
is
Horizon
r = 0Ringsingularity
Ergosphere
CTCs
Rot
atio
nax
is
Horizon
Turnaround
r = 0Ringsingularity
Figure 9.5 Geometry of (upper) an extremal (a = M) Kerr black hole, and (lower) an extremal Kerr-Newmanblack hole with charge Q = 0.8M and spin parameter a = 0.6M .
conservation of rest mass becomes
gµν (πµ − qAµ) (πν − qAν) = −m2 . (9.33)
120 Kerr-Newman Black Hole
Finally, Hamilton-Jacobi replaces the generalized momentum with its derivative with respect to the action,
πµ = ∂S/∂xµ,
gµν
(
∂S
∂xµ− qAµ
)(
∂S
∂xν− qAν
)
= −m2 . (9.34)
Equation (9.34) is the Hamilton-Jacobi equation for a test particle of rest mass m and charge q moving in a
prescribed background with metric gµν and electromagnetic potential Aµ.
For the Boyer-Linqust geometry, two integrals of motion follow immediately from the fact that the metric
is independent of time t and azimuthal angle φ. These imply conservation of energy E and azimuthal angular
momentum Lz,
πt =∂S
∂t= −E , πφ =
∂S
∂φ= Lz . (9.35)
Thus for Boyer-Linquist the Hamiton-Jacobi equation (9.34) becomes
gtt(−E− qAt)2 + 2gtφ(−E− qAt)(Lz − qAφ)+ gφφ(Lz− qAφ)2 + grr
(
∂S
∂r
)2
+ gθθ
(
∂S
∂θ
)2
= −m2 . (9.36)
Substituting in the Boyer-Linquist metric and the electromagnetic potential (9.22) brings equation (9.36) to
− P2
∆+
(
aE sin θ − Lz
sin θ
)2
+ ∆
(
∂S
∂r
)2
+
(
∂S
∂θ
)2
= −m2ρ2 , (9.37)
where P , a function of radius r, is
P ≡ ER2 − aLz − qQr . (9.38)
Separating variables in equation (9.37) gives
−∆
(
∂S
∂r
)2
+P 2
∆−m2r2 =
(
∂S
∂θ
)2
+
(
aE sin θ − Lz
sin θ
)2
+m2a2 cos2θ = K , (9.39)
with K a separation constant. The separation constant K provides the desired fourth integral of motion, and
the solution is now essentially complete. The separated equation (9.39) implies
πr =∂S
∂r=
√R
∆, (9.40a)
πθ =∂S
∂θ=√
Θ , (9.40b)
where R and Θ are radial and angular potentials
R ≡ P 2 −(
K +m2r2)
∆ , (9.41a)
Θ ≡ K −(
aE sin θ − Lz
sin θ
)2
−m2a2 cos2θ . (9.41b)
The middle expression of equation (9.39) shows that the constant K is necessarily positive, hitting zero for
9.15 Trajectories of test particles in the Kerr-Newman geometry 121
a massless particle, m = 0, moving along the polar axis, Lz = 0 and θ = 0 or π. However, it is common to
replace the constant K by the Carter integral Q
K = Q+ (aE − Lz)2 , (9.42)
which has the property that Q = 0 for orbits in the equatorial plane, θ = π/2. In terms of the Carter integral
Q, the radial and angular potentials R and Θ are
R = P 2 −[
Q+ (aE − Lz)2 +m2r2
]
∆ , (9.43a)
Θ = Q− cos2θ
[
a2(m2 − E2) +L2
z
sin2θ
]
. (9.43b)
Converting the generalized covariant momenta, equations (9.35) and (9.40), to ordinary contravariant mo-
menta, pµ = gµν(πν − qAµ), yields
pt =1
ρ2
[
PR2
∆− a
(
aE sin2θ − Lz
)
]
, (9.44a)
pr =1
ρ2
√R , (9.44b)
pθ =1
ρ2
√Θ , (9.44c)
pφ =1
ρ2
[
aP
∆− aE +
Lz
sin2θ
]
. (9.44d)
Equations (9.44) give the general solution for the 4-momentum pµ of a test particle of mass m and charge q
moving in the Kerr-Newman geometry, in terms of its integrals of motion E, Lz, and Q.
The following exercise shows that among ideal black holes, the Schwarzschild geometry is exceptional, not
typical, in having a gravitationally attractive singularity.
Exercise 9.1 Near the Kerr-Newman ring singularity. Explore the behaviour of trajectories of test
particles in the vicinity of the Kerr-Newman singularity, where ρ → 0. Under what conditions does a test
particle reach the singularity?
1. Argue that for a particle to reach the singularity at θ = π/2, positivity of the angular potential Θ requires
that
Q ≥ 0 . (9.45)
2. Argue that for a particle to reach the singularity at r = 0, positivity of the radial potential R requires
that
Q2(aE − Lz)2 + (Q2 + a2)Q ≤ 0 . (9.46)
3. Schwarzschild case: show that if Q = 0 and a = 0, then a particle reaches the singularity provided that
the mass of the black hole is positive, M > 0.
122 Kerr-Newman Black Hole
4. Reissner-Nordstrom case: show that if Q2 > 0 and a = 0, then a particle can reach the singularity only if
it has zero angular momentum, Q = Lz = 0, and if the particle’s charge-to-mass exceeds unity,
q2
m2≥ 1 . (9.47)
5. Kerr case: show that if Q = 0 but a2 > 0, then a particle can reach the singularity only if Q = 0, and
provided that the mass of the black hole is positive, M > 0.
6. Kerr-Newman case: show that if Q2 > 0 and a2 > 0, then a particle can reach the singularity only if
Lz = aE and Q = 0, and if the particle’s charge-to-mass is large enough,
q2
m2≥ Q2 + a2
Q2, (9.48)
which generalizes the Reissner-Nordstrom condition (9.47).
9.16 Penrose process
Trajectories in the Kerr-Newman geometry can have negative energy E outside the horizon. It is possible to
reduce the mass M of the black hole by dropping negative energy particles into the black hole. This process
of extracting mass-energy from the black hole is called the Penrose process.
Exercise 9.2 Negative energy trajectories outside the horizon. Under what conditions can test
particles have negative energy trajectories, E < 0, outside the horizon?
1. Argue that outside the horizon, the positivity of the horizon function ∆, and of the radial and angular
potentials R and Θ, equations (9.41) implies that P , equation (9.38), satisfies
P 2 ≥(
K +m2r2)
∆ ≥[
(
aE sin θ − Lz
sin θ
)2
+m2ρ2
]
∆ . (9.49)
2. Argue that the condition (9.49) implies by continuity that for a massive particle P must be strictly positive
outside the horizon. Extend your argument to a massless particle by taking a massless particle as a massive
particle in the limit of large energy.
3. Argue that the positivity of P implies that aLz + qQr must be negative for the energy E to be negative.
Show that, more stringently, negative E requires that
aLz + qQr ≤ −√
(
L2z
sin2θ+m2ρ2
)
∆ . (9.50)
4. Argue that for an uncharged particle, q = 0, negative energy trajectories exist only inside the ergosphere.
5. Do negative energy trajectories exist outside the ergosphere for a charged particle?
9.17 Constant latitude trajectories in the Kerr-Newman geometry 123
6. For the Penrose process to work, the negative energy particle must fall through the horizon, where ∆ = 0.
Does this happen?
Exercise 9.3 When can objects go forwards or backwards in time t?
9.17 Constant latitude trajectories in the Kerr-Newman geometry
A trajectory is at constant latitude if it is at constant polar angle θ,
θ = constant . (9.51)
Constant latitude orbits occur where the angular potential Θ, equation (9.43b), not only vanishes, but is an
extremum,
Θ =dΘ
dθ= 0 , (9.52)
the derivative being taken with the constants of motion E, Lz, and Q of the orbit being held fixed. The
condition Θ = 0 simply sets the value of the Carter integral Q. Solving dΘ/dθ = 0 yields the condition
between energy E and angular momentum Lz
E = ±√
1 +L2
z
a2 sin4θ. (9.53)
Solutions at any polar angle θ and any angular momentum Lz exist, ranging from E = ±1 at Lz = 0, to
E = ±Lz/(a sin2θ) at Lz → ±∞. The solutions with Lz = 0 are those of the freely-falling observers that
define the Doran coordinate system, §9.13. The solutions with Lz →∞ define the principal null congruences
discussed in §9.18.
9.18 Principal null congruence
A congruence is a space-filling, non-overlapping set of geodesics. In the Kerr-Newman geometry there is
a special set of null geodesics, the ingoing and outgoing principal null congruences, with respect to
which the symmetries of the geometry are especially apparent. Photons that hold steady on the horizon are
members of the outgoing principal null congruence. The energy-momentum tensor is diagonal in a locally
inertial frame aligned with the ingoing or outgoing principal null congruence. The Weyl tensor, decomposed
into spin components in the locally inertial frame of the principal null congruences, contains only spin-0
components.
The Boyer-Linquist metric (9.1) is specifically constructed so that the Boyer Linquist tetrad is aligned with
the principal null tetrad. Along the principal null congruences, the final two terms of the Boyer-Linquist
124 Kerr-Newman Black Hole
line element (9.1) vanish
dθ = dφ− a
R2dt = 0 . (9.54)
Solving the null condition ds2 = 0 on the rest of the metric yields the photon 4-velocity vµ ≡ dxµ/dλ on the
principal null congruences
vt =R2
∆, vr = ±1 , vθ = 0 , vφ =
a
∆. (9.55)
In the regions outside the outer horizon or inside the inner horizon, the ± sign in front of vr is + for outgoing,
− for ingoing geodesics. Between the outer and inner horizons, vr is negative in the Black Hole region, and
positive in the White Hole region, while vt and vφ are negative for ingoing, positive for outgoing geodesics.
The angular momentum per unit energy Jz ≡ Lz/|E| of photons along the principal null congruences is not
zero, but is
Jz = a sin2θ (9.56)
with the same sign for both ingoing and outgoing geodesics.
9.19 Circular orbits in the Kerr-Newman geometry
An orbit can be termed circular if it is at constant radius r,
r = constant . (9.57)
It is convenient to call such an orbit circular even if the orbit is at finite inclination (not confined to the
equatorial plane) about a rotating black hole, and therefore follows the surface of a spheroid (in Boyer-
Lindquist coordinates).
Orbits turn around in r, reaching periapsis or apoapsis, where the radial potential R, equation (9.43a),
vanishes. Circular orbits occur where the radial potential R not only vanishes, but is an extremum,
R =dRdr
= 0 , (9.58)
the derivative being taken with the constants of motion E, Lz, and Q of the orbit being held fixed. Circular
orbits may be either stable or unstable. The stability of a circular orbit is determined by the sign of the
second derivative of the potential
d2Rdr2
, (9.59)
with + for stable, − for unstable circular orbits. Marginally stable orbits occur where d2R/dr2 = 0.
Circular orbits occur not only in the equatorial plane, but at general inclinations. The inclination of an
orbit can be characterized by the minimum polar angle θmin to which it extends. An astronomer would call
π/2− θmin the inclination angle of the orbit. It is convenient to define an inclination parameter α by
α ≡ cos2θmin , (9.60)
9.19 Circular orbits in the Kerr-Newman geometry 125
which lies in the interval [0, 1]. Equatorial orbits, at θ = π/2, correspond to α = 0, while polar orbits, those
that go over the poles at θ = 0 and π, correspond to α = 1.
9.19.1 General solution for circular orbits
The general solution for circular orbits of a test particle of arbitrary electric charge q in the Kerr-Newman
geometry is as follows.
The rest mass m of the test particle can be set equal to unity, m = 1, without loss of generality. Circular
orbits of particles with zero rest mass, m = 0, discussed later in this section, occur in cases where the circular
orbits for massive particles attain infinite energy and angular momentum.
In the radial potential R, equation (9.43a), eliminate the Carter integral Q in favour of the inclination
parameter α, equation (9.60), using equation (9.43b)
Q = α
[
a2(1− E2) +L2
z
1− α
]
. (9.61)
Furthermore, eliminate the energy E in favour of P , equation (9.38). The radial derivatives dnR/drn must
be taken before E is replaced by P , since E is a constant of motion, whereas P varies with r. The physical
motivation for replacing E with P lies in the sign of P . Solutions with positive P correspond to orbits in the
Universe, Wormhole, or Antiverse parts of the Kerr-Newman geometry in the Penrose diagram of Figure 9.3,
while solutions with negative P correspond to orbits in their Parallel counterparts. If only the Universe
region is considered, then P is necessarily positive. By contrast, the energy E can be either positive or
negative in the same region of the Kerr-Newman geometry (the energy E is negative for orbits of sufficiently
large negative angular momentum Lz inside the ergosphere of the Universe).
The condition R = 0 is a quadratic equation in Lz, whose solutions are
Lz =1
r2 + a2α
[
a(1− α)(P + qQr)±R2√
(1− α) [P 2/∆− (r2 + a2α)]]
. (9.62)
Substituting the two (±) expressions (9.62) for Lz into dR/dr, and setting the product of the resulting two
expressions for dR/dr equal to zero, yields a quartic equation for P/∆:
p0 + p1(P/∆) + p2(P/∆)2 + p3(P/∆)3 + p4(P/∆)4 = 0 , (9.63)
with coefficients
p0 ≡ r2(r2 + a2α)2 , (9.64a)
p1 ≡ − 2qQr(r2 − a2α)(r2 + a2α) , (9.64b)
p2 ≡ − 2r2(r2 + a2α)(r2 − 3Mr + 2Q2 + a2α+ a2αM/r) + q2Q2(r2 − a2α)2 , (9.64c)
p3 ≡ 2qQr(r2 − a2α)(r2 − 3Mr + 2Q2 + 2a2 − a2α+ a2αM/r) , (9.64d)
p4 ≡[
r6 − 6Mr5 + (9M2+4Q2+2a2α)r4 − 4M(3Q2+a2)r3
+ (4Q4−6a2αM+4a2Q2+a4α2)r2 + 2a2α(2Q2+2a2−a2α)Mr + a4α2M2]
. (9.64e)
126 Kerr-Newman Black Hole
The quartic (9.63) is the condition for an orbit at radius r to be circular. Physical solutions must be real.
The quartic (9.63) has either zero, two, or four real solutions at any one radius r. Numerically, it is better
to solve the quartic (9.63) for the reciprocal ∆/P rather than P/∆, since the vanishing of 1/P defines the
location of circular orbits of massless particles.
−3 −2 −1 0 1 2 3 4 5 6 7 8 9−3
−2
−1
0
1
2
3
Radius r/M
∆/P
Outer
horizon
Innerhorizon
Figure 9.6 Values of ∆/P for circular orbits at radius r of a charged particle about a Kerr-Newman blackhole. The values ∆/P are real roots of the quartic (9.63); there are either zero, two, or four real roots atany one radius. The parameters are representative: a particle of charge-to-mass q/m = 2.4 on an orbit ofinclination parameter α = 0.5 about a black hole of charge Q = 0.5M and spin parameter a = 0.5M . Solid(green) lines indicate stable orbits; dashed (brown) lines indicate unstable orbits. Positive ∆/P orbits occurin Universe, Wormhole, and Antiverse regions; negative ∆/P orbits occur in their Parallel counterparts;zero ∆/P orbits are null. The fact that the particle is charged breaks the symmetry between positive andnegative ∆/P . If the charge of the particle were flipped, q/m = −2.4, then the diagram would be reflectedabout the horizontal axis (the sign of ∆/P would flip).
The angular momentum Lz, energy E, and stability d2R/dr2 of a circular orbit are, in terms of a solution
P/∆ of the quartic (9.63),
Lz = ± 1
r2 + a2α
√
(1− α) [l−1(∆/P ) + l0 + l1(P/∆) + l2(P/∆)2] , (9.65a)
E =1
2[(∆/P ) + qQ/r + (1−M/r)(P/∆)] , (9.65b)
d2Rdr2
=2
(r2 + a2α)2[
q−1(∆/P ) + q0 + q1(P/∆) + q2(P/∆)2]
, (9.65c)
9.19 Circular orbits in the Kerr-Newman geometry 127
where the coefficients li and qi are
l−1 = qQrR2(r2 + a2α) , (9.66a)
l0 = −R2(r2 + a2α)(2Mr −Q2)− q2Q2(r4 − a4α) , (9.66b)
l1 = − qQr[
2r4 − 5Mr3 + 3(Q2+a2)r2 − a2(1+α)Mr + a2(Q2+αQ2+a2−a2α)
+ 3a4αM/r − a4α(Q2+a2)/r2]
, (9.66c)
l2 =[
3Mr3 − 2Q2r2 + a2(1+α)Mr − a2(1+α)Q2 − a4αM/r]
∆ , (9.66d)
and
q−1 = 2qQr(r2 − a2α)(r2 + a2α) , (9.67a)
q0 = − 4(r2 + a2α)(Mr3 −Q2r2 − a2αMr)− q2Q2(r2 − a2α)2 , (9.67b)
q1 = − qQr[
r4 − 4Mr3 + 3(Q2+a2−2a2α)r2 + 12a2αMr − a2α(6Q2+6a2−a2α)
− a4α2(Q2+a2)/r2]
, (9.67c)
q2 =(
3Mr3 − 4Q2r2 − 6a2αMr − a4α2M/r)
∆ . (9.67d)
The sign of the angular momentum Lz in equation (9.65a) should be chosen such that the relations (9.38)
for P and (9.65b) for E hold. This choice of sign becomes ambiguous for a = 0; but this is as it should be,
since either sign of Lz is valid for a = 0, where the black hole is spherically symmetric, and therefore defines
no preferred direction.
The expressions (9.62) and (9.65a) for Lz are equal on a circular orbit. The advantage of the latter
expression (9.65a) will become apparent below, where it is found that for particles of zero electric charge,
q = 0, one circular orbit is always prograde, aLz > 0, while the other is always retrograde, aLz < 0.
For non-zero a, the reality of a solution P/∆ of the quartic (9.63) is a necessary and sufficient condition for
a corresponding circular orbit to exist. In particular, the argument of the square root in the expression (9.65a)
for Lz is guaranteed to be positive. For zero a, however, the quartic (9.65a), which reduces in this case to
the square of a quadratic, admits real solutions that do not correspond to a circular orbit. For these invalid
solutions, the argument of the square root in the expression (9.65a) for Lz is negative. Thus for zero a, a
necessary and sufficient condition for a circular orbit to exist is that the solutions for both P/∆ and Lz be
real.
9.19.2 Circular orbits for massless particles
Circular orbits for massless particles, m = 0, or null circular orbits, follow from the solutions for massive
particles in the case where the energy and angular momentum on the circular orbit become infinite, which
occurs when P → ±∞. Except at horizons, where ∆ = 0, the solution for P from the quartic (9.63) diverges
when the ratio p4/p0 of the highest to lowest order coefficients vanishes. The ratio p4/p0, equations (9.64),
factors as
p4/p0 =F+F−
(r2 + a2α)2, (9.68)
128 Kerr-Newman Black Hole
where
F± ≡ r2 − 3Mr + 2Q2 + a2α(1 +M/r) ± 2a√
(1− α)(Mr −Q2 − a2αM/r) . (9.69)
A null circular orbit thus occurs at a radius r such that
F+ = 0 or F− = 0 , (9.70)
with + for prograde (aLz > 0) orbits, − for retrograde (aLz < 0) orbits. The location of null circular
orbits are independent of the charge q of the particle, since F± are independent of charge q. The angular
momentum Jz per unit energy on the null circular orbit is, from equations (9.65a) and (9.65b) in the limit
P → ±∞,
Jz ≡ Lz/|E| = ±2√
(1− α)l2(r2 + a2α)2(1−M/r)
. (9.71)
The case where F+ or F− vanishes at a horizon is special. This occurs when the black hole is extremal,
M2 = Q2 + a2. A circular orbit exists at the horizon of an extremal black hole provided that the charge
squared Q2 and inclination parameter α are not too large, the precise condition being
a4α2 + 6(Q2+a2)α− (Q2+a2)(Q2−3a2) ≤ 0 . (9.72)
The circular orbit is non-null, since the vanishing of ∆/P no longer implies that P diverges if ∆ = 0, as is
true on the horizon. A careful analysis shows that the limiting value of P/√
∆ is finite for a circular orbit
at the horizon of an extremal black hole, so in fact P = 0 for such an orbit.
Since there are null geodesics, the ingoing or outgoing principal null geodesics, that hold steady on the
horizon, one might have expected that there would always be solutions for null circular orbits on the horizon,
but this is false. The resolution of the paradox is that massless particles experience no proper time along
their geodesics. If a massive particle is put on the horizon on a relativistic geodesic, then the massive particle
necessarily falls off the horizon in a finite proper time: it is impossible for the geodesic to hold steady on
the horizon. The only exception is that, as discussed in the previous paragraph, an extremal black hole may
have circular orbits at its horizon; but these orbits have P = 0, and are not null.
9.19.3 Circular orbits for particles with zero electric charge
For a particle with zero electric charge, q = 0, the quartic condition (9.63) for a circular orbit reduces to a
quadratic in (P/∆)2. Solving the quadratic for the reciprocal (∆/P )2 yields two possible solutions
(∆/P )2 =F±
r2 + a2α, (9.73)
where F± are defined by equation (9.69), with + for prograde (aLz > 0) orbits, − for retrograde (aLz < 0)
orbits. The sign of P is positive in the Universe, Wormhole, and Antiverse of Figure 9.3, negative in their
Parallel counterparts. For zero electric charge, the expressions (9.65) for the angular momentum Lz, energy
9.19 Circular orbits in the Kerr-Newman geometry 129
E, and stability d2R/dr2 of a circular orbit simplify to
Lz = ± 1
r2 + a2α
√
(1− α) [l0 + l2(P/∆)2] , (9.74a)
E =1
2[(∆/P ) + (1 −M/r)(P/∆)] , (9.74b)
d2Rdr2
=2
(r2 + a2α)2[
q0 + q2(P/∆)2]
. (9.74c)
The coefficients li and qi in equations (9.74) reduce from the expressions (9.66) and (9.67) to
l0 = −R2(r2 + a2α)(2Mr −Q2) , (9.75a)
l2 =[
3Mr3 − 2Q2r2 + a2(1+α)Mr − a2(1+α)Q2 − a4αM/r]
∆ , (9.75b)
and
q0 = − 4(r2 + a2α)(Mr3 −Q2r2 − a2αMr) , (9.76a)
q2 =(
3Mr3 − 4Q2r2 − 6a2αMr − a4α2M/r)
∆ . (9.76b)
.0 .5 1.0 1.5 2.0−2.0
−1.5
−1.0
−.5
.0
.5
1.0
1.5
2.0
Radius r/M
P
Outer
horizon
Innerhorizon
Figure 9.7 Values of P for circular orbits at radius r in the equatorial plane of a near-extremal Kerr blackhole, with black hole spin parameter a = 0.999M . The diagram illustrates that as the orbital radius rapproaches the horizon, P first approaches zero, but then increases sharply to infinity, corresponding to nullcircular orbits. In the case of an exactly extremal black hole, P goes as to zero at the horizon, there is noincrease of P to infinity, and no null circular orbit. Solid (green) lines indicate stable orbits; dashed (brown)lines indicate unstable orbits.
9.19.4 Equatorial circular orbits in the Kerr geometry
The case of greatest practical interest to astrophysicists is that of circular orbits in the equatorial plane of
an uncharged black hole, the Kerr geometry.
130 Kerr-Newman Black Hole
For circular orbits in the equatorial plane, α = 0, of an uncharged black hole, Q = 0, the solution (9.73)
simplifies to
(∆/P )2 =F±r2
(9.77)
where F±, equation (9.69), reduce to
F± ≡ r2 − 3Mr ± 2a√Mr , (9.78)
with + for prograde (aLz > 0) orbits, − for retrograde (aLz < 0) orbits.
As discussed above, null circular orbits occur where F± = 0, except in the special case that the circular
orbit is at the horizon, which occurs when the black hole is extremal. In the limit where the Kerr black hole
is near but not exactly extremal, a → |M |, null circular orbits occur at r → M (prograde) and r → 4M
(retrograde). For an exactly extremal Kerr black hole, a = |M |, the (prograde) circular orbit at the horizon
is no longer null. The situation of a near extremal Kerr black hole is illustrated by Figure 9.7.
It is generally argued that the inner edge of an accretion disk is likely to occur at the innermost stable
equatorial circular orbit. An orbit at this point has marginal stability, d2R/dr2 = 0. Simplifying the stability
d2R/dr2 from equation (9.74c) to the case of equatorial orbits, α = 0, and zero black hole charge, Q = 0,
yields the condition of marginal stability
r2 − 6Mr − 3a2 ± 8a√Mr = 0 . (9.79)
The + (prograde) orbit has the smaller radius, and so defines the innermost stable circular orbit. For an
extremal Kerr black hole, a = |M |, marginally stable circular equatorial orbits are at r = M (prograde) and
r = 9M (retrograde).
9.19.5 Circular orbits in the Reissner-Nordstrom geometry
Circular orbits of particles in the Reissner-Nordstrom geometry follow from those in the Kerr-Newman
geometry in the limit of a non-rotating black hole, a = 0. For a non-rotating black hole, an orbit can be
taken without loss of generality to circulate right-handedly in the equatorial plane, θ = π/2, so that α = 0
and the azimuthal angular momentum Lz equals the positive total angular momentum L. For non-equatorial
orbits, the relation between azimuthal and total angular momentum is Lz = ±√
1− αL.
For a non-rotating black hole, a = 0, the quartic condition (9.63) for a circular orbit of a particle of rest
mass m = 1 and electric charge q reduces to the square of a quadratic,
r2 − qQr(P/∆)−(
r2 − 3Mr + 2Q2)
(P/∆)2 = 0 . (9.80)
Solving the quadratic (9.80) for the reciprocal ∆/P yields two solutions
∆/P =qQ
2r±√
1− 3M
r+
2Q2
r2+q2Q2
4r2. (9.81)
The sign of P is positive in the Universe, Wormhole, and Antiverse parts of the Reissner-Nordstrom geometry
9.19 Circular orbits in the Kerr-Newman geometry 131
in the Penrose diagram of Figure 8.1, negative in their Parallel counterparts. The angular momentum L,
energy E, and stability d2R/dr2 of a circular orbit are, in terms of a solution ∆/P of the quadratic (9.80),
L =√
P 2/∆− r2 , (9.82)
E =P
r2+qQ
r, (9.83)
d2Rdr2
= 2(
r2 − 6Mr + 5Q2 + q2Q2)
− 2
(
1− 6M
r+
6Q2
r2
)
P 2
∆. (9.84)
For massless particles, circular orbits occur where the solution (9.81) for ∆/P vanishes, which occurs when
r2 − 3Mr + 2Q2 = 0 , (9.85)
independent of the charge q of the particle. The condition (9.85) is consistent with the Kerr-Newman
condition for a null circular orbit, the vanishing of F± given by equation (9.69). However, for Kerr-Newman,
the argument of the square root on the right hand side of equation (9.69) for F± must be positive, even in
the limit of infinitesimal a. In the limit of small a, this requires that Mr −Q2 ≥ 0. If the charge Q of the
Reissner-Nordstrom black hole lies in the standard range 0 ≤ Q2 ≤ M2, then one of the solutions of the
quadratic (9.85) lies outside the outer horizon, while the other lies between the outer and inner horizons.
As one might hope, the additional condition Mr −Q2 ≥ 0 eliminates the undesirable solution between the
horizons, leaving only the solution outside the horizon, which is
r =3M
2
(
1 +
√
1− 8Q2
9
)
for 0 ≤ Q2 ≤M2 . (9.86)
In (unphysical) cases Q2 < 0 or M2 < Q2 ≤ (9/8)M2, both solutions of equation (9.85) are valid.
PART FOUR
HOMOGENEOUS, ISOTROPIC COSMOLOGY
Concept Questions
1. What does it mean that the Universe is expanding?
2. Does the expansion affect the solar system or the Milky Way?
3. How far out do you have to go before the expansion is evident?
4. What is the Universe expanding into?
5. In what sense is the Hubble constant constant?
6. Does our Universe have a center, and if so where is it?
7. What evidence suggests that the Universe at large is homogeneous and isotropic?
8. How can the CMB be construed as evidence for homogeneity and isotropy given that it provides information
only about a 2D surface on the sky?
9. What is thermodynamic equilibrium? What evidence suggests that the early Universe was in thermody-
namic equilibrium?
10. What are cosmological parameters?
11. What cosmological parameters can or cannot be measured from the power spectrum of fluctuations of the
CMB?
12. FRW Universes are characterized as closed, flat, or open. Does flat here mean the same as flat Minkowski
space?
13. What is it that astronomers call dark matter?
14. What is the primary evidence for the existence of non-baryonic cold dark matter?
15. How can astronomers detect dark matter in galaxies or clusters of galaxies?
16. How can cosmologists claim that the Universe is dominated by not one but two distinct kinds of mysterious
mass-energy, dark matter and dark energy, neither of which has been observed in the laboratory?
17. What key property or properties distinguish dark energy from dark matter?
18. Does the Universe conserve entropy?
19. Does the annihilation of electron-positron pairs into photons generate entropy in the early Universe, as its
temperature cools through 1 MeV?
20. How does the wavelength of light change with the expansion of the Universe?
21. How does the temperature of the CMB change with the expansion of the Universe?
136 Concept Questions
22. How does a blackbody (Planck) distribution change with the expansion of the Universe? What about a
non-relativistic distribution? What about a semi-relativistic distribution?
23. What is the horizon of our Universe?
24. What happens beyond the horizon of our Universe?
25. What caused the Big Bang?
26. What happened before the Big Bang?
27. What will be the fate of the Universe?
What’s important?
1. The CMB indicates that the early (≈ 400,000 year old) Universe was (a) uniform to a few ×10−5, and (b)
in thermodynamic equilibrium. This indicates that
the Universe was once very simple .
It is this simplicity that makes it possible to model the early Universe with some degree of confidence.
2. The power spectrum of fluctuations of the CMB has enabled precise measurements of cosmological pa-
rameters, excepting the Hubble constant.
3. There is a remarkable concordance of evidence from a broad range of astronomical observations — su-
pernovae, big bang nucleosynthesis, the clustering of galaxies, the abundances of clusters of galaxies,
measurements of the Hubble constant from Cepheid variables, the ages of the oldest stars.
4. Observational evidence is consistent with the predictions of the theory of inflation in its simplest form —
the expansion of the Universe, the spatial flatness of the Universe, the near uniformity of temperature
fluctuations of the CMB (the horizon problem), the presence of acoustic peaks and troughs in the power
spectrum of fluctuations of the CMB, the near power law shape of the power spectrum at large scales, its
spectral index (tilt), the gaussian distribution of fluctuations at large scales.
5. What is non-baryonic dark matter?
6. What is dark energy? What is its equation of state w ≡ p/ρ, and how does w evolve with time?
10
Homogeneous, Isotropic Cosmology
10.1 Observational basis
Since 1998, observations have converged on a Standard Model of Cosmology, a spatially flat universe domi-
nated by dark energy and by non-baryonic dark matter.
1. The Hubble diagram (distance versus redshift) of galaxies indicates that the Universe is expanding (Hubble
1929).
2. The Cosmic Microwave Background (CMB).
• Near black body spectrum, with T0 = 2.725± 0.001 K (Fixsen & Mather 2002).
• Dipole ⇒ the solar system is moving at 365 kms−1 through the CMB.
• After dipole subtraction, the temperature of the CMB over the sky is uniform to a few parts in 105.
• The power spectrum of temperature T fluctuations shows a scale-invariant spectrum at large scales,
and prominent acoustic peaks at smaller scales. Allows measurement of the amplitude As and tilt ns of
primordial fluctuations, the curvature density Ωk, and the proper densities Ωch2 of non-baryonic cold
dark matter and Ωbh2 of baryons. Does not measure Hubble constant h ≡ H0/(100 kms−1 Mpc−1).
• The power spectra of E and B polarization fluctuations, and the various cross power spectra (only T -E
should be non-vanishing).
3. The Hubble diagram of Type Ia (thermonuclear) supernovae indicates that the Universe is accelerating.
This points to the dominance of gravitationally repulsive dark energy, with ΩΛ ≈ 0.75. The amount of
dark energy is consistent with observations from the CMB indicating that the Universe is spatially flat,
Ω ≈ 1, and observations from CMB, galaxy clustering, and clusters of galaxies indicating that the density
in gravitationally attractive matter is only Ωm ≈ 0.25.
4. Observed abundances of light elements H, D, 3He, He, and Li are consistent with the predictions of big
bang nucleosynthesis (BBN) provideed that Ωb ≈ 0.04, in good agreement with measurements from the
CMB.
5. The clustering of matter (dark and bright) shows a power spectrum in good agreement with the Standard
Model:
• galaxies;
• the Lyman alpha forest;
10.2 Cosmological Principle 139
• gravitational lensing.
Historically, the principle evidence for non-baryonic cold dark matter is comparison between the power
spectra of galaxies versus CMB. How can tiny fluctuations in the CMB grow into the observed fluctuations
in matter today in only the age of the Universe? Answer: non-baryonic dark matter that begins to cluster
before Recombination, when the CMB was released.
6. The abundance of galaxy clusters as a function of redshift.
7. The ages of the oldest stars, in globular clusters. The Hubble constant yields an estimate of the age of
the Universe that is older with dark energy than without. The ages of the oldest stars agree with the age
of the Universe with dark enery, but are older than the Universe without dark energy.
8. Ubiquitous evidence for dark matter, deduced from sizes and velocities (or in the case of gravitational
lensing, the gravitational potential) of various objects.
• The Local Group of galaxies.
• Rotation curves of spiral galaxies.
• The temperature and distribution of x-ray gas in elliptical galaxies.
• The temperature and distribution of x-ray gas in clusters of galaxies.
• Gravitational lensing by clusters of galaxies.
9. The Bullet cluster is a rare example that supports the notion that the dark matter is non-baryonic. In the
Bullet cluster, two clusters recently passed through each other. The baryonic matter, as measured from
x-ray emission of hot gas, appears displaced from the dark matter, as measured from weak gravitational
lensing.
10.2 Cosmological Principle
The cosmological principle states that the Universe at large is
• homogeneous (has spatial translation symmetry),
• isotropic (has spatial rotation symmetry).
The primary evidence for this is the uniformity of the temperature of the CMB, which, after subtraction
of the dipole produced by the motion of the solar system through the CMB, is constant over the sky to a
few parts in 105. Confirming evidence is the statistical uniformity of the distribution of galaxies over large
scales.
The cosmological principle allows that the Universe evolves in time, as observations surely indicate — the
Universe is expanding, galaxies, quasars, and galaxy clusters evolve with redshift, and the temperature of
the CMB is undoubtedly decreasing as the Universe expands.
140 Homogeneous, Isotropic Cosmology
10.3 Friedmann-Robertson-Walker metric
Universes satisfying the cosmological principle are described by the Friedmann-Robertson-Walker (FRW)
metric, equation (10.25) below. The metric, and the associated Einstein equations, which are known as the
Friedmann equations, are set forward in the next several sections, §§10.4–10.9.
10.4 Spatial part of the FRW metric: informal approach
The cosmological principle implies that
the spatial part of the FRW metric is a 3D hypersphere (10.1)
where in this context the term hypersphere is to be construed as including not only cases of positive curvature,
which have finite positive radius of curvature, but also cases of zero and negative curvature, which have
infinite and imaginary radius of curvature.
w
r = Rsinχ
R
xy
r// =
RχRdχ
Rsinχdφ
χ
φ
Figure 10.1 Embedding diagram of the FRW geometry.
Figure 10.1 shows an embedding diagram of a 3D hypersphere in 4D Euclidean space. The horizontal
directions in the diagram represent the normal 3 spatial x, y, z dimensions, with one dimension z suppressed,
10.4 Spatial part of the FRW metric: informal approach 141
while the vertical dimension represents the 4th spatial dimension w. The 3D hypersphere is a set of points
x, y, z, w satisfying(
x2 + y2 + z2 + w2)1/2
= R = constant . (10.2)
An observer is sitting at the north pole of the diagram, at 0, 0, 0, 1. A 2D sphere (which forms a 1D circle
in the embedding diagram of Figure 10.1) at fixed distance surrounding the observer has geodesic distance
r‖ defined by
r‖ ≡ proper distance to sphere measured along a radial geodesic , (10.3)
and circumferential radius r defined by
r ≡(
x2 + y2 + z2)1/2
, (10.4)
which has the property that the proper circumference of the sphere is 2πr. In terms of r‖ and r, the spatial
metric is
dl2 = dr2‖ + r2do2 (10.5)
where do2 ≡ dθ2 + sin2θ dφ2 is the metric of a unit 2-sphere.
Introduce the angle χ illustrated in the diagram. Evidently
r‖ = Rχ ,
r = R sinχ . (10.6)
In terms of the angle χ, the spatial metric is
dl2 = R2(
dχ2 + sin2χdo2)
(10.7)
which is one version of the spatial FRW metric. The metric resembles the metric of a 2-sphere of radius R,
which is not surprising since the same construction, with Figure 10.1 interpreted as the embedding diagram
of a 2D sphere in 3D, yields the metric of a 2-sphere. Indeed, the construction iterates to give the metric of
an n-dimensional sphere of arbitrarily many dimensions n.
Instead of the angle χ, the metric can be expressed in terms of the circumferential radius r. It follows
from equations (10.6) that
r‖ = R sin−1(r/R) (10.8)
whence
dr‖ =dr
√
1− r2/R2
=dr√
1−Kr2(10.9)
where K is the curvature
K ≡ 1
R2. (10.10)
142 Homogeneous, Isotropic Cosmology
In terms of r, the spatial FRW metric is then
dl2 =dr2
1−Kr2 + r2do2 . (10.11)
The embedding diagram Figure 10.1 is a nice prop for the imagination, but it is not the whole story. The
curvature K in the metric (10.11) may be not only positive, corresponding to real finite radius R, but also
zero or negative, corresponding to infinite or imaginary radius R. The possibilities are called closed, flat,
and open:
K
> 0 closed R real ,
= 0 flat R→∞ ,
< 0 open R imaginary .
(10.12)
10.5 Comoving coordinates
The metric (10.11) is valid at any single instant of cosmic time t. As the Universe expands, the 3D spatial
hypersphere (whether closed, flat, or open) expands. In cosmology it is highly advantageous to work in
comoving coordinates that expand with the Universe. Why? First, it is helpful conceptually and math-
ematically to think of the Universe as at rest in comoving coordinates. Second, linear perturbations, such
as those in the CMB, have wavelengths that expand with the Universe, and are therefore fixed in comoving
coordinates.
In practice, cosmologists introduce the cosmic scale factor a(t)
a(t) ≡ measure of the size of the Universe, expanding with the Universe (10.13)
which is proportional to but not necessarily equal to the radius R of the Universe. The cosmic scale factor
a can be normalized in any arbitrary way. The most common convention adopted by cosmologists is to
normalize it to unity at the present time,
a0 = 1 , (10.14)
where the 0 subscript conventionally signifies the present time.
Comoving geodesic and circumferential radial distances x‖ and x are defined in terms of the proper geodesic
and circumferential radial distances r‖ and r by
ax‖ ≡ r‖ , ax ≡ r . (10.15)
Objects expanding with the Universe remain at fixed comoving positions x‖ and x. In terms of the comoving
circumferential radius x, the spatial FRW metric is
dl2 = a2
(
dx2
1− κx2+ x2do2
)
, (10.16)
10.6 Spatial part of the FRW metric: more formal approach 143
where the curvature constant κ, a constant in time and space, is related to the curvature K, equation (10.10),
by
κ ≡ a2K . (10.17)
Alternatively, in terms of the geodesic comoving radius x‖, the spatial FRW metric is
dl2 = a2(
dx2‖ + x2do2
)
, (10.18)
where
x =
sin(κ1/2x‖)
κ1/2κ > 0 closed ,
x‖ κ = 0 flat ,
sinh(|κ|1/2x‖)
|κ|1/2κ < 0 open .
(10.19)
For some purposes it is convenient to normalize the cosmic scale factor a so that κ = 1, 0, or −1. In this
case the spatial FRW metric may be written
dl2 = a2(
dχ2 + x2do2)
, (10.20)
where
x =
sin(χ) κ = 1 closed ,
χ κ = 0 flat ,
sinh(χ) κ = −1 open .
(10.21)
Exercise 10.1 By a suitable transformation of the comoving radial coordinate x, bring the spatial FRW
metric (10.16) to the “isotropic” form
dl2 =a2
(
1 + 14κX
2)2
(
dX2 +X2do2)
. (10.22)
What is the relation between X and x?
10.6 Spatial part of the FRW metric: more formal approach
A more formal approach to the derivation of the spatial FRW metric from the cosmological principle starts
with the proposition that the spatial components Gij of the Einstein tensor at fixed scale factor a (all time
derivatives of a set to zero) should be proportional to the metric tensor
Gij = K gij (i, j = 1, 2, 3) . (10.23)
144 Homogeneous, Isotropic Cosmology
Without loss of generality, the spatial metric can be taken to be of the form
dl2 = f(r) dr2 + r2do2 . (10.24)
Imposing the condition (10.23) on the metric (10.24) recovers the spatial FRW metric (10.11).
10.7 FRW metric
The full Friedmann-Robertson-Walker spacetime metric is
ds2 = − dt2 + a(t)2(
dx2
1− κx2+ x2do2
)
(10.25)
where t is cosmic time, which is the proper time experienced by comoving observers, who remain at rest
in comoving coordinates dx = dθ = dφ = 0. Any of the alternative versions of the comoving spatial FRW
metric, equations (10.16), (10.18), (10.20), or (10.22). may be used as the spatial part of the FRW spacetime
metric (10.25).
10.8 Einstein equations for FRW metric
The Einstein equations for the FRW metric (10.25) are
−Gtt = 3
(
κ
a2+a2
a2
)
= 8πGρ ,
Gxx = Gθ
θ = Gφφ = − κ
a2− a2
a2− 2 a
a= 8πGp , (10.26)
where overdots represent differentiation with respect to cosmic time t, so that for example a ≡ da/dt. Note
the trick of one index up, one down, to remove, modulo signs, the distorting effect of the metric on the
Einstein tensor. The Einstein equations (10.26) rearrange to give Friedmann’s equations
a2
a2=
8πGρ
3− κ
a2,
a
a= −4πG
3(ρ+ 3p) .
(10.27)
Friedmann’s two equations (10.27) are fundamental to cosmology. The first one relates the curvature κ of
the Universe to the expansion rate a/a and the density ρ. The second one relates the acceleration a/a to
the density ρ plus 3 times the pressure p.
10.9 Newtonian “derivation” of Friedmann equations 145
10.9 Newtonian “derivation” of Friedmann equations
10.9.1 Energy equation
Model a piece of the Universe as a ball of radius a and mass M = 43πρa
3. Consider a small mass m attracted
by this ball. Conservation of the kinetic plus potential energy of the small mass m implies
1
2ma2 − GMm
a= −κmc
2
2, (10.28)
where the quantity on the right is some constant whose value is not determined by this Newtonian treatment,
but which GR implies is as given. The energy equation (10.28) rearranges to
a2
a2=
8πGρ
3− κc2
a2, (10.29)
which reproduces the first Friedmann equation.
10.9.2 First law of thermodynamics
For adiabatic expansion, the first law of thermodynamics is
dE + p dV = 0 . (10.30)
With E = ρV and V = 43πa
3, the first law (10.30) becomes
d(ρa3) + p da3 = 0 , (10.31)
or, with the derivative taken with respect to cosmic time t,
ρ+ 3(ρ+ p)a
a= 0 . (10.32)
Differentiating the first Friedmann equation in the form
a2 =8πGρa2
3− κc2 (10.33)
gives
2aa =8πG
3
(
ρa2 + 2ρaa)
, (10.34)
and substituting ρ from the first law (10.32) reduces this to
2aa =8πG
3aa (− ρ− 3p) . (10.35)
Hence
a
a= −4πG
3(ρ+ 3p) , (10.36)
which reproduces the second Friedmann equation.
146 Homogeneous, Isotropic Cosmology
10.9.3 Comment on the Newtonian derivation
The above Newtonian derivation of Friedmann’s equations is only heuristic. A different result could have
been obtained if different assumptions had been made. If for example the Newtonian gravitational force law
ma = −GMm/a2 were taken as correct, then it would follow that a/a = − 43πGρ, which is missing the all
important 3p contribution (without which there would be no inflation or dark energy) to Friedmann’s second
equation.
It is notable that the first law of thermodynamics is built in to the Friedmann equations. This implies that
entropy is conserved in FRW Universes. This remains true even when the mix of particles changes, as happens
for example during the epoch of electron-positron annihilation, or during big bang nucleosynthesis. How
then does entropy increase in the real Universe? Through fluctuations away from the perfect homogeneity
and isotropy assumed by the FRW metric.
10.10 Hubble parameter
The Hubble parameter H(t) is defined by
H ≡ a
a. (10.37)
The Hubble parameter H varies in cosmic time t, but is constant in space at fixed cosmic time t.
The value of the Hubble parameter today is called the Hubble constant H0 (the subscript 0 signifies
the present time). The Hubble constant is measured from Cepheids and Type Ia supernova to be (Riess et
al. 2005, astro-ph/0503159)
H0 = 73± 4(stat)± 5(sys) kms−1 Mpc−1 . (10.38)
The distance d to an object that is receding with the expansion of the universe is proportional to the
cosmic scale factor, d ∝ a, and its recession velocity v is consequently proportional to a. The result is
Hubble’s law relating the recession velocity v and distance d of distant objects
v = H0d . (10.39)
Since it takes light time to travel from a distant object, and the Hubble parameter varies in time, the linear
relation (10.39) breaks down at cosmological distances.
We, in the Milky Way, reside in an overdense region of the Universe that has collapsed out of the general
Hubble expansion of the Universe. The local overdense region of the Universe that has just turned around
from the general expansion and is beginning to collapse for the first time is called the Local Group of
galaxies. The Local Group consists of about 40 or so galaxies, mostly dwarf and irregular galaxies. It
contains two major spiral galaxies, Andromeda (M31) and the Milky Way, and one mid-sized spiral galaxy
Triangulum (M33). The Local Group is about 1 Mpc in radius.
10.11 Critical density 147
Because of the ubiquity of the Hubble constant in cosmological studies, cosmologists often parameterize
it by the quantity h defined by
h ≡ H0
100 km s−1 Mpc−1 . (10.40)
10.11 Critical density
The critical density ρcrit is defined to be the density required for the Universe to be flat, κ = 0. According
to the first of Friedmann equations (10.27), this sets
ρcrit ≡3H2
8πG. (10.41)
The critical density ρcrit, like the Hubble parameter H , evolves with time.
10.12 Omega
Cosmologists designate the ratio of the actual density ρ of the Universe to the critical density ρcrit by the
fateful letter Ω, the final letter of the Greek alphabet,
Ω ≡ ρ
ρcrit. (10.42)
With no subscript, Ω denotes the total mass-energy density in all forms. A subscript x on Ωx denotes
mass-energy density of type x.
The curvature density ρk, which is not really a form of mass-energy but it is sometimes convenient to treat
Table 10.1 Cosmic inventory
Species (2008)
Dark energy (Λ) ΩΛ 0.72 ± 0.02Non-baryonic cold dark matter (CDM) Ωc 0.234 ± 0.02Baryonic matter Ωb 0.046 ± 0.002Neutrinos Ων < 0.014Photons (CMB) Ωγ 5 × 10−5
Total Ω 1.005 ± 0.006
Curvature Ωk −0.005 ± 0.006
148 Homogeneous, Isotropic Cosmology
it as though it were, is defined by
ρk ≡ −3κc2
8πGa2(10.43)
and correspondingly Ωk ≡ ρk/ρcrit. According to the first of Friedmann’s equations (10.27), the curvature
density Ωk satisfies
Ωk = 1− Ω . (10.44)
Table 10.1 gives 2008 measurements of Ω in various species, obtained by combining 5-year WMAP CMB
measurements with a variety of other astronomical evidence, including supernovae, big bang nucleosynthesis,
galaxy clustering, weak lensing, and local measurements of the Hubble constant H0.
10.13 Redshifting
The spatial translation symmetry of the FRW metric implies conservation of generalized momentum. As you
will show in a problem set, a particle that moves along a geodesic in the radial direction, so that dθ = dφ = 0,
has 4-velocity uν satisfying
ux‖= constant . (10.45)
This conservation law implies that the proper momentum p‖ of a radially moving particle decays as
p‖ ≡ madx‖dτ∝ 1
a, (10.46)
which is true for both massive and massless particles.
It follows from equation (10.46) that light observed on Earth from a distant object will be redshifted by
a factor
1 + z =a0
a, (10.47)
where a0 is the present day cosmic scale factor. Cosmologists often refer to the redshift of an epoch, since
the cosmological redshift is an observationally accessible quantity that uniquely determines the cosmic time
of emission.
10.14 Types of mass-energy
The energy-momentum tensor Tµν of an FRW Universe is necessarily homogeneous and isotropic, by as-
sumption of the cosmological principle, taking the form (note yet again the trick of one index up and one
10.15 Evolution of the cosmic scale factor 149
down to remove the distorting effect of the metric)
T µν =
T tt 0 0 0
0 T rr 0 0
0 0 T θθ 0
0 0 0 T φφ
=
−ρ 0 0 0
0 p 0 0
0 0 p 0
0 0 0 p
. (10.48)
Table 10.2 gives equations of state p/ρ for generic species of mass-energy, along with (ρ + 3p)/ρ, which
determines the gravitational attraction per unit energy, and how the mass-energy varies with cosmic scale
factor, ρ ∝ an.
Table 10.2 Properties of universes dominated by various species
Species p/ρ (ρ+ 3p)/ρ ρ ∝
Radiation 1/3 2 a−4
Matter 0 1 a−3
Curvature “−1/3” “0” a−2
Vaccum −1 −2 a0
As commented in §23.16 above, the first law of thermodynamics for adiabatic expansion is built into Fried-
mann’s equations. In fact the law represents covariant conservation of energy-momentum for the system as
a whole
DµTµν = 0 . (10.49)
As long as species do not convert into each other (for example, no annihilation), covariant energy-momentum
conservation holds individually for each species, so the first law applies to each species individually, deter-
mining how its energy density ρ varies with cosmic scale factor a. Figure 10.2 illustrates how the energy
densities ρ of various species evolve as a function of scale factor a.
10.15 Evolution of the cosmic scale factor
Given how the energy density ρ of each species evolves with cosmic scale factor a, the first Friedmann
equation then determines how the cosmic scale factor a(t) itself evolves with cosmic time t. The evolution
equation for a(t) can be cast as an equation for the Hubble parameter H ≡ a/a, which in view of the
definition (10.41) of the critical density can be written
H(t)
H0=
[
ρcrit(t)
ρcrit(t0)
]1/2
. (10.50)
150 Homogeneous, Isotropic Cosmology
Cosmic scale factor a
Mas
s-en
ergy
dens
ity
ρ
ργ ∝ a−4
ρm ∝ a−3
ρk ∝ a−2
ρΛ = constant
Figure 10.2 Behavior of the mass-energy density ρ of various species as a function of cosmic time t.
Given the definition (10.43) of the curvature density as the critical density minus the total density, the
critical density ρcrit is itself the sum of the densities ρ of all species including the curvature density
ρcrit = ρk +∑
species x
ρx . (10.51)
Integrating equation (10.50) gives cosmic time t as a function of cosmic scale factor a
t =
∫
da
aH. (10.52)
For example, in the case that the density is comprised of radiation, matter, and vacuum, the critical density
is
ρcrit = ργ + ρm + ρk + ρΛ , (10.53)
and equation (10.50) is
H(t)
H0=(
Ωγa−4 + Ωma
−3 + Ωka−2 + ΩΛ
)1/2, (10.54)
where Ωx represents its value at the present time. The time t, equation (10.52), is then
t =1
H0
∫
da
a (Ωγa−4 + Ωma−3 + Ωka−2 + ΩΛ)1/2
, (10.55)
10.16 Conformal time 151
which is an elliptical integral of the 3rd kind.
If one single species in particular dominates the mass-energy density, then equation (10.55) integrates
easily to give the results in the following table.
Table 10.3 Evolution of cosmic scale factor in universes dominated by various species
Dominant Species a ∝
Radiation t1/2
Matter t2/3
Curvature tVaccum eHt
10.16 Conformal time
Especially when doing cosmological perturbation theory, it is convenient to use conformal time η defined
by (with units c temporarily restored)
a dη ≡ c dt (10.56)
with respect to which the FRW metric is
ds2 = a(η)2(
− dη2 +dx2
1− κx2+ x2do2
)
. (10.57)
The term conformal refers to a metric that is multiplied by an overall factor, the conformal factor. In the
FRW metric (10.57), the cosmic scale factor a is the conformal factor.
Conformal time η has the property that the speed of light is one in conformal coordinates: light moves
unit comoving distance per unit conformal time. In particular, light moving radially towards an observer at
x‖ = 0, with dθ = dφ = 0, satisfies
dx‖dη
= −1 . (10.58)
10.17 Looking back along the lightcone
Since light moves at unit velocity in conformal coordinates, an object at geodesic distance x‖ that emits
light at conformal time ηemit is observed at conformal time ηobs given by
x‖ = ηobs − ηemit . (10.59)
152 Homogeneous, Isotropic Cosmology
The comoving geodesic distance x‖ to an object is
x‖ =
∫ ηobs
ηemit
dη =
∫ tobs
temit
c dt
a=
∫ aobs
aemit
c da
a2H=
∫ z
0
c dz
H, (10.60)
where the last equation assumes the relation 1 + z = 1/a, valid as long as a is normalized to unity at the
observer (us) at the present time aobs = a0 = 1. In the case that the density is comprised of (curvature and)
radiation, matter, and vacuum, equation (10.60) gives
x‖ =c
H0
∫ 1
1/(1+z)
da
a2 (Ωγa−4 + Ωma−3 + Ωka−2 + ΩΛ)1/2
, (10.61)
which is an elliptical integral of the 1st kind. Given the geodesic comoving distance x‖, the circumferential
comoving distance x then follows as
x =sinh(Ω
1/2k H0x‖/c)
Ω1/2k H0/c
. (10.62)
To second order in redshift z,
x ≈ x‖ ≈c
H0
[
z − z2(
Ωγ + 34Ωm + 1
2Ωk
)
+ ...]
. (10.63)
The geodesic and circumferential distances x‖ and x differ at order z3.
10.18 Horizon
Light can come from no more distant point than the Big Bang. This distant point defines the horizon of
our Universe, which is located at infinite redshift, z =∞. Equation (10.60) gives the geodesic distance from
us at redshift zero to the horizon as
x‖(horizon) =
∫ ∞
0
c dz
H(10.64)
where again the cosmic scale factor has been normalized to unity at the present time, a0 = 1.
Equation (10.64) formally defines the event horizon of the Universe, but the cosmological scale over which
objects can continue to affect each other causally is typically smaller than this (much smaller, post-inflation).
It is thus common to define the cosmological horizon distance at any time as
cosmological horizon distance ≡ c
H(10.65)
which is roughly the scale over which objects can remain in causal contact.
Exercise 10.2 Then versus now.
10.18 Horizon 153
10−40 10−30 10−20 10−10 100 1010 102010−40
10−30
10−20
10−10
100
1010
1020
103010−50 10−40 10−30 10−20 10−10 100 1010
Age of the Universe (seconds)
Age of the Universe (years)
Size
ofth
eU
nive
rse
(met
ers)
Figure 10.3 Cosmic scale factor a and cosmological horizon distance c/H as a function of cosmic time t.
1. Prove that∫ ∞
0
xn−1 dx
ex + 1=(
1− 21−n)
∫ ∞
0
xn−1 dx
ex − 1. (10.66)
[Hint: Use the fact that (ex + 1)(ex − 1) = (e2x − 1).] Hence argue that the ratios of energy, entropy,
and number densities of relativistic fermionic (f) to relativistic bosonic (b) species in thermodynamic
equilibrium at the same temperature are
ρf
ρb=sf
sb=
7
8,
nf
nb=
3
4. (10.67)
[Hint: The proper entropy density of each relativistic species is s = (ρ+ p)/T = (4/3)ρ/T .]
2. Weak interactions were fast enough to keep neutrinos in thermodynamic equilibrium with photons, elec-
trons, and positrons up to just before ee annihilation, but then neutrinos decoupled. Argue that conser-
vation of comoving entropy implies
a3T 3(
gγ +7
8ge
)
= T 3γ gγ , (10.68a)
a3T 3 gν = T 3ν gν , (10.68b)
154 Homogeneous, Isotropic Cosmology
10−40 10−30 10−20 10−10 100 1010 1020
10−5
100
105
1010
1015
1020
1025
1030
1035
10−50 10−40 10−30 10−20 10−10 100 1010
Age of the Universe (years)
Age of the Universe (seconds)
Rad
iati
onT
empe
ratu
reof
the
Uni
vers
e(K
elvi
n)
10−40 10−30 10−20 10−10 100 1010 1020
10−30
10−20
10−10
100
1010
1020
1030
1040
1050
1060
1070
1080
1090
10100
10−50 10−40 10−30 10−20 10−10 100 1010
Age of the Universe (years)
Age of the Universe (seconds)
Mas
s-E
nerg
yD
ensi
tyof
the
Uni
vers
e(k
g/m
3 )
Figure 10.4 (Top) Temperature T , and (bottom) mass-energy density ρ, of the Universe as a function ofcosmic time t.
10.18 Horizon 155
where the left hand sides refer to quantities before ee annihilation, which happened at T ∼ 1 MeV ≈ 1010 K,
and the right hand sides to quantities after ee annihilation (including today). Deduce the ratio of neutrino
to photon temperatures today,
Tν
Tγ. (10.69)
Does the temperature ratio (10.69) depend on the number of neutrino types? What is the neutrino
temperature today in K, if the photon temperature today is 2.725 K?
3. The energy, entropy, and number densities of relativistic particles today are, with units restored (energy
density ρ in units energy/volume; entropy and number density s and n in units 1/volume),
ρ = gρ,0
π2(kT0)4
30c3~3, s = gs,0
2π2(kT0)3
45c3~3, n = gn,0
ζ(3)(kT0)3
π2c3~3, (10.70)
where T0 = 2.725 K is the CMB temperature today, ζ(3) = 1.2020569 is a Riemann zeta function, and gρ,0 ,
gs,0 , and gn,0 denote the energy-, entropy-, and number-weighted effective number of relativistic species
today, normalized to 1 per bosonic degree of freedom. What are the arithmetic values of gρ,0 , gs,0 , and
gn,0 if the relativistic species consist of photons and three species of neutrino? What is the energy density
Ωr of relativistic particles today relative to the critical density? [Hint: Don’t forget to take into account
the fact that the neutrino temperature today differs from the photon temperature.]
4. Evidence for neutrino oscillations from the MINOS experiment (2008, http://www-numi.fnal.gov/
PublicInfo/forscientists.html) indicates that at least one neutrino type has mass mν>∼ 0.05 eV. At
what redshift zν would such a neutrino become non-relativistic? If neutrinos are non-relativistic, what is
the neutrino density Ων relative to the critical density, in terms of the sum of the neutrino masses∑
mν?
Which of the effective number of relativistic species today gρ,0 , gs,0 , and gn,0 is changed if some neutrinos
are non-relativistic today?
5. Use entropy conservation to argue that the ratio of the photon temperature T at redshift z in the early
Universe to the photon temperature T0 today is
T
T0= (1 + z)
(
gs,0
gs
)1/3
. (10.71)
What is gs in terms of the numbers gb and gf of relativistic boson and fermion types, if all species were
at the same temperature T ?
Solution. The ratio of neutrino to photon temperatures post ee annihilation is
Tν
Tγ=
(
gγ
gγ + 78 ge
)1/3
=
(
4
11
)1/3
. (10.72)
The Cosmic Neutrino Background temperature is
Tν =
(
4
11
)1/3
2.725 K = 1.945 K . (10.73)
With 2 bosonic degrees of freedom from photons, and 6 fermionic degrees of freedom from 3 relativistic
156 Homogeneous, Isotropic Cosmology
neutrino types, the effective energy-, entropy-, and number-weighted number of relativistic degrees of freedom
is
gρ,0 = gγ +
(
Tν
Tγ
)47
8gν = 2 +
(
4
11
)4/37
86 = 3.36 , (10.74a)
gs,0 = gγ +
(
Tν
Tγ
)37
8gν = 2 +
4
11
7
86 =
43
11= 3.91 , (10.74b)
gn,0 = gγ +
(
Tν
Tγ
)33
4gν = 2 +
4
11
3
46 =
40
11= 3.64 . (10.74c)
The redshift at which a neutrino of mass mν becomes non-relativistic is
1 + zν =mν
Tν= 300
( mν
0.05 eV
)
. (10.75)
If some neutrinos are non-relativistic, then the neutrino density Ων today is related to the sum∑
mν of
neutrino masses by
Ων =8πG
∑
mνnν
3H20
= 5.3× 10−4
( ∑
mν
0.05 eV
)(
h
0.71
)−2
. (10.76)
gρ,0 is changed if some neutrinos are non-relativistic today, but gs,0 and gn,0 remain unchanged. If just one
of the neutrino types is massive, and the other two are relativistic, then
gρ,0 = 2 +
(
4
11
)4/37
84 = 2.91 . (10.77)
The radiation density Ωr today, including photons and neutrinos, is
Ωr =8πGρr
3c2H20
= 1.236× 10−5gρ,0h−2 = 8.2× 10−5
( gρ,0
3.36
)
(
h
0.71
)−2
. (10.78)
If the temperatures of all species are equal, then the entropy-weighted effective number of relativistic species
is
gs = gb +7
8gf . (10.79)
PART FIVE
TETRAD APPROACH TO GENERAL RELATIVITY
Concept Questions
1. The vierbein has 16 degrees of freedom instead of the 10 degrees of freedom of the metric. What do the
extra 6 degrees of freedom correspond to?
2. Tetrad transformations are defined to be Lorentz transformations. Don’t general coordinate transfor-
mations already include Lorentz transformations as a particular case, so aren’t tetrad transformations
redundant?
3. What does coordinate gauge-invariant mean? What does tetrad gauge-invariant mean?
4. Is the coordinate metric gµν tetrad gauge-invariant?
5. What does a directed derivative ∂m mean physically?
6. Is the directed derivative ∂m coordinate gauge-invariant?
7. Is the tetrad metric γmn coordinate gauge-invariant? Is it tetrad gauge-invariant?
8. What is the tetrad-frame 4-velocity um of a person at rest in an orthonormal tetrad frame?
9. If the tetrad frame is accelerating (not in free-fall), which of the following is true/false?
a. Does the tetrad-frame 4-velocity um of a person continuously at rest in the tetrad frame change with
time? ∂0um = 0? D0u
m = 0?
b. Do the tetrad axes γγm change with time? ∂0γγm = 0? D0γγm = 0?
c. Does the tetrad metric γmn change with time? ∂0γmn = 0? D0γmn = 0?
d. Do the covariant components um of the 4-velocity of a person continuously at rest in the tetrad frame
change with time? ∂0um = 0? D0um = 0?
10. Suppose that p = γγmpm is a 4-vector. Is the proper rate of change of the proper components pm measured
by an observer equal to the directed time derivative ∂0pm or to the covariant time derivative D0p
m? What
about the covariant components pm of the 4-vector? [Hint: The proper contravariant components of the
4-vector measured by an observer are pm ≡ γγm · p where γγm are the contravariant locally inertial rest
axes of the observer. Similarly the proper covariant components are pm ≡ γγm · p.]
11. A person with two eyes separated by proper distance δξn observes an object. The observer observes the
photon 4-vector from the object to be pm. The observer uses the difference δpm in the two 4-vectors
detected by the two eyes to infer the binocular distance to the object. Is the difference δpm in photon
160 Concept Questions
4-vectors detected by the two eyes equal to the directed derivative δξn∂npm or to the covariant derivative
δξnDnpm?
12. Suppose that pm is a tetrad 4-vector. Parallel-transport the 4-vector by an infinitesimal proper distance
δξn. Is the change in pm measured by an ensemble of observers at rest in the tetrad frame equal to the
directed derivative δξn∂npm or to the covariant derivative δξnDnp
m? [Hint: What if “rest” means that the
observer at each point is separately at rest in the tetrad frame at that point? What if “rest” means that
the observers are mutually at rest relative to each other in the rest frame of the tetrad at one particular
point?]
13. What is the physical significance of the fact that directed derivatives fail to commute?
14. Physically, what do the tetrad connection coefficients Γkmn mean?
15. What is the physical significance of the fact that Γkmn is antisymmetric in its first two indices (if the
tetrad metric γmn is constant)?
16. Are the tetrad connections Γkmn coordinate gauge-invariant?
What’s important?
This part of the notes describes the tetrad formalism of GR.
1. Why tetrads? Because physics is clearer in a locally inertial frame than in a coordinate frame.
2. The primitive object in the tetrad formalism is the vierbein emµ, in place of the metric in the coordinate
formalism.
3. Written suitably, for example as equation (11.9), a metric ds2 encodes not only the metric coefficients gµν ,
but a full (inverse) vierbein emµ, through ds2 = γmn e
mµdx
µ enνdx
ν .
4. The tetrad road from vierbein to energy-momentum is similar to the coordinate road from metric to
energy-momentum, albeit a little more complicated.
5. In the tetrad formalism, the directed derivative ∂m is the analog of the coordinate partial derivative ∂/∂xµ
of the coordinate formalism. Directed derivatives ∂m do not commute, whereas coordinate derivatives
∂/∂xµ do commute.
11
The tetrad formalism
11.1 Tetrad
A tetrad (greek foursome) γγm(x) is a set of axes
γγm ≡ γγ0,γγ1,γγ2,γγ3 (11.1)
attached to each point xµ of spacetime. The common case is that of an orthonormal tetrad, where the axes
form a locally inertial frame at each point, so that the scalar products of the axes constitute the Minkowski
metric ηmn
γγm · γγn = ηmn . (11.2)
However, other tetrads prove useful in appropriate circumstances. There are spinor tetrads, null tetrads
(notably the Newman-Penrose double null tetrad), and others (indeed, the basis of coordinate tangent
vectors gµ is itself a tetrad). In general, the tetrad metric is some symmetric matrix γmn
γγm · γγn ≡ γmn . (11.3)
Andrew’s convention:
latin (black) dummy indices label tetrad frames.
greek (brown) dummy indices label coordinate frames.
Why introduce tetrads?
1. The physics is more transparent when expressed in a locally inertial frame (or some other frame adapted
to the physics), as opposed to the coordinate frame, where Salvador Dali rules.
2. If you want to consider spin- 12 particles and quantum physics, you better work with tetrads.
3. For good reason, much of the GR literature works with tetrads, so it’s useful to understand them.
11.2 Vierbein
The vierbein (German four-legs, or colloquially, critter) emµ is defined to be the matrix that transforms
between the tetrad frame and the coordinate frame (note the placement of indices: the tetrad index m comes
11.3 The metric encodes the vierbein 163
first, then the coordinate index µ)
γγm = emµ gµ . (11.4)
The vierbein is a 4× 4 matrix, with 16 independent components. The inverse vierbein emµ is defined to be
the matrix inverse of the vierbein emµ, so that
emµ em
ν = δνµ , em
µ enµ = δn
m . (11.5)
Thus equation (11.4) inverts to
gµ = emµ γγm . (11.6)
11.3 The metric encodes the vierbein
The scalar spacetime distance is
ds2 = gµν dxµ dxν = gµ · gν dx
µ dxν = γmn em
µ en
ν dxµ dxν (11.7)
from which it follows that the coordinate metric gµν is
gµν = γmn em
µ en
ν . (11.8)
The shorthand way in which metric’s are commonly written encodes not only a metric but also an inverse
vierbein, hence a tetrad. For example, the Schwarzschild metric
ds2 = −(
1− 2M
r
)
dt2 +
(
1− 2M
r
)−1
dr2 + r2dθ2 + r2 sin2θ dφ2 (11.9)
takes the form (11.7) with an orthonormal (Minkowski) tetrad metric γmn = ηmn, and an inverse vierbein
encoded in the differentials
e0µdxµ =
(
1− 2M
r
)1/2
dt , (11.10a)
e1µdxµ =
(
1− 2M
r
)−1/2
dr , (11.10b)
e2µdxµ = r dθ , (11.10c)
e3µdxµ = r sin θ dφ , (11.10d)
Explicitly, the inverse vierbein of the Schwarzschild metric is the diagonal matrix
emµ =
(1− 2M/r)1/2 0 0 0
0 (1− 2M/r)−1/2 0 0
0 0 r 0
0 0 0 r sin θ
, (11.11)
164 The tetrad formalism
and the corresponding vierbein is (note that, because the tetrad index is always in the first place and the
coordinate index is always in the second place, the matrices as written are actually inverse transposes of
each other, not just inverses)
emµ =
(1− 2M/r)−1/2 0 0 0
0 (1− 2M/r)1/2 0 0
0 0 1/r 0
0 0 0 1/(r sin θ)
. (11.12)
Concept question 11.1 Schwarzschild vierbein. The components e0t and e1
r of the Schwarzschild
vierbein (11.12) are imaginary inside the horizon. What does this mean? Is the vierbein still valid inside
the horizon? ⋄
11.4 Tetrad transformations
Tetrad transformations are transformations that preserve the fundamental property of interest, for example
the orthonormality, of the tetrad. For all tetrads of interest in these notes, which includes not only orthonor-
mal tetrads, but also spinor tetrads and null tetrads (but not coordinate-based tetrads), tetrad transfor-
mations are Lorentz transformations. Hereafter these notes will presume that a tetrad transformation is a
Lorentz transformation. The Lorentz transformation may be, and usually is, a different transformation at
each point. Tetrad transformations rotate the tetrad axes γγk at each point by a Lorentz transformation
Lkm, while keeping the background coordinates xµ unchanged:
γγk → γγ′k = Lk
m γγm . (11.13)
In the case that the tetrad axes γγk are orthonormal, with a Minkowski metric, the Lorentz transformation
matrices Lkm in equation (11.13) take the familiar special relativistic form, but the linear matrices Lk
m in
equation (11.13) signify a Lorentz transformation in any case.
In all the cases of interest, including orthonormal, spinor, and null tetrads, the tetrad metric γmn is
constant. Lorentz transformations are precisely those transformations that leave the tetrad metric unchanged
γ′kl = γγ′k · γγ′
l = LkmLl
n γγm · γγn = LkmLl
n γmn = γkl . (11.14)
Exercise 11.2 Generators of Lorentz transformations are antisymmetric. From the condition that
the tetrad metric γmn is unchanged by a Lorentz transformation, show that the generator of an infinitesimal
Lorentz transformation is an antisymmetric matrix. Is this true only for an orthonormal tetrad, or is it true
more generally?
Solution. An infinitesimal Lorentz transformation is the sum of the unit matrix and an infinitesimal piece
∆Lkm, the generator of the infinitesimal Lorentz transformation,
Lkm = δm
k + ∆Lkm . (11.15)
11.5 Tetrad Tensor 165
Under such an infinitesimal Lorentz transformation, the tetrad metric transforms to
γ′kl = (δmk + ∆Lk
m)(δnl + ∆Ll
n)γmn ≈ γkl + ∆Lkl + ∆Llk , (11.16)
which by proposition equals the original tetrad metric γkl, equation (11.14). It follows that
∆Lkl + ∆Llk = 0 , (11.17)
that is, the generator ∆Lkl is antisymmetric, as claimed. ⋄
11.5 Tetrad Tensor
In general, a tetrad-frame tensor Akl...mn... is an object that transforms under tetrad (Lorentz) transforma-
tions (11.13) as
A′kl...mn... = Lk
aLlb ... Lm
cLnd ... Aab...
cd... . (11.18)
11.6 Raising and lowering indices
In the coordinate approach to GR, coordinate indices were lowered and raised with the coordinate metric
gµν and its inverse gµν . In the tetrad formalism there are two kinds of indices, tetrad indices and coordinate
indices, and they flip around as follows:
1. Lower and raise coordinate indices with the coordinate metric gµν and its inverse gµν ;
2. Lower and raise tetrad indices with the tetrad metric γmn and its inverse γmn;
3. Switch between coordinate and tetrad frames with the vierbein emµ and its inverse em
µ.
The kinds of objects for which this flippery is valid are called tensors. Tensors with only tetrad indices,
such as the tetrad axes γγm or the tetrad metric γmn, are called tetrad tensors, and they remain unchanged
under coordinate transformations. Tensors with only coordinate indices, such as the coordinate tangent
axes gµ or the coordinate metric gµν , are called coordinate tensors, and they remain unchanged under tetrad
transformations. Tensors may also be mixed, such as the vierbein emµ. And of course just because something
has an index, greek or latin, does not make it a tensor: a tensor is a tensor if any only if it transforms like
a tensor.
11.7 Gauge transformations
Gauge transformations are transformations of the coordinates or tetrad. Such transformations do not
change the underlying spacetime.
Quantities that are unchanged by a coordinate transformation are coordinate gauge-invariant. Quan-
tities that are unchanged under a tetrad transformation are tetrad gauge-invariant. For example, tetrad
tensors are coordinate gauge-invariant, while coordinate tensors are tetrad gauge-invariant.
166 The tetrad formalism
Tetrad transformations have the 6 degrees of freedom of Lorentz transformations, with 3 degrees of freedom
in spatial rotations, and 3 more in Lorentz boosts. General coordinate transformations have 4 degrees of
freedom. Thus there are 10 degrees of freedom in the choice of tetrad and coordinate system. The 16 degrees
of freedom of the vierbein, minus the 10 degrees of freedom from the transformations of the tetrad and
coordinates, leave 6 physical degrees of freedom in spacetime, the same as in the coordinate approach to GR,
which is as it should be.
11.8 Directed derivatives
Directed derivatives ∂m are defined to be the directional derivatives along the axes γγm
∂m ≡ γγm · ∂ = γγm · gµ ∂
∂xµ= em
µ ∂
∂xµis a tetrad 4-vector . (11.19)
The directed derivative ∂m is independent of the choice of coordinates, as signalled by the fact that it has
only a tetrad index, no coordinate index.
Unlike coordinate derivatives ∂/∂xµ, directed derivatives ∂m do not commute. Their commutator is
[∂m, ∂n] =
[
emµ ∂
∂xµ, en
ν ∂
∂xν
]
= emµ ∂en
ν
∂xµ
∂
∂xν− en
ν ∂emµ
∂xν
∂
∂xµ
= (dknm − dk
mn) ∂k is not a tetrad tensor (11.20)
where dlmn ≡ γlk dkmn is the vierbein derivative
dlmn ≡ γlk ek
κ enν ∂em
κ
∂xνis not a tetrad tensor . (11.21)
Since the vierbein and inverse vierbein are inverse to each other, an equivalent definition of dlmn in terms of
the inverse vierbein is
dlmn ≡ − γlk emµ en
ν ∂ek
µ
∂xνis not a tetrad tensor . (11.22)
11.9 Tetrad covariant derivative
The derivation of tetrad covariant derivatives Dm follows precisely the analogous derivation of coordinate
covariant derivatives Dµ. The tetrad-frame formulae look entirely similar to the coordinate-frame formu-
lae, with the replacement of coordinate partial derivatives by directed derivatives, ∂/∂xµ → ∂m, and the
replacement of coordinate-frame connections by tetrad-frame connections Γκµν → Γk
mn. There are two things
to be careful about: first, unlike coordinate partial derivatives, directed derivatives ∂m do not commute;
11.9 Tetrad covariant derivative 167
and second, neither tetrad-frame nor coordinate-frame connections are tensors, and therefore it should be
no surprise that the tetrad-frame connections Γlmn are not related to the coordinate-frame connections
Γλµν by the ‘usual’ vierbein transformations. Rather, the tetrad and coordinate connections are related by
equation (11.32).
If Φ is a scalar, then ∂mΦ is a tetrad 4-vector. The tetrad covariant derivative of a scalar is just the
directed derivative
DmΦ = ∂mΦ is a tetrad 4-vector . (11.23)
If Am is a tetrad 4-vector, then ∂nAm is not a tetrad tensor, and ∂nAm is not a tetrad tensor. But the
4-vector A = γγmAm, being by construction invariant under both tetrad and coordinate transformations, is
a scalar, and its directed derivative is therefore a 4-vector
∂nA = ∂n(γγmAm) is a tetrad 4-vector
= γγm∂nAm + (∂nγγm)Am
= γγm∂nAm + Γk
mnγγk Am (11.24)
where the tetrad-frame connection coefficients, Γkmn, also known as Ricci rotation coefficients (or, in
the context of Newman-Penrose tetrads, spin coefficients) are defined by
∂nγγm ≡ Γkmn γγk is not a tetrad tensor . (11.25)
Equation (11.24) shows that
∂nA = γγk(DnAk) is a tetrad tensor (11.26)
where DnAk is the covariant derivative of the contravariant 4-vector Ak
DnAk ≡ ∂nA
k + ΓkmnA
m is a tetrad tensor . (11.27)
Similarly,
∂nA = γγk(DnAk) (11.28)
where DnAk is the covariant derivative of the covariant 4-vector Ak
DnAk ≡ ∂nAk − ΓmknAm is a tetrad tensor . (11.29)
In general, the covariant derivative of a tensor is
DaAkl...mn... = ∂aA
kl...mn... + Γk
baAbl...mn... + Γl
baAkb...mn... + ...− Γb
maAkl...bn... − Γb
naAkl...mb... − ... (11.30)
with a positive Γ term for each contravariant index, and a negative Γ term for each covariant index.
168 The tetrad formalism
11.10 Relation between tetrad and coordinate connections
The relation between the tetrad connections Γkmn and their coordinate counterparts Γκ
µν follows from
Γkmnγγk = ∂nγγm = en
ν ∂emκgκ
∂xνis not a tetrad tensor
= enν ∂em
κ
∂xνgκ + en
ν emκ ∂gκ
∂xν
= dkmn ekκ gκ + en
ν emκ Γλ
κν gλ . (11.31)
Thus the relation is
Γlmn − dlmn = elλ em
µ enν Γλµν is not a tetrad tensor (11.32)
where
Γlmn ≡ γlk Γkmn . (11.33)
11.11 Torsion tensor
The torsion tensor Smkl , which GR assumes to vanish, is defined in the usual way by the commutator of
the covariant derivative acting on a scalar Φ
[Dk, Dl] Φ = Smkl ∂mΦ is a tetrad tensor . (11.34)
The expression (11.29) for the covariant derivatives coupled with the commutator (11.20) of directed deriva-
tives shows that the torsion tensor is
Smkl = Γm
kl − Γmlk − dm
kl + dmlk is a tetrad tensor (11.35)
where dmkl are the vierbein derivatives defined by equation (11.21). The torsion tensor Sm
kl is antisymmetric
in k ↔ l, as is evident from its definition (11.34).
11.12 No-torsion condition
GR assumes vanishing torsion. Then equation (11.35) implies the no-torsion condition
Γmkl − dmkl = Γmlk − dmlk is not a tetrad tensor . (11.36)
In view of the relation (11.32) between tetrad and coordinate connections, the no-torsion condition (11.36) is
equivalent to the usual symmetry condition Γµκλ = Γµλκ on the coordinate frame connections, as it should
be.
11.13 Antisymmetry of the connection coefficients 169
11.13 Antisymmetry of the connection coefficients
The directed derivative of the tetrad metric is
∂nγlm = ∂n(γγl · γγm)
= γγl · ∂nγγm + γγm · ∂nγγl
= Γlmn + Γmln . (11.37)
In most cases of interest, including orthonormal, spinor, and null tetrads, the tetrad metric is chosen to be a
constant. For example, if the tetrad is orthonormal, then the tetrad metric is the Minkowski metric, which
is constant, the same everywhere. If the tetrad metric is constant, then all derivatives of the tetrad metric
vanish, and then equation (11.37) shows that the tetrad connections are antisymmetric in their first two
indices
Γlmn = −Γmln . (11.38)
This antisymmetry reflects the fact that Γlmn is the generator of a Lorentz transformation for each n.
11.14 Connection coefficients in terms of the vierbein
In the general case of non-constant tetrad metric, and non-vanishing torsion, the following manipulation
∂nγlm + ∂mγln − ∂lγmn = Γlmn + Γmln + Γlnm + Γnlm − Γmnl − Γnml (11.39)
= 2 Γlmn − Slmn − Smnl − Snml − dlmn + dlnm − dmnl + dmln − dnml + dnlm
implies that the tetrad connections Γlmn are given in terms of the derivatives ∂nγlm of the tetrad metric,
the torsion Slmn, and the vierbein derivatives dlmn by
Γlmn = 12 (∂nγlm + ∂mγln − ∂lγmn + Slmn + Smnl + Snml
+ dlmn − dlnm + dmnl − dmln + dnml − dnlm) is not a tetrad tensor . (11.40)
If torsion vanishes, as GR assumes, and if furthermore the tetrad metric is constant, then equation (11.40)
simplifies to the following expression for the tetrad connections in terms of the vierbein derivatives dlmn
defined by (11.21)
Γlmn = 12 (dlmn − dlnm + dmnl − dmln + dnml − dnlm) is not a tetrad tensor . (11.41)
This is the formula that allows connection coefficients to be calculated from the vierbein.
170 The tetrad formalism
11.15 Riemann curvature tensor
The Riemann curvature tensor Rklmn is defined in the usual way by the commutator of the covariant
derivative acting on a contravariant 4-vector. In the presence of torsion,
[Dk, Dl]Am ≡ SnklDnAm +RklmnA
n is a tetrad tensor . (11.42)
If torsion vanishes, as GR assumes, then the definition (11.42) reduces to
[Dk, Dl]Am ≡ RklmnAn is a tetrad tensor . (11.43)
The expression (11.29) for the covariant derivative coupled with the torsion equation (11.34) yields the
following formula for the Riemann tensor in terms of connection coefficients, for the general case of non-
vanishing torsion:
Rklmn = ∂kΓmnl − ∂lΓmnk + ΓamlΓank − Γa
mkΓanl + (Γakl − Γa
lk − Sakl)Γmna is a tetrad tensor . (11.44)
The formula has extra terms (Γakl − Γa
lk − Sakl)Γmna compared to the usual formula for the coordinate-frame
Riemann tensor Rκλµν . If torsion vanishes, as GR assumes, then
Rklmn = ∂kΓmnl − ∂lΓmnk + ΓamlΓank − Γa
mkΓanl + (Γakl − Γa
lk)Γmna is a tetrad tensor . (11.45)
The symmetries of the tetrad-frame Riemann tensor are the same as those of the coordinate-frame Riemann
tensor. For vanishing torsion, these are
R([kl][mn]) , (11.46)
Rklmn +Rknlm +Rkmnl = 0 . (11.47)
Exercise 11.3 Riemann tensor. From the definition (11.42), derive the expression (11.44) for the
Riemann tensor. Show that, in addition to the antisymmetry in kl which follows immediately from the
definition (11.42), the Riemann tensor Rklmn is antisymmetric in the indices mn. [Hint: Start by expanding
out the definition (11.42) using the definition (11.30) of the covariant derivative. You will find it easier to
derive an expression for the Riemann tensor with one index raised, such as Rklmn, but you should resist the
temptation to leave it there, because the symmetries of the Riemann tensor are obscured when one index is
raised. To switch to all lowered indices, you will need to convert terms such as ∂kΓnml by
∂kΓnml = ∂k(γnpΓpml) = γnp ∂kΓpml + Γpml ∂kγ
np . (11.48)
You should show that the directed derivative ∂kγnp in this expression is related to tetrad connections through
a formula similar to equation (11.37)
∂kγnp = −Γnp
k − Γpnk , (11.49)
which you should recognize as equivalent to Dkγnp = 0. To complete the derivation, show that
∂k(Γmnl + Γnml)− ∂l(Γmnk + Γnmk) = [∂k, ∂l]γmn = (Γalk − Γa
kl + Sakl)(Γmna + Γnma) . (11.50)
The antisymmetry of Rklmn in mn follows from equation (11.50).] ⋄
11.16 Ricci, Einstein, Bianchi 171
11.16 Ricci, Einstein, Bianchi
The usual suite of formulae leading to Einstein’s equations apply. Since all the quantities are tensors, and
all the equations are tensor equations, their form follows immediately from their coordinate counterparts.
Ricci tensor:
Rkm ≡ γlnRklmn . (11.51)
Ricci scalar:
R ≡ γkmRkm . (11.52)
Einstein tensor:
Gkm ≡ Rkm − 12 γkm R . (11.53)
Einstein’s equations:
Gkm = 8πGTkm . (11.54)
Bianchi identities:
DkRlmnp +DlRmknp +DmRklnp = 0 , (11.55)
which most importantly imply covariant conservation of the Einstein tensor, hence conservation of energy-
momentum
DkTkm = 0 . (11.56)
11.17 Electromagnetism
11.17.1 Electromagnetic potential and field
The electromagnetic field is derivable from an electromagnetic 4-potential Am,
Am = (φ,A) . (11.57)
The electromagnetic field is a bivector field (an antisymmetric tensor) Fmn,
Fmn ≡ DnAm −DmAn . (11.58)
If torsion vanishes, as general relativity assumes, then the covariant derivatives in the definition (11.58) can
be replaced by directed derivatives,
Fmn = ∂nAm − ∂mAn . (11.59)
If the tetrad is orthonormal, then the traditional description of the electromagnetic field in terms of electric
and magnetic fields E and B is valid:
E = − ∂tA−∇φ , B = ∇×A , (11.60)
172 The tetrad formalism
where ∇ ≡ ∂x, ∂y, ∂z denotes the spatial tetrad directed derivative 3-vector. In an orthonormal tetrad
γγt,γγx,γγy,γγz, the 6 components of the electromagnetic field Fmn are related to the electric and magnetic
fields E = Ex, Ey, Ez and B = Bx, By, Bz by
Fmn =
0 Ex Ey Ez
−Ex 0 −Bz By
−Ey Bz 0 −Bx
−Ez −By Bx 0
, Fmn =
0 −Ex −Ey −Ez
Ex 0 −Bz By
Ey Bz 0 −Bx
Ez −By Bx 0
. (11.61)
11.17.2 Lorentz force law
In the presence of an electromagnetic field Fmn, the general relativistic equation of motion for the 4-velocity
um ≡ dxm/dτ of a particle of mass m and charge q is modified by the addition of a Lorentz force qFnmun
mDum
Dτ= qFn
mun . (11.62)
In the absence of gravitational fields, so D/Dτ = d/dτ , and with um = ut1,v where v is the 3-velocity,
the spatial components of equation (11.62) reduce to [note that d/dt = (1/ut) d/dτ ]
mdutv
dt= q (E + v ×B) (11.63)
which is the classical special relativistic Lorentz force law. The signs in the expression (11.61) for Fmn in
terms of E = Ex, Ey, Ez and B = Bx, By, Bz are arranged to agree with the classical law (11.63).
11.17.3 Maxwell’s equations
The source-free Maxwell’s equations are
DlFmn +DmFnl +DnFlm = 0 , (11.64)
while the sourced Maxwell’s equations are
DmFmn = 4πjn , (11.65)
where jn is the electric 4-current. The sourced Maxwell’s equations (11.65) coupled with the antisymmetry
of the electromagnetic field tensor Fmn ensure conservation of electric charge
Dnjn = 0 . (11.66)
In flat space with a Minkowski metric, the covariant derivatives simplify to ordinary derivatives, Dm →∂m → ∂/∂xm, and the source-free Maxwell’s equations (11.64) reduce to the traditional form
∇ ·B = 0 , ∇×E +∂B
∂t= 0 , (11.67)
11.17 Electromagnetism 173
while the sourced Maxwell’s equations reduce to
∇ ·E = 4πq , ∇×B − ∂E
∂t= 4πj , (11.68)
where the electric charge density q and the electric current density j are the time and space components of
the electric 4-current
jn = q, j . (11.69)
11.17.4 Electromagnetic energy-momentum tensor
The energy-momentum tensor Tmne of an electromagnetic field Fmn is
Tmne =
1
4π
(
FmkF
nk − 1
4γmnFklF
kl
)
. (11.70)
12
∗More on the tetrad formalism
This chapter presents some more advanced aspects of the tetrad formalism. It discusses spinor tetrads
(§12.1) and Newman-Penrose tetrads (§12.2). The chapter goes on to show how the fields that describe
electromagnetic (§12.3) and gravitational (§12.4) waves have a natural and insightful complex structure that
is brought out in a Newman-Penrose tetrad. The Newman-Penrose formalism provides a natural context for
the Petrov classification of the Weyl tensor (§12.5), and for the derivation of the Raychaudhuri equations
(§12.6) which imply the focussing theorem (§12.7) that is a key ingredient of the Penrose-Hawking singularity
theorems.
12.1 Spinor tetrad formalism
In quantum mechanics, fundamental particles have spin. The 3 generations of leptons (electrons, muons,
tauons, and their respective neutrino partners) and quarks (up, strange, top, and their down, charm, and
bottom partners) have spin 12 (in units ~ = 1). The carrier particles of the electromagnetic force (photons),
the weak force (the W± and Z bosons), and the colour force (the 8 gluons), have spin 1. The carrier of the
gravitational force, the graviton, is expected to have spin 2, though as of 2010 no gravitational wave, let
alone its quantum, the graviton, has been detected.
General relativity is a classical, not quantum, theory. Nevertheless the spin properties of classical waves,
such as electromagnetic or gravitational waves, are already apparent classically.
12.1.1 Spinor tetrad
A systematic way to project objects into spin components is to work in a spinor tetrad. As will become
apparent below, equation (12.5), spin describes how an object transforms under rotation about some preferred
axis. In the case of an electromagnetic or gravitational wave, the natural preferred axis is the direction of
propagation of the wave. With respect to the direction of propagation, electromagnetic waves prove to have
two possible spins, or helicities, ±1, while gravitational waves have two possible spins, or helicities, ±2. A
preferred axis might also be set by an experimenter who chooses to measure spin along some particular
12.1 Spinor tetrad formalism 175
direction. The following treatment takes the preferred direction to lie along the z-axis γγz, but there is no
loss of generality in making this choice.
Start with an orthonormal tetrad γγt,γγx,γγy,γγz. If the preferred tetrad axis is the z-axis γγz , then the
spinor tetrad axes γγ+,γγ− are defined to be complex combinations of the transverse axes γγx,γγy,
γγ+ ≡ 1√2(γγx + iγγy) , (12.1a)
γγ− ≡ 1√2(γγx − iγγy) . (12.1b)
The tetrad metric of the spinor tetrad γγt,γγz,γγ+,γγ− is
γmn =
−1 0 0 0
0 1 0 0
0 0 0 1
0 0 1 0
. (12.2)
Notice that the spinor axes γγ+,γγ− are themselves null, γγ+ ·γγ+ = γγ− ·γγ− = 0, whereas their scalar product
with each other is non-zero γγ+ · γγ− = 1. The null character of the spinor axes is what makes spin especially
well-suited to describing fields, such as electromagnetism and gravity, that propagate at the speed of light.
An even better trick in dealing with fields that propagate at the speed of light is to work in a Newman-Penrose
tetrad, §12.2, in which all 4 tetrad axes are taken to be null.
12.1.2 Transformation of spin under rotation about the preferred axis
Under a right-handed rotation by angle χ about the preferred axis γγz, the transverse axes γγx and γγy
transform as SIGN?!
γγx → cosχγγx − sinχγγy ,
γγy → sinχγγx + cosχγγy . (12.3)
It follows that the spinor axes γγ+ and γγ− transform under a right-handed rotation by angle χ about γγz as
γγ± → e±iχ γγ± . (12.4)
The transformation (12.4) identifies the spinor axes γγ+ and γγ− as having spin +1 and −1 respectively.
12.1.3 Spin
More generally, an object can be defined as having spin s if it varies by
esiχ (12.5)
under a right-handed rotation by angle χ about the preferred axis γγz . Thus an object of spin s is unchanged
by a rotation of 2π/s about the preferred axis. A spin-0 object is symmetric about the γγz axis, unchanged
by a rotation of any angle about the axis. The γγz axis itself is spin-0, as is the time axis γγt.
176 ∗More on the tetrad formalism
The components of a tensor in a spinor tetrad inherit spin properties from that of the spinor basis. The
general rule is that the spin s of any tensor component is equal to the number of + covariant indices minus
the number of − covariant indices:
spin s = number of + minus − covariant indices . (12.6)
12.1.4 Spin flip
Under a reflection through the y-axis, the spinor axes swap:
γγ+ ↔ γγ− , (12.7)
which may also be accomplished by complex conjugation. Reflection through the y-axis, or equivalently
complex conjugation, changes the sign of all spinor indices of a tensor component
+↔ − . (12.8)
In short, complex conjugation flips spin, a pretty feature of the spinor formalism.
12.1.5 Spin versus spherical harmonics
In physical problems, such as in cosmological perturbations, or in perturbations of spherical black holes, or
in the hydrogen atom, spin often appears in conjunction with an expansion in spherical harmonics. Spin
should not be confused with spherical harmonics.
Spin and spherical harmonics appear together whenever the problem at hand has a symmetry under the 3D
special orthogonal group SO(3) of spatial rotations (special means of unit determinant; the full orthogonal
group O(3) contains in addition the discrete transformation corresponding to reflection of one of the axes,
which flips the sign of the determinant). Rotations in SO(3) are described by 3 Euler angles θ, φ, χ. Spin
is associated with the Euler angle χ. The usual spherical harmonics Yℓm(θ, φ) are the spin-0 eigenfunctions
of SO(3). The eigenfunctions of the full SO(3) group are the spin harmonics SIGN?
sYℓm(θ, φ, χ) = Θℓms(θ, φ, χ)eimφeisχ . (12.9)
12.1.6 Spinor components of the Einstein tensor
With respect to a spinor tetrad, the components of the Einstein tensor Gmn are
Gmn =
Gtt Gtz Gt+ Gt−
Gtz Gzz Gz+ Gz−
Gt+ Gz+ G++ G+−
Gt− Gz− G+− G−−
. (12.10)
12.2 Newman-Penrose tetrad formalism 177
From this it is apparent that the 10 components of the Einstein tensor decompose into 4 spin-0 components,
4 spin-±1 components, and 2 spin-±2 components:
−2 : G−− ,
−1 : Gt− , Gz− ,
0 : Gtt , Gtz , Gzz , G+− ,
+1 : Gt+ , Gz+ ,
+2 : G++ .
(12.11)
The 4 spin-0 components are all real; in particular G+− is real since G∗+− = G−+ = G+−. The 4 spin-±1 and
2 spin-±2 components comprise 3 complex components
G∗++
= G−− , G∗t+ = Gt− , G∗
z+= Gz− . (12.12)
In some contexts, for example in cosmological perturbation theory, REALLY? the various spin components
are commonly referred to as scalar (spin-0), vector (spin-±1), and tensor (spin-±2).
12.2 Newman-Penrose tetrad formalism
The Newman-Penrose formalism (E. T. Newman & R. Penrose, 1962, “An Approach to Gravitational Ra-
diation by a Method of Spin Coefficients,” J. Math. Phys. 3, 566–579; E. (Ted) Newman & R. Penrose,
2009, “Spin-coefficient formalism,” Scholarpedia, 4(6), 7445, http://www.scholarpedia.org/article/
Newman-Penrose_formalism) provides a particularly powerful way to deal with fields that propagate at the
speed of light. The Newman-Penrose formalism adopts a tetrad in which the two axes γγv (outgoing) and
γγu (ingoing) along the direction of propagation are chosen to be lightlike, while the two axes γγ+ and γγ−
transverse to the direction of propagation are chosen to be spinor axes.
Sadly, the literature on the Newman-Penrose formalism is characterized by an arcane and random notation
whose principal purpose seems to be to perpetuate exclusivity for an old-boys club of people who understand
it. This is unfortunate given the intrinsic power of the formalism. A. Held (1974, “A formalism for the
investigation of algebraically special metrics. I,” Commun. Math. Phys. 37, 311–326) comments that the
Newman-Penrose formalism presents “a formidable notational barrier to the uninitiate.” For example, the
tetrad connections Γkmn are called “spin coefficients,” and assigned individual greek letters that obscure
their transformation properties. Do not be fooled: all the standard tetrad formalism presented in Chapter 11
carries through unaltered. One ill-born child of the notation that persists in widespread use is ψ2−s for the
spin s component of the Weyl tensor, equations (12.49).
Gravitational waves are commonly characterized by the Newman-Penrose (NP) components of the Weyl
tensor. The NP components of the Weyl tensor are sometimes referred to as the NP scalars. The designation
as NP scalars is potentially misleading, because the NP components of the Weyl tensor form a tetrad-frame
tensor, not a set of scalars (though of course the tetrad-frame Weyl tensor is, like any tetrad-frame quantity, a
coordinate scalar). The NP components do become proper quantities, and in that sense scalars, when referred
178 ∗More on the tetrad formalism
to the frame of a particular observer, such as a gravitational wave telescope, observing along a particular
direction. However, the use of the word scalar to describe the components of a tensor is unfortunate.
12.2.1 Newman-Penrose tetrad
A Newman-Penrose tetrad γγv,γγu,γγ+,γγ− is defined in terms of an orthonormal tetrad γγt,γγx,γγy,γγz by
γγv ≡ 1√2(γγt + γγz) , (12.13a)
γγu ≡ 1√2(γγt − γγz) , (12.13b)
γγ+ ≡ 1√2(γγx + iγγy) , (12.13c)
γγ− ≡ 1√2(γγx − iγγy) , (12.13d)
or in matrix form
γγv
γγu
γγ+
γγ−
=1√2
1 0 0 1
1 0 0 −1
0 1 i 0
0 1 −i 0
=
γγt
γγx
γγy
γγz
. (12.14)
Just as each of the spinor axes γγ+ and γγ− individually specifies not one but two distinct directions γγx and
γγy, so also each of the null axes γγv and γγv individually specifies not one but two distinct directions γγt and
γγz. The ingoing null axis γγu may be obtained from the outgoing null axis γγv, and vice versa, by the parity
operation of inverting all the spatial axes.
All four tetrad axes are null
γγv · γγv = γγu · γγu = γγ+ · γγ+ = γγ− · γγ− = 0 . (12.15)
In a profound sense, the null, or lightlike, character of each the four NP axes explains why the NP formalism is
well adapted to treating fields that propagate at the speed of light. The tetrad metric of the Newman-Penrose
tetrad γγv,γγu,γγ+,γγ− is
γmn =
0 −1 0 0
−1 0 0 0
0 0 0 1
0 0 1 0
. (12.16)
12.3 Electromagnetic field tensor 179
12.2.2 Boost weight
A boost along the γγz axis multiplies the outgoing and ingoing axes γγv and γγu by a blueshift factor ǫ and its
reciprocal
γγv → ǫγγv ,
γγu → (1/ǫ)γγu . (12.17)
If the observer boosts by velocity v in the γγz direction away from the source, then the blueshift factor is the
special relativistic Doppler shift factor
ǫ =
(
1− v1 + v
)1/2
. (12.18)
The exponent n of the power ǫn by which an object changes under a boost along the γγz axis is called its
boost weight. Thus γγv has boost weight +1, and γγu has boost weight −1. The spinor axes γγ± both have
boost weight 0. The NP components of a tensor inherit their boost weight properties from those of the NP
basis. The general rule is that the boost weight n of any tensor component is equal to the number of v
covariant indices minus the number of u covariant indices:
boost weight n = number of v minus u covariant indices . (12.19)
12.3 Electromagnetic field tensor
12.3.1 Complexified electromagnetic field tensor
The electromagnetic field Fmn is a bivector, and as such has a natural complex structure. The real part
of the electromagnetic bivector field is the electric field E, which changes sign under spatial inversion,
while the imaginary part is the magnetic field B, which remains unchanged under spatial inversion. In an
orthonormal tetrad γγt,γγx,γγy,γγz, the electromagnetic bivector can be written as though it were a vector
with 6 components:
F =(
E B)
=(
Ex Ey Ez Bx By Bz
)
=(
Ftx Fty Ftz Fzy Fxz Fyx
)
. (12.20)
The bivector has 3 electric and 3 magnetic components:
electric bivector indices: tx, ty, tz ,
magnetic bivector indices: zy, xz, yx .(12.21)
The natural complex structure motivates defining the complexified electromagnetic field tensor Fkl to be
the complex combination
Fkl ≡1
2(Fkl + ∗Fkl) =
1
2
(
δmk δ
nl +
i
2εkl
mn
)
Fmn is a tetrad tensor , (12.22)
180 ∗More on the tetrad formalism
where ∗Fkl denotes the Hodge dual of Fkl
∗Fkl ≡i
2εkl
mn Fmn . (12.23)
Here εklmn is the totally antisymmetric tensor (see Exercise 12.1). The overall factor of 12 on the right hand
sides of equations (12.22) is introduced so that the complexification operator Pmnkl ≡ 1
2 (δmk δ
nl + i
2 εklmn)
is a projection operator, satisfying P 2 = P . The definitions (12.22) and (12.23) of the complexified and
dual electromagnetic field tensors Fkl and ∗Fkl are valid in any frame, not just an orthonormal frame or
a Newman-Penrose frame. In an orthonormal frame, the dual ∗F has the structure, in the same notation
as (12.20),
∗F = i(
B −E)
. (12.24)
In an orthonormal frame the complexified electromagnetic field F , equation (12.22), then has the structure
F = 12
(
1 −i)
(E + iB) . (12.25)
Thus the complexified electromagnetic field effectively embodies the electric and magnetic fields in the
complex 3-vector combination E + iB. The complexified electromagnetic field is self-dual
∗F = F . (12.26)
One advantage of this approach is that Maxwell’s equations (11.64) and (11.65) combine into a single
complex equation
DmFmn = 4πjn (12.27)
whose real and imaginary parts represent respectively the source and source-free Maxwell’s equations.
In some cases, such as the Kerr-Newman geometry expressed in the Boyer-Lindquist orthonormal tetrad,
equation (15.33), the complexified electromagnetic field makes manifest the inner elegance of the geometry.
Exercise 12.1 Totally antisymmetric tensor. In an orthonormal tetrad γγm where γγ0 points to the
future and γγ1, γγ2, γγ3 are right-handed, the contravariant totally antisymmetric tensor εklmn is defined by
(this is the opposite sign from MTW’s notation)
εklmn ≡ [klmn] (12.28)
and hence
εklmn = −[klmn] (12.29)
where [klmn] is the totally antisymmetric symbol
[klmn] ≡
+1 if klmn is an even permutation of 0123 ,
−1 if klmn is an odd permutation of 0123 ,
0 if klmn are not all different .
(12.30)
12.3 Electromagnetic field tensor 181
Argue that in a general basis gµ the contravariant totally antisymmetric tensor εκλµν is
εκλµν = ekκel
λemµen
ν εklmn = e [κλµν] (12.31)
while its covariant counterpart is
εκλµν = −(1/e) [κλµν] (12.32)
where e ≡ |emµ| is the determinant of the vierbein. ⋄
12.3.2 Newman-Penrose components of the electromagnetic field
With respect to a NP null tetrad γγv,γγu,γγ+,γγ−, equation (12.13), the electromagnetic field Fmn has 3
distinct complex components, here denoted φs, of spins respectively s = −1, 0, and +1 in accordance with
the rule (12.6):
−1 : φ−1 ≡ Fu− ,
0 : φ0 ≡ 12 (Fuv + F+−) ,
+1 : φ1 ≡ Fv+ .
(12.33)
The complex conjugates φ∗s of the 3 NP components of the electromagnetic field are
φ∗−1 = Fu+ ,
φ∗0 = 12 (Fuv − F+−) ,
φ∗1 = Fv− ,
(12.34)
whose spins have the opposite sign, in accordance with the rule (12.8) that complex conjugation flips spin.
The above convention that the index s on the NP component φs labels its spin differs from the standard
convention, where the spin s component is capriciously denoted φ1−s (e.g. S. Chandrasekhar, 1983, The
Mathematical Theory of Black Holes, Clarendon Press, Oxford, 1983):
−1 : φ2 ,
0 : φ1 ,
+1 : φ0 .
(standard convention, not followed here) (12.35)
In terms of the electric and magnetic fields E and B in the parent orthonormal tetrad of the NP tetrad,
the 3 complex NP components φs of the electromagnetic field are
φ−1 = 12 [Ex + iBx − i(Ey + iBy)] ,
φ0 = 12 (Ez + iBz) ,
φ1 = 12 [Ex + iBx + i(Ey + iBy)] .
(12.36)
Equations (12.36) show that the NP components of the electromagnetic field contain the electric and mag-
netic fields in the complex combination E + iB, just like the complexified electromagnetic field Fkl, equa-
182 ∗More on the tetrad formalism
tion (12.25). Explicitly, the NP components are related to the components Fkl of the complexified electro-
magnetic field by
φ−1 = Ftx − iFty ,
φ0 = Ftz ,
φ1 = Ftx + iFty .
(12.37)
Part of the power of the NP formalism arises from the fact that it exploits the natural complex structure of
the electromagnetic bivector field.
12.3.3 Newman-Penrose components of the complexified electromagnetic field
The non-vanishing NP components of the complexified electromagnetic field Fkl defined by equation (12.22)
are
Fu− = φ−1 ,
Fuv = F+− = φ0 ,
Fv+ = φ1 ,
(12.38)
whereas components with bivector indices v− or u+ vanish,
Fv− = Fu+ = 0 . (12.39)
The rule that complex conjugation flips spin fails here because the complexification operator in equa-
tion (12.22) breaks the rule. Equations (12.38) and (12.39) show that the complexified electromagnetic
field in an NP tetrad contains just 3 distinct non-vanishing complex components, and those components are
precisely equal to the complex spin components φs.
12.3.4 Propagating components of electromagnetic waves
An oscillating electric charge emits electromagnetic waves. Similarly, an electromagnetic wave incident on an
electric charge causes it to oscillate. An electromagnetic wave moving away from a source is called outgoing,
while a wave moving towards a source is called ingoing.
It can be shown that only the spin −1 NP component φ−1 of an outgoing electromagnetic wave propagates,
carrying electromagnetic energy to infinity:
φ−1 : propagating, outgoing . (12.40)
This propagating, outgoing −1 component has spin −1, but its complex conjugate has spin +1, so effec-
tively both spin components, or helicities, or circular polarizations, of an outgoing electromagnetic wave are
embodied in the single complex component φ−1. The remaining 2 complex NP components φ0 and φ1 of an
outgoing wave are short range, describing the electromagnetic field near the source.
Similarly, only the spin +1 component φ1 of an ingoing electromagnetic wave propagates, carrying energy
from infinity:
φ1 : propagating, ingoing . (12.41)
12.4 Weyl tensor 183
The isolation of each propagating mode into a single complex NP mode, incorporating both helicities, is
simpler than the standard picture of oscillating orthogonal electric and magnetic fields.
12.4 Weyl tensor
The Weyl tensor is the trace-free part of the Riemann tensor,
Cklmn ≡ Rklmn − 12 (γkmRln − γknRlm + γlnRkm − γlmRkn) + 1
6 (γkmγln − γknγlm)R . (12.42)
By construction, the Weyl tensor vanishes when contracted on any pair of indices. Whereas the Ricci and
Einstein tensors vanish identically in any region of spacetime containing no energy-momentum, Tmn = 0,
the Weyl tensor can be non-vanishing. Physically, the Weyl tensor describes tidal forces and gravitational
waves.
12.4.1 Complexified Weyl tensor
The Weyl tensor is is, like the Riemann tensor, a symmetric matrix of bivectors. Just as the electromagnetic
bivector Fkl has a natural complex structure, so also the Weyl tensor Cklmn has a natural complex structure.
The properties of the Weyl tensor emerge most plainly when that complex structure is made manifest.
In an orthonormal tetrad γγt,γγx,γγy,γγz, the Weyl tensor Cklmn can be written as a 6 × 6 symmetric
bivector matrix, organized as a 2× 2 matrix of 3× 3 blocks, with the structure
C =
(
CEE CEB
CBE CBB
)
=
Ctxtx Ctxty Ctxtz Ctxzy Ctxxz Ctxyx
Ctytx ... ... ... ... ...
Ctztx ... ... ... ... ...
Czytx ... ... ... ... ...
Cxztx ... ... ... ... ...
Cyxtx ... ... ... ... ...
, (12.43)
where E denotes electric indices, B magnetic indices, per the designation (12.21). The condition of being
symmetric implies that the 3 × 3 blocks CEE and CBB are symmetric, while CBE = C⊤EB. The cyclic
symmetry (11.47) of the Riemann, hence Weyl, tensor implies that the off-diagonal 3 × 3 block CEB (and
likewise CBE) is traceless.
The natural complex structure motivates defining a complexified Weyl tensor Cklmn by
Cklmn ≡1
4
(
δpkδ
ql +
i
2εkl
pq
)(
δrmδ
sn +
i
2εmn
rs
)
Cpqrs is a tetrad tensor (12.44)
analogously to the definition (12.22) of the complexified electromagnetic field. The definition (12.44) of the
complexified Weyl tensor Cklmn is valid in any frame, not just an orthonormal frame. In an orthonormal
184 ∗More on the tetrad formalism
frame, if the Weyl tensor Cklmn is organized according to the structure (12.43), then the complexified Weyl
tensor Cklmn defined by equation (12.44) has the structure
C =1
4
(
1 −i−i −1
)
(CEE − CBB + i CEB + i CBE) . (12.45)
Thus the independent components of the complexified Weyl tensor Cklmn constitute a 3 × 3 complex sym-
metric traceless matrix CEE − CBB + i(CEB + CBE), with 5 complex degrees of freedom. Although the
complexified Weyl tensor Cklmn is defined, equation (12.44), as a projection of the Weyl tensor, it nevertheless
retains all the 10 degrees of freedom of the original Weyl tensor Cklmn.
The same complexification projection operator applied to the trace (Ricci) parts of the Riemann tensor
yields only the Ricci scalar multiplied by that unique combination of the tetrad metric that has the symme-
tries of the Riemann tensor. Thus complexifying the trace parts of the Riemann tensor produces nothing
useful.
12.4.2 Newman-Penrose components of the Weyl tensor
With respect to a NP null tetrad γγv,γγu,γγ+,γγ−, equation (12.13), the Weyl tensor Cklmn has 5 distinct
complex components, here denoted ψs, of spins respectively s = −2, −1, 0, +1, and +2:
−2 : ψ−2 ≡ Cu−u− ,
−1 : ψ−1 ≡ Cuvu− = C+−u− ,
0 : ψ0 ≡ 12 (Cuvuv + Cuv+−) = 1
2 (C+−+− + Cuv+−) = Cv+−u ,
+1 : ψ1 ≡ Cvuv+ = C−+v+ ,
+2 : ψ2 ≡ Cv+v+ .
(12.46)
The complex conjugates ψ∗s of the 5 NP components of the Weyl tensor are:
ψ∗−2 = Cu+u+ ,
ψ∗−1 = Cuvu+ = C−+u+ ,
ψ∗0 = 1
2 (Cuvuv + Cuv−+) = 12 (C−+−+ + Cuv−+) = Cv−+u ,
ψ∗1 = Cvuv− = C+−v− ,
ψ∗2 = Cv−v− .
(12.47)
whose spins have the opposite sign, in accordance with the rule (12.8) that complex conjugation flips spin.
The above expressions (12.46) and (12.47) account for all the NP components Cklmn of the Weyl tensor but
four, which vanish identically:
Cv+v− = Cu+u− = Cv+u+ = Cv−u− = 0 . (12.48)
The above convention that the index s on the NP component ψs labels its spin differs from the standard con-
vention, where the spin s component of the Weyl tensor is impenetrably denoted ψ2−s (e.g. S. Chandrasekhar
12.4 Weyl tensor 185
1983):
−2 : ψ4 ,
−1 : ψ3 ,
0 : ψ2 ,
+1 : ψ1 ,
+2 : ψ0 .
(standard convention, not followed here) (12.49)
With respect to a triple of bivector indices ordered as u−, uv,+v, the NP components of the Weyl tensor
constitute the 3× 3 complex symmetric matrix IS THIS NP OR COMPLEX NP?
Cklmn =
ψ−2 ψ−1 ψ0
ψ−1 ψ0 ψ1
ψ0 ψ1 ψ2
. (12.50)
12.4.3 Newman-Penrose components of the complexified Weyl tensor
The non-vanishing NP components of the complexified Weyl tensor Cklmn defined by equation (12.44) are
Cu−u− = ψ−2 ,
Cuvu− = C+−u− = ψ−1 ,
Cuvuv = C+−+− = Cuv+− = Cv+−u = ψ0 ,
Cvuv+ = C−+v+ = ψ1 ,
Cv+v+ = ψ2 .
(12.51)
whereas any component with either of its two bivector indices equal to v− or u+ vanishes. As with the
complexified electromagnetic field, the rule that complex conjugation flips spin fails here because the com-
plexification operator breaks the rule. Equations (12.51) show that the complexified Weyl tensor in an NP
tetrad contains just 5 distinct non-vanishing complex components, and those components are precisely equal
to the complex spin components ψs.
12.4.4 Components of the complexified Weyl tensor in an orthonormal tetrad
The complexified Weyl tensor forms a 3 × 3 complex symmetric traceless matrix in any frame, not just an
NP frame. In an orthonormal frame, with respect to a triple of bivector indices tx, ty, tz, the complexified
Weyl tensor Cklmn can be expressed in terms of the NP spin components ψs as
Cklmn =
ψ012 (ψ1 − ψ−1) − i
2 (ψ1 + ψ−1)12 (ψ1 − ψ−1) − 1
2ψ0 + 14 (ψ2 + ψ−2) − i
4 (ψ2 − ψ−2)
− i2 (ψ1 + ψ−1) − i
4 (ψ2 − ψ−2) − 12ψ0 − 1
4 (ψ2 + ψ−2)
. (12.52)
186 ∗More on the tetrad formalism
12.4.5 Propagating components of gravitational waves
For outgoing gravitational waves, only the spin −2 component ψ−2 (the one conventionally called ψ4) prop-
agates, carrying gravitational waves from a source to infinity:
ψ−2 : propagating, outgoing . (12.53)
This propagating, outgoing −2 component has spin −2, but its complex conjugate has spin +2, so effectively
both spin components, or helicities, or circular polarizations, of an outgoing gravitational wave are embodied
in the single complex component. The remaining 4 complex NP components (spins −1 to 2) of an outgoing
gravitational wave are short range, describing the gravitational field near the source.
Similarly, only the spin +2 component ψ2 of an ingoing gravitational wave propagates, carrying energy
from infinity:
ψ2 : propagating, ingoing . (12.54)
12.5 Petrov classification of the Weyl tensor
As seen above, the complexified Weyl tensor is a complex symmetric traceless 3 × 3 matrix. If the matrix
were real symmetric (or complex Hermitian), then standard mathematical theorems would guarantee that
it would be diagonalizable, with a complete set of eigenvalues and eigenvectors. But the Weyl matrix is
complex symmetric, and there is no such theorem.
The mathematical theorems state that a matrix is diagonalizable if and only if it has a complete set of
linearly independent eigenvectors. Since there is always at least one distinct linearly independent eigenvector
associated with each distinct eigenvalue, if all eigenvalues are distinct, then necessarily there is a complete
set of eigenvectors, and the Weyl tensor is diagonalizable. However, if some of the eigenvalues coincide, then
there may not be a complete set of linearly independent eigenvectors, in which case the Weyl tensor is not
diagonalizable.
The Petrov classification, tabulated in Table 12.1, classifies the Weyl tensor in accordance with the number
of distinct eigenvalues and eigenvectors. The normal form is with respect to an orthonormal frame aligned
with the eigenvectors to the extent possible. The tetrad with respect to which the complexified Weyl tensor
takes its normal form is called the Weyl principal tetrad. The Weyl principal tetrad is unique except in
cases D, O, and N. For Types D and N, the Weyl principal tetrad is unique up to Lorentz transformations
that leave the eigen-bivector γγtz unchanged, which is to say, transformations generated by the Lorentz rotor
exp(ζγγtz) where ζ is complex.
The Kerr-Newman geometry is Type D. General spherically symmetric geometries are Type D. The
Friedmann-Robertson-Walker geometry is Type O. Plane gravitational waves are Type N.
12.6 Raychaudhuri equations and the Sachs optical scalars 187
expansion θ rotation ω shear σ
Figure 12.1 Illustrating how the Sachs optical scalars, the expansion θ, the rotation ω, and the shear σ,defined by equations (12.60), characterize the rate at which a bundle of light rays changes shape as itpropagates. The bundle of light is coming vertically upward out of the paper.
12.6 Raychaudhuri equations and the Sachs optical scalars
Consider a light ray. Let the γγv null axis lie along the worldline of the light ray. Choose the Newman-
Penrose tetrad so that the tetrad axes γγm are parallel-transported along the path of the light ray. You can
think of a bunch of observers arrayed along the ray each observing the same image, unprecessed, unboosted,
Table 12.1 Petrov classification of the Weyl tensor
Petrov Distinct Distinct Normal formtype eigenvalues eigenvectors of the complexified Weyl tensor
I 3 3
0
@
ψ0 0 00 − 1
2ψ0 + 1
2ψ2 0
0 0 − 12ψ0 − 1
2ψ2
1
A
D 2 3
0
@
ψ0 0 00 − 1
2ψ0 0
0 0 − 12ψ0
1
A
II 2 2
0
@
ψ0 0 00 − 1
2ψ0 + 1
4ψ2 − i
4ψ2
0 − i4ψ2 − 1
2ψ0 − 1
4ψ2
1
A
O 1 3
0
@
0 0 00 0 00 0 0
1
A
N 1 2
0
@
0 0 00 1
4ψ2 − i
4ψ2
0 − i4ψ2 − 1
4ψ2
1
A
III 1 1
0
@
0 12ψ1 − i
2ψ1
12ψ1 0 0
− i2ψ1 0 0
1
A
188 ∗More on the tetrad formalism
unredshifted. Mathematically, this means that the tetrad axes along the worldline of the light ray satisfy
∂vγγm = 0 . (12.55)
By definition of the tetrad connections, this is equivalent to the conditions that the Newman-Penrose tetrad
connections with final index v all vanish
Γkmv = 0 . (12.56)
The conditions (12.56) constitute a set of 6 conditions which define the Lorentz transformation of the tetrad
axes along the worldline of the light ray. Given the conditions (12.56), the usual expression (11.45) for the
Riemann tensor implies that the rate of change ∂vΓmnl of each of the 18 remaining tetrad connections Γmnl,
those with final index l 6= v, along the worldline of the light ray satisfies
∂vΓmnl + ΓkvlΓmnk = Rvlmn . (12.57)
As commented after the definition (12.13) of the NP tetrad, the null axis γγv ≡ (γγt + γγz)/√
2 defines not a
single direction, but rather a 2D surface spanned by the two directions γγt and γγz. Orthogonal to this surface
is the 2D surface spanned by the transverse axes γγx and γγy, or equivalently by the spinor axes γγ+ and γγ−.
Of the 18 tetrad connections Γmnl with l 6= v, four are embodied in the extrinsic curvature Kab defined by
Kab ≡ γγa · ∂bγγv = Γavb for a, b = +,− . (12.58)
The extrinsic curvature (12.58) describes how the null axis γγv varies over the 2D surface spanned by the
transverse axes γγ+ and γγ−. For the extrinsic curvature Kab, the evolution equations (12.57) become
∂vKab +K+
bKa+ +K−bKa− = Rvbav . (12.59)
The Sachs optical scalars constitute the components of the extrinsic curvatureKab defined by equation (12.58).
Conventionally, the Sachs optical scalars consist of the expansion θ, the rotation ω, and the complex shear
σ, defined in terms of the extrinsic curvature (12.58) by
θ + iω ≡ K+− , (12.60a)
σ ≡ K++ , (12.60b)
whose complex conjugates are θ − iω ≡ K−+ and σ∗ = K−−. Resolved into real and imaginary parts, the
definitions (12.60) of the Sachs optical scalars are
θ ≡ 12 (K+− +K−+) , (12.61a)
ω ≡ 12i (K+− −K−+) , (12.61b)
Reσ ≡ 12 (K++ +K−−) , (12.61c)
Imσ ≡ 12i (K++ −K−−) . (12.61d)
Physically, the Sachs scalars characterize how the shape of a bundle of light rays evolves as it propagates,
as illustrated in Figure 12.1. The expansion represents how fast the bundle expands, the rotation how fast
12.7 Focussing theorem 189
it rotates, and the shear how fast its ellipticity is changing. The amplitude and phase of the complex shear
represent the amplitude and phase of the major axis of the shear ellipse.
The derivation from equations (12.59) of the evolutionary equations governing the Sachs scalars θ, ω, and
σ is left as Exercise 12.2. The equation (12.62b) for the shear σ shows that the shear changes only in the
presence of a non-vanishing Weyl tensor, that is, in the presence of tidal forces. The equation (12.63b) for
the rotation ω shows that if the rotation is initially zero, then it will remain zero along the path of the light
bundle. Thus geodesic motion cannot by itself generate rotation from nothing; but non-geodesic processes,
such the electromagnetic scattering of light, can generate rotation. The equation (12.63a) for the expansion
θ is commonly called the Raychaudhuri equation (A. K. Raychaudhuri, 1955, “Relativistic cosmology,
I,” Phys. Rev. 98, 1123). The Raychaudhuri equation is the basis of the focussing theorem, §12.7, which is
a central ingredient of the Penrose-Hawking singularity theorems.
Exercise 12.2 Raychaudhuri equations
Show that equations (12.59) imply the following evolutionary equations for the Sachs scalars defined by (12.60):
(∂v + θ + iω)(θ + iω) + σσ∗ + 12Gvv = 0 , (12.62a)
(∂v + 2θ)σ + Cv+v+ = 0 , (12.62b)
where Gmn and Cklmn denote the Einstein and Weyl tensors as usual. Show that the first (12.62a) of these
equations is equivalent to the two equations
∂vθ + (θ2 − ω2) + σσ∗ + 12Gvv = 0 , (12.63a)
(∂v + 2θ)ω = 0 . (12.63b)
⋄
12.7 Focussing theorem
The Raychaudhuri equation (12.63a) provides the basis for the focussing theorem, which is a key ingredient
of the singularity theorems introduced by Penrose (R. Penrose, 1965, “Gravitational collapse and spacetime
singularities,” Phys. Rev. Lett. 14, 57–59) and elaborated extensively by Hawking (S. W. Hawking and G.
F. R. Ellis, 1975, The large scale structure of space-time, Cambridge University Press).
If the rotation ω vanishes, then equation (12.63a) for the expansion θ simplifies to
∂vθ + θ2 + σσ∗ + 12Gvv = 0 . (12.64)
The terms θ2 and σσ∗ are necessarily positive. The NP component Gvv of the Einstein tensor is related to
the components in the parent orthonormal tetrad by
Gvv = 12Gtt +Gtz + 1
2Gzz . (12.65)
190 ∗More on the tetrad formalism
Boosted along the z-direction into the center-of-mass frame, where Gtz = 0, equation (12.65) reduces to
Gvv = 4π(ρ+ pz) (12.66)
where ρ is the energy density and pz the pressure along the z-direction. If it is true that
ρ+ pz ≥ 0 , (12.67)
then the rotation-free Raychaudhuri equation (12.64) shows that the expansion θ must always decrease.
13
∗The 3+1 (ADM) formalism
Einstein’s equations constitute a set of 10 coupled second-order partial differential equations. Solving these
equations in a general fashion presents a formidable challenge. The 3+1, or ADM, formalism devised by
R. Arnowitt, S. Deser, & C. W. Misner (1959, Phys. Rev. 116, 1322–1330; 1962, “The dynamics of general
relativity,” in Gravitation: an introduction to current research, ed. L. Witten, 227–265) offers an insightful
and systematic way to proceed. The formalism is widely used in numerical general relativity. For reviews,
see L. Lehner (2001, “Numerical relativity: a review,” CQG 18, R25–86, gr-qc/0106072), and H. Shinkai
(2008, “Formulations of the Einstein equations for numerical simulations,” APCTP winter school on black
hole astrophysics, arXiv:0805.0068).
The central idea of the ADM formalism is to recast the Einstein equations into Hamiltonian form. The
Hamiltonian approach identifies “canonical momenta” conjugate to the “coordinates,” and converts the
equations of motion from second order partial differential equations in the coordinates into coupled first
order partial differential equations in the coordinates and momenta. The Hamiltonian H of a system is
its “energy” expressed in terms of the coordinates and momenta. In quantum mechanics, equating the
time translation operator to the Hamiltonian operator, i~ ∂/∂t = H , determines the evolution in time t of
the system. To implement the Hamiltonian approach in general relativity, it is necessary to identify one
coordinate, the time coordinate t, as having a special status. The system of Einstein (and other) equations
is evolved by integrating from one spacelike hypersurface of constant time, t = constant, to the next.
The 3+1 formalism provides answers to several basic questions about the dynamical structure of Einstein’s
equations:
1. Are there natural coordinates for the gravitational field, and what are they? Answer: Yes. The natural
coordinates are the 6 components of the spatial metric gαβ on the hypersurfaces of constant time. The
3+1 formalism shows that only these 6 of the 10 components of the 4D metric gµν are governed by time
evolution equations. The remaining 4 degrees of freedom in the metric represent gauge freedoms associated
with general coordinate transformations.
2. What happens to the 4 degrees of freedom associated with general coordinate transformations? Answer:
They are accomodated into the lapse α and shift βα, which express the rate at which the unit timelike
normal γγ0 to the spatial hypersurfaces of constant time marches through the coordinates, equations (13.9).
192 ∗The 3+1 (ADM) formalism
In the 3+1 formalism, the lapse and shift are specifiable arbitrarily, and are not governed by time evolution
equations.
3. What are the 6 momenta conjugate to the 6 coordinates gαβ? Answer: The momenta are, up to a factor,
the components of the trace-modified extrinsic curvature Kαβ − gαβK, equation (13.36). The extrinsic
curvature Kαβ , whose tetrad-frame expression is defined by equation (13.12b), is a symmetric 3×3 matrix
that describes how the unit timelike normal γγ0 varies over the spatial hypersurfaces of constant time.
4. What is the structure of the 10 Einstein equations in the 3+1 formalism? Answer: The 6 spatial com-
ponents of the Einstein equations provide dynamical time evolution equations for the 6 momenta. The
4 remaining components of the Einstein equations, the time-time and time-space components, prove to
be constraint equations, called the Hamiltonian (or scalar) and momentum (or vector) constraints. The
constraint equations specify conditions that must be arranged to be satisfied on the initial hypersurface
of constant time, but thereafter the constraints are automatically satisfied (modulo numerical error and
instabilities).
5. How is covariant conservation of energy-momentum expressed? Answer: The fact that the Hamiltonian
and momentum constraints continue to be satisfied as time advances expresses covariant conservation of
energy-momentum, as guaranteed by the contracted Bianchi identities.
In this chapter, the coordinate time index is t, while the tetrad time index is 0. Early-alphabet greek
(brown) letters α, β, ..., denote 3D spatial coordinate indices, while mid-alphabet greek letters κ, λ, ...,
denote 4D spacetime coordinate indices. Early-alphabet latin (black) letters a, b, ..., denote 3D spatial
tetrad indices, while mid-alphabet latin (black) letters k, l, ..., denote 4D spacetime tetrad indices.
13.1 ADM tetrad
The ADM formalism splits the spacetime coordinates xµ into a time coordinate t and spatial coordinates
xα, α = 1, 2, 3,
xµ ≡ t, xα . (13.1)
At each point of spacetime, the spacelike hypersurface of constant time t has a unique future-pointing unit
normal γγ0, defined to have unit length and to be orthogonal to the spatial tangent axes gα,
γγ0 · γγ0 = −1 , γγ0 · gα = 0 α = 1, 2, 3 . (13.2)
The central element of the ADM approach is to work in a tetrad frame γγm consisting of this time axis γγ0,
together with three spatial tetrad axes γγa that are orthogonal to the tetrad time axis γγ0, and therefore lie
in the 3D spatial hypersurface of constant time,
γγ0 · γγa = 0 a = 1, 2, 3 . (13.3)
The tetrad metric γmn in the ADM formalism is thus
γmn =
( −1 0
0 γab
)
, (13.4)
13.2 Traditional ADM approach 193
and the inverse tetrad metric γmn is correspondingly
γmn =
( −1 0
0 γab
)
, (13.5)
whose spatial part γab is the inverse of γab. Given the conditions (13.2) and (13.3), the vierbein emµ and
inverse vierbein emµ take the form
emµ =
(
1/α βα/α
0 eaα
)
, emµ =
(
α 0
−eaαβ
α eaα
)
, (13.6)
where α and βα are the lapse and shift (see next paragraph), and eaα and ea
α represent the spatial vierbein
and inverse vierbein, which are inverse to each other, eaαeb
α = δca. The ADM metric is
ds2 = −α2dt2 + gαβ (dxα − βαdt)(
dxβ − ββdt)
, (13.7)
where gαβ is the spatial coordinate metric
gαβ = γabea
αebβ . (13.8)
Essentially all the tetrad formalism developed in Chapter 11 carries through, subject only to the condi-
tions (13.2) and (13.3).
The vierbein coefficient α is called the lapse, while βα is called the shift. Physically, the lapse α is the
rate at which the proper time τ of the tetrad rest frame elapses per unit coordinate time t, while the shift βα
is the velocity at which the tetrad rest frame moves through the spatial coordinates xα per unit coordinate
time t,
α =dτ
dt, βα =
dxα
dt. (13.9)
These relations (13.9) follow from the fact that the 4-velocity in the tetrad rest frame is by definition
um ≡ 1, 0, 0, 0, so the coordinate 4-velocity uµ ≡ emµum of the tetrad rest frame is
dxµ
dτ≡ uµ = et
µ =1
α1, βα . (13.10)
13.2 Traditional ADM approach
The traditional ADM approach sets the spatial tetrad axes γγa equal to the spatial coordinate tangent axes
gα,
γγa = gα (traditional ADM) , (13.11)
equivalent to choosing the spatial vierbein to be the unit matrix, eaα = δα
a . The traditional ADM approach
may be termed semi-tetrad, since it works with a tetrad time axis γγ0 together with coordinate spatial axes
gα. It is natural however to extend the ADM approach into a full tetrad approach, allowing the spatial
tetrad axes γγa to be chosen more generally, subject only to the condition (13.3) that they be orthogonal to
194 ∗The 3+1 (ADM) formalism
the tetrad time axis, and therefore lie in the hypersurface of constant time t. For example, the spatial tetrad
γγa can be chosen to form 3D orthonormal axes, γab ≡ γγa · γγb = δab, so that the full 4D tetrad metric γmn is
Minkowski.
This chapter follows the full tetrad approach to the ADM formalism, but all the results hold for traditional
case where the spatial tetrad axes are set equal to the coordinate spatial axes, equation (13.11).
13.3 Spatial tetrad vectors and tensors
Since the tetrad time axis γγ0 in the ADM formalism is defined uniquely by the choice of hypersurfaces
of constant time t, there is no freedom of tetrad transformations of the time axis distinct from temporal
coordinate transformations (no distinct freedom of Lorentz boosts). However, there is still freedom of tetrad
transformations of the spatial tetrad axes (spatial rotations).
A covariant spatial tetrad vector Aa is defined in the usual way as a vector that transforms like the
spatial tetrad axes γγa. Similarly a covariant spatial tetrad tensor is a tensor that transforms like products
of the spatial tetrad axes γγa. Indices on spatial tetrad vectors and tensors are raised with the inverse spatial
tetrad metric γγab, and lowered with the spatial tetrad metric γγab.
A temporal coordinate transformation changes the hypersurface of constant time t, and therefore changes
its unit normal, the tetrad time axis γγ0, and correspondingly all the spatial tetrad axes γγa.
13.4 ADM connections, gravity, and extrinsic curvature
Since the tetrad time axis γγ0 is a spatial tetrad scalar, its directed time derivative ∂0γγ0 is a spatial tetrad
scalar, while its directed spatial derivatives ∂bγγ0 form a spatial tetrad vector. It follows that the connections
Γm0n defined by the directed derivatives of γγ0 are spatial tetrad tensors. These connections play an important
role in the ADM formalism, and they are given special names and symbols, the gravity κa, and the extrinsic
curvature Kab (the remaining components of the directed derivatives of γγ0 vanish, Γ000 = Γ00a = 0):
κa ≡ γγa · ∂0γγ0 = Γa00 is a spatial tetrad vector , (13.12a)
Kab ≡ γγa · ∂bγγ0 = Γa0b is a spatial tetrad tensor . (13.12b)
The gravity κa is justly named because the geodesic equation shows that it is minus the acceleration expe-
rienced in the tetrad rest frame, where um = 1, 0, 0, 0,dua
dτ= −κa . (13.13)
The extrinsic curvature Kab describes how the unit normal γγ0 to the 3-dimensional spatial hypersurface
changes over the hypersurface, and can therefore be regarded as embodying the curvature of the 3-dimensional
spatial hypersurface embedded in the 4-dimensional spacetime. From equation (11.40) with vanishing torsion,
13.5 ADM Riemann, Ricci, and Einstein tensors 195
it follows that the gravity and the extrinsic curvature are
κa = d00a , (13.14a)
Kab = 12 (∂0γab − dab0 − dba0 + da0b + db0a) , (13.14b)
where the relevant vierbein derivatives dlmn are
d00a =1
α∂aα , da0b =
1
αeaα ∂bβ
α , dab0 = − γac ebβ ∂0e
cβ . (13.15)
Equation (13.14b) shows that the extrinsic curvature is symmetric,
Kab = Kba . (13.16)
The non-vanishing tetrad connections are, from the general formula (11.40) with vanishing torsion,
Γa00 = −Γ0a0 = κa , (13.17a)
Γa0b = −Γ0ab = Kab , (13.17b)
Γab0 = Kab + dab0 − da0b , (13.17c)
Γabc = same as eq. (11.40) . (13.17d)
The connections (13.17a) and (13.17b) form, as commented above, a spatial tetrad vector κa and tensor
Kab, but the remaining connections (13.17c) and (13.17d) are not spatial tetrad tensors. Note that the
purely spatial tetrad connections Γabc, like the spatial tetrad axes γγa, transform under temporal coordinate
transformations despite the absence of temporal indices.
13.5 ADM Riemann, Ricci, and Einstein tensors
The ADM Riemann tensor Rklmn inherits from the standard tetrad formalism the property of being a full
tetrad tensor, its components transforming like products of the tetrad axes γγm. Of course, since the ADM
tetrad time axis γγ0 is tied to the coordinates, a tetrad transformation of the time axis requires a simultaneous
coordinate transformation consistent with it. The usual expression (11.45) for the tetrad-frame Riemann
tensor with vanishing torsion yields the ADM Riemann tensor
R0a0b = −D0Kab −KacKcb +
1
αDaDbα , (13.18a)
R0abc = DcKab −DbKac , (13.18b)
Rabcd = KacKbd −KadKbc +R(3)abcd . (13.18c)
Here R(3)abcd is the spatial tetrad Riemann tensor considered confined to the 3D spatial hypersurface (given by
equation (11.45) with all time components discarded), and Dm denotes the usual 4D tetrad-frame covariant
derivative, with the understanding that when acting on a spatial tensor such as Kab, all time components
196 ∗The 3+1 (ADM) formalism
of the spatial tensor are to be considered equal to zero. Thus the 4D tetrad covariant derivative of a spatial
tetrad vector Aa is
DmAa = ∂mAa − ΓbamAb , (13.19)
in which the possible Γ0amA0 term is considered to vanish. When acting on a spatial tetrad tensor, the spatial
part Da of the 4D tetrad covariant derivative Dm involves only spatial connections, and is identical to the
3D spatial tetrad covariant derivative D(3) considered confined to the 3D spatial hypersurface,
Da ≡ D(3)a when acting on a spatial tetrad tensor . (13.20)
The covariant tetrad time derivative D0 acting on a spatial tetrad tensor yields a 3D spatial tetrad tensor,
but the full 4D tetrad covariant derivative Dm acting on a spatial tetrad tensor does not yield a 4D tetrad
tensor. It is important to bear these fine distinctions in mind when integrating by parts, as done in §13.6.
The Ricci tensor Rkm ≡ γlnRklmn is, like the Riemann tensor, a tetrad tensor. Its components are
R00 = − ∂0K −KabKab +
1
αDaD
aα , (13.21a)
R0a = DbKab − ∂aK , (13.21b)
Rab = D0Kab +KKab −1
αDaDbα+R
(3)ab , (13.21c)
where K ≡ γabKab is the trace of the extrinsic curvature, and R(3)ac ≡ γbdR
(3)abcd is the Ricci tensor confined
to the 3D spatial hypersurface. The Ricci scalar R ≡ γkmRkm is
R = 2 ∂0K +K2 +KabKab − 2
αDaD
aα+R(3) , (13.22)
where R(3) ≡ γabR(3)ab is the Ricci scalar confined to the 3D spatial hypersurface.
The Einstein tensor Gkm ≡ Rkm − 12γkmR is
G00 = 12
(
K2 −KabKab +R(3)
)
, (13.23a)
G0a = DbKab −DaK , (13.23b)
Gab = D0Kab +KKab −1
αDaDbα− γab
(
∂0K + 12K
2 + 12KcdK
cd − 1
αDcD
cα
)
+G(3)ab , (13.23c)
where G(3)ab ≡ R
(3)ab − 1
2γabR(3) is the Einstein tensor confined to the 3D spatial hypersurface.
13.6 ADM action
In this chapter up to this point, the Einstein tensor and other quantities have been expressed in terms of
other things, but no equation of motion has been invoked. The ADM philosophy is to derive the gravitational
equations of motion — Einstein’s equations — in Hamiltonian form. The starting point is to write down the
gravitational action, extremization of which will yield equations of motion.
13.6 ADM action 197
The Hilbert gravitational action Sg is
Sg =
∫ tf
ti
Lg dx4 =1
16π
∫ tf
ti
Rdx4 , (13.24)
with scalar Lagrangian Lg equal to a normalization factor times the Ricci scalar R,
Lg =1
16πR . (13.25)
The integration in the action (13.24) is over a 4-volume from an initial hypersurface of constant time ti to
a final hypersurface of constant time tf . The integration measure dx4 in the integral (13.24) denotes the
scalar 4-volume element1. With respect to coordinates t, xα, the scalar 4-volume element is
dx4 = αdtdx30 = (α/e) dt d3x , (13.26)
where dx30 is the tetrad time component (the component along the timelike normal γγ0 to the spatial hy-
persurface) of the vector 3-volume element dx3k, and e = |ea
β | is the determinant of the spatial vierbein, so
that 1/e = |eaβ | is the determinant of the inverse spatial vierbein. Instead of the scalar Lagrangian Lg, it is
equally possible to work with the Lagrangian density Lg,
Sg =
∫ tf
ti
Lg dt d3x with Lg ≡ Lg α/e . (13.27)
Either way, the important thing is to be careful to get factors in the integration measure right when inte-
grating by parts. The following development works with the scalar Lagrangian Lg, in which case a term
integrates over a scalar 4-volume V to a 3D surface integral over the 3-boundary ∂V of the volume provided
that the term is a covariant 4-divergence,∫
V
D ·A dx4 =
∮
∂V
A · dx3 . (13.28)
The least action principle demands that the conditions on the initial and final hypersurfaces ti and tf be
considered fixed, and asserts that the path followed by the system between the fixed initial and final conditions
is such that the action is minimized. The equations of motion that result from extremizing the action are
unaffected by terms in the scalar Lagrangian that are covariant 4-divergences, since these integrate to surface
terms that are asserted to be fixed, and therefore unchanged by variation. The ADM expression (13.22) for
the Ricci scalar involves two terms, 2∂0K and (2/α)DaDaα, that contain second derivatives of the vierbein
coefficients, and therefore demand to be integrated by parts to bring the Lagrangian into Hamiltonian form,
depending only on first derivatives. To integrate the ∂0K term by parts, it is necessary to express it in terms
of a covariant 4-divergence, which is accomplished by
∂0K = DmKm −K2 , (13.29)
1 Technically, the scalar volume element dx4 is the quadvector, or pseudoscalar, 4-volume element, the differential 4-form14!
εklmn dxk∧ dxl
∧ dxm∧ dxn. Likewise the vector 3-volume element dx3
k is the trivector, or pseudovector, 3-volume
element, the differential 3-form 13!
εklmn dxl∧ dxm
∧ dxn.
198 ∗The 3+1 (ADM) formalism
where Km ≡ K, 0, 0, 0. Similarly, the (1/α)DaDaα term is converted to a 4-divergence by
1
αDaD
aα = Dmκm , (13.30)
where κm ≡ 0, κa with κa defined by equation (13.12a).
Inserting the ADM expression (13.22) for the Ricci scalar into the Hilbert action (13.24), and integrating
the ∂0K and (1/α)DaDaα terms by parts, yields the ADM gravitational action
Sg =1
8π
[∫
K dx30
]tf
ti
+1
8π
∫
∂V
κa dx3a +
1
16π
∫ tf
ti
(
KabKab −K2 +R(3)
)
dx4 . (13.31)
For the purpose of extremizing the action, the surface terms can be discarded. Thus the action to be
extremized is the one with the ADM Lagrangian
LADM =1
16π
(
KabKab −K2 +R(3)
)
. (13.32)
According to the usual procedure, conjugate momenta are obtained as partial derivatives of the Lagrangian
with respect to velocities. In the present instance, the velocities are time derivatives of the coordinates. Now
the only things containing time derivatives in the ADM Lagrangian (13.32) are those contained in the terms
involving the extrinsic curvature Kab (the spatial Ricci scalar R(3) contains no time derivatives). From the
expression (13.14b) for the extrinsic curvature, together with equations (13.15), the time derivatives in the
extrinsic curvature are the directed time derivatives ∂0γab of the spatial tetrad metric, and the directed time
derivatives ∂0ecβ of the spatial inverse vierbein. However, these time derivatives appear only the combination
∂0γab − dab0 − dba0 = eaαeb
β∂0(γcd ecα e
dβ) = ea
αebβ ∂0gαβ . (13.33)
Thus the ADM Lagrangian picks out the natural coordinates as being the spatial components gαβ of the
coordinate metric, since only time derivatives of these appear in the ADM Lagrangian. An expression for
the extrinsic curvature that demonstrates explicitly its dependence on the time derivatives of the spatial
coordinate metric gαβ is, from manipulating equation (13.14b),
Kab =1
2α
(
eaαeb
β ∂gαβ
∂t−Daebt −Dbeat
)
. (13.34)
Conjugate momenta are obtained by differentiating the Lagrangian with respect to the velocities. The
derivatives of the extrinsic curvature with respect to the velocities gαβ ≡ ∂gαβ/∂t are
∂Kab
∂gαβ=
1
2αea
αebβ . (13.35)
The conjugate momenta παβ are therefore
παβ ≡ α ∂LADM
∂gαβ=
1
16π
(
Kαβ − gαβK)
, (13.36)
13.7 ADM equations of motion 199
in which the factor of α is introduced to convert the conjugate momenta into a tensor, as opposed to a tensor
density. Projected into the tetrad frame, the conjugate momenta are
πab = eaαe
bβ π
αβ =1
16π
(
Kab − γabK)
. (13.37)
The ADM Lagrangian (13.32) can be rewritten in terms of the conjugate momenta (13.37) as
LADM = 2Kabπab − G0
0
8π, (13.38)
in which the 3D Ricci scalar R(3) has been eliminated in favour of the time-time component G00 of the tetrad-
frame Einstein tensor by equation (13.23a). Substituting the expression (13.34) for the extrinsic curvature
Kab brings the ADM Lagrangian to
LADM =1
α
(
eaαeb
β ∂gαβ
∂t−Daebt −Dbeat
)
πab − G00
8π. (13.39)
The two terms Daebt and Dbeat can be combined into one because of the symmetry of πab, and integrated
by parts to give the ADM action
SADM ≡∫ tf
ti
LADM dx4 = −2
∫
∂V
ebtπab dx3
a +
∫ tf
ti
(
eaαeb
β ∂gαβ
∂tπab + 2 eatDbπ
ab − αG00
8π
)
dtdx30 .
(13.40)
Once again, for the purposes of extremizing the action, the surface term can be discarded. Eliminating the
Dbπab term in favour of the time-space part G0
a of the tetrad-frame Einstein tensor by equation (13.23b)
produces
SADM =
∫ tf
ti
(
eaαeb
β ∂gαβ
∂tπab − ea
tG0
a
8π− e0t
G00
8π
)
dtdx30 . (13.41)
This is the ADM gravitational action in desired Hamiltonian form. Compactly,
SADM =
∫ tf
ti
(
∂gαβ
∂tπαβ −H
)
dtdx30 (13.42)
where H is the Hamiltonian
H ≡ emtG
0m
8π=G0
t
8π. (13.43)
13.7 ADM equations of motion
Equations of motion follow from extremizing the ADM action (13.42) in Hamiltonian form. To accomplish
this, the Hamiltonian H must be expressed in terms of the coordinates and momenta. The expression (13.42)
for the ADM action shows that the natural dynamical coordinates and momenta of the system are the
200 ∗The 3+1 (ADM) formalism
components of the spatial coordinate-frame metric gαβ and their conjugate momenta παβ . In terms of these,
the Hamiltonian is
H =1
8π
(
αG00 + βαG0
α
)
= α
16π[
παβπαβ − (πγ
γ )2]
− R(3)
16π
− βα
2Dβπαβ
, (13.44)
where the expressions for G00 and G0
α are from equations (13.23a) and (13.23b) recast into coordinate-frame
quantities, the extrinsic curvature Kαβ being eliminated in favour of the momenta παβ by equation (13.36).
The expression (13.44) for the Hamiltonian depends not only on the coordinates gαβ and momenta παβ , but
also on the lapse α and shift βα, which are independent of gαβ and παβ . Consequently the lapse and shift must
also be treated as additional coordinates. The Hamiltonian (13.44) depends on the lapse and shift linearly,
the quantities in braces in equation (13.44) being independent of the lapse and shift. The Hamiltonian (13.44)
also depends on spatial derivatives ∂gαβ/∂xγ of the coordinates through its dependence on the Riemann
3-scalar R(3) and on the spatial connections Γαβγ associated with the covariant spatial coordinate derivative
Dβ in the term Dβπαβ . However, the spatial derivatives ∂gαβ/∂xγ are to be considered as determined by
the coordinates gαβ as a function of the spatial coordinates xγ on a spatial hypersurface of constant time t,
not as independent coordinates to be varied separately.
Variation of the ADM action (13.42) with respect to the lapse α, the shift βα, the coordinates gαβ, and
the momenta παβ , gives
δSADM =
[∫
παβδgαβ dx30
]tf
ti
(13.45)
+
∫ tf
ti
∂H
∂αδα+
∂H
∂βαδβα −
(
∂παβ
∂t+
∂H
∂gαβ
)
δgαβ +
(
∂gαβ
∂t− ∂H
∂παβ
)
δπαβ
dtdx30 .
The least action principle asserts that the action is minimized along the actual path taken by the system
between fixed initial and final conditions at ti and tf . Setting the variation (13.45) of the action equal to
zero with respect to arbitrary variations δα and δβα of the lapse and shift yields the constraint equations
∂H
∂α= 0 ,
∂H
∂βα= 0 . (13.46)
Setting the variation (13.45) of the action equal to zero with respect to arbitrary variations δgαβ and δπαβ
of the coordinates and momenta yields Hamilton’s equations
∂παβ
∂t= − ∂H
∂gαβ,
∂gαβ
∂t=
∂H
∂παβ. (13.47)
13.8 Constraints and energy-momentum conservation
14
∗The geometric algebra
The geometric algebra is an intuitively appealing formalism that draws together several mathematical threads
relevant to special and general relativity:
1. What is the best way to conceptualize Lorentz transformations (§14.6), and to implement them on a
computer (§14.20)?
2. How is it that bivectors in 4D spacetime have a natural complex structure, which has been seen to be the
heart of the Newman-Penrose formalism? See §14.17.
3. How can spin- 12 objects be incorporated into general relativity? What is a Dirac spinor? See §14.23.
4. How and why do differential forms work?
This chapter starts by setting up the geometric algebra in n-dimensional Euclidean space Rn, then gener-
alizes to Minkowski space, where the geometric algebra is called the spacetime algebra. The 4D spacetime
algebra proves to be identical to the Clifford algebra of the Dirac γ-matrices (which explains the adoption
of the symbol γγm to denote the basis vectors of a tetrad). Although the formalism is presented initially
in Euclidean or Minkowski space, everything generalizes immediately to general relativity, where the basis
vectors γγm form the basis of an orthonormal tetrad at each point of spacetime.
One convention adopted here, which agrees with the convention adopted by OpenGL and the computer
graphics industry, but is opposite to the standard physics convention, is that a rotor R rotates a multivector
a as a→ RaR, equation (14.35). This, along with the standard definition (14.16) for the pseudoscalar, has
the consequence that a right-handed rotation corresponds to R = eiθ/2 with θ increasing, and that rotations
accumulate to the right, that is, a rotation R followed by a rotation S is the product RS. By contrast, in the
standard physics convention a → RaR, a right-handed rotation corresponds to R = e−iθ/2, and rotations
accumulate to the left, that is, R followed by S is SR. The convention adopted here also means that a Weyl
or Dirac spinor ϕ is isomorphic to a scaled reverse rotor R, not to R as in the standard physics convention.
In this chapter, boldface denotes a multivector. A rotor is written in normal (not bold) face as a reminder
that, even though a rotor is an even member of the geometric algebra, it can also be regarded as a spin- 12
object with a transformation law (14.37) different from that (14.35) of multivectors. Later Latin indices
m,n, ... run over both time and space indices 0, 1, 2, 3, while earlier Latin indices i, j, k run over spatial
indices 1, 2, 3 only.
202 ∗The geometric algebra
a
θa
b
a
b
c
Figure 14.1 Multivectors of grade 1, 2, and 3: a vector a (left), a bivector a∧ b (middle), and a trivectora∧ b∧ c (right).
14.1 Products of vectors
In 3-dimensional Euclidean space R3, there are two familiar ways of taking the product of two vectors, the
scalar product and the vector product.
1. The scalar product a · b, also known as the dot product or inner product, of two vectors a and b is
a scalar of magnitude |a| |b| cos θ, where |a| and |b| are the lengths of the two vectors, and θ the angle
between them. The scalar product is commutative, a · b = b · a.
2. The vector product, a×b, also known as the cross product, is a vector of magnitude |a| |b| sin θ, directed
perpendicular to both a and b, such that a, b, and a× b form a right-handed set. The vector product is
anticommutative, a× b = −b× a.
The definition of the scalar product continues to work fine in a Euclidean space of any dimension, but the
definition of the vector product works only in three dimensions, because in two dimensions there is no vector
perpendicular to two vectors, and in four or more dimensions there are many vectors perpendicular to two
vectors. It is therefore useful to define a more general version, the outer product (H. Grassmann, 1862 Die
Ausdehnungslehre, Berlin) that works in Euclidean space Rn of any dimension.
3. The outer product a∧ b, also known as the wedge product, of two vectors a and b is a bivector, a
multivector of dimension 2, or grade 2. The bivector a∧b is the directed 2-dimensional area, of magnitude
|a| |b| sin θ, of the parallelogram formed by the vectors a and b, as illustrated in Figure 14.1. The bivector
has an orientation, or handedness, defined by circulating the parallelogram first along a, then along b.
The outer product is anticommutative, a∧ b = −b∧a, like its forebear the vector product.
The outer product can be repeated, so that (a∧ b)∧ c is a trivector, a directed volume, a multivector
of grade 3. The magnitude of the trivector is the volume of the parallelepiped defined by the vectors a, b,
and c, illustrated in Figure 14.1. The outer product is by construction associative, (a∧ b)∧ c = a∧ (b∧ c).
Associativity, together with anticommutativity of bivectors, implies that the trivector a∧b∧ c is totally
antisymmetric under permutations of the three vectors, that is, it is unchanged under even permutations, and
changes sign under odd permutations. The ordering of an outer product thus defines one of two handednesses.
14.2 Geometric product 203
It is a familiar concept that a vector a can be regarded as a geometric object, a directed length, independent
of the coordinates used to describe it. The components of a vector change when the reference frame changes,
but the vector itself remains the same physical thing. In the same way, a bivector a∧ b is a directed area,
and a trivector a∧b∧ c is a directed volume, both geometric objects with a physical meaning independent
of the coordinate system.
In two dimensions the triple outer product of any three vectors is zero, a∧ b∧ c = 0, because the volume
of a parallelepiped confined to a plane is zero. More generally, in n-dimensional space Rn, the outer product
of n+ 1 vectors is zero
a1 ∧a2 ∧ · · · ∧an+1 = 0 (n dimensions) . (14.1)
14.2 Geometric product
The inner and outer products offer two different ways of multiplying vectors. However, by itself neither
product conforms to the usual desideratum of multiplication, that the product of two elements of a set be
an element of the set. Taking the inner product of a vector with another vector lowers the dimension by
one, while taking the outer product raises the dimension by one.
H. Grassmann H. (1877, “Der ort der Hamilton’schen quaternionen in der audehnungslehre,” Math. Ann.
12, 375) and W. K. Clifford (1878, “Applications of Grassmann’s extensive algebra,” Am. J. Math. 1, 350)
resolved the problem by defining a multivector as any linear combination of scalars, vectors, bivectors, and
objects of higher grade. Let γγ1,γγ2, ...,γγn form an orthonormal basis for n-dimensional Euclidean space Rn.
A multivector in n = 2 dimensions is then a linear combination of
1 ,
1 scalar
γγ1 , γγ2 ,
2 vectors
γγ1 ∧γγ2 ,
1 bivector(14.2)
forming a linear space of dimension 1 + 2 + 1 = 4 = 22. Similarly, a multivector in n = 3 dimensions is a
linear combination of
1 ,
1 scalar
γγ1 , γγ2 , γγ3 ,
3 vectors
γγ1 ∧γγ2 , γγ2 ∧γγ3 , γγ3 ∧γγ1 ,
3 bivectors
γγ1 ∧γγ2 ∧γγ3 ,
1 trivector(14.3)
forming a linear space of dimension 1 + 3 + 3 + 1 = 8 = 23. In general, multivectors in n dimensions form a
linear space of dimension 2n, with n!/[m!(n−m)!] distinct basis elements of grade m.
A multivector a in n-dimensional Euclidean space Rn can thus be written as a linear combination of basis
elements
a =∑
distinct i,j,...,m⊆1,2,...,naij...m γγi ∧γγj ∧ ...∧γγm (14.4)
the sum being over all 2n distinct subsets of 1, 2, ..., n. The index on each component aij...m is a totally
antisymmetric quantity, reflecting the total antisymmetry of γγi ∧γγj ∧ ...∧γγm.
The point of introducing multivectors is to allow multiplication to be defined so that the product of two
204 ∗The geometric algebra
multivectors is a multivector. The key trick is to define the geometric product ab of two vectors a and b
to be the sum of their inner and outer products:
ab = a · b + a∧ b . (14.5)
That is a seriously big trick, and if you buy a ticket to it, you are in for a seriously big ride. As a particular
example of (14.5), the geometric product of any element γγi of the orthonormal basis with itself is a scalar,
and with any other element of the basis is a bivector:
γγiγγj =
1 (i = j)
γγi ∧γγj (i 6= j) .(14.6)
Conversely, the rules (14.6), plus distributivity, imply the multiplication rule (14.5). A generalization of the
rule (14.6) completes the definition of the geometric product:
γγiγγj ...γγm = γγi ∧γγj ∧ ...∧γγm (i, j, ...,m all distinct) . (14.7)
The rules (14.6) and (14.7), along with the usual requirements of associativity and distributivity, combined
with commutativity of scalars and anticommutativity of pairs of γγi, uniquely define multiplication over the
space of multivectors. For example, the product of the bivector γγ1 ∧γγ2 with the vector γγ1 is
(γγ1 ∧γγ2)γγ1 = γγ1γγ2γγ1 = −γγ2γγ1γγ1 = −γγ2 . (14.8)
Sometimes it is convenient to denote the wedge product (14.7) of distinct basis elements by the abbreviated
symbol γγij...m
γγij...m ≡ γγi ∧γγj ∧ ...∧γγm (i, j, ...,m all distinct) . (14.9)
By construction, γγij...m is antisymmetric in its indices. The product of two general multivectors a = aαγγα
and b = bαγγα, with paired indices implicitly summed over distinct subsets of 1, ..., n, is
ab = aαbβγγαγγβ . (14.10)
Does the geometric algebra form a group under multiplication? No. One of the defining properties of a
group is that every element should have an inverse. But, for example,
(1 + γγ1)(1− γγ1) = 0 (14.11)
shows that neither 1 + γγ1 nor 1− γγ1 has an inverse.
14.3 Reverse
The reverse of any basis element is defined to be the reversed product
γγi ∧γγj ∧ ...∧γγm ≡ γγm ∧ ...∧γγj ∧γγi . (14.12)
14.4 The pseudoscalar and the Hodge dual 205
The reverse a of any multivector a is the multivector obtained by reversing each of its components. Reversion
leaves unchanged all multivectors whose grade is 0 or 1, modulo 4, and changes the sign of all multivectors
whose grade is 2 or 3, modulo 4. For example, scalars and vectors are unchanged by reversion, but bivectors
and trivectors change sign. Reversion satisfies
a + b = a + b , (14.13)
ab = ba . (14.14)
Among other things, it follows that the reverse of any product of multivectors is the reversed product, as
you would hope:
ab ... c = c ... ba . (14.15)
14.4 The pseudoscalar and the Hodge dual
Orthogonal to any m-dimensional subspace of n-dimensional space is an (n−m)-dimensional space, called
the Hodge dual space. For example, the Hodge dual of a bivector in 2 dimensions is a 0-dimensional
object, a pseudoscalar. Similarly, the Hodge dual of a bivector in 3 dimensions is a 1-dimensional object, a
pseudovector.
Define the pseudoscalar in in n dimensions to be
in ≡ γγ1 ∧γγ2 ∧ ...∧γγn (14.16)
with reverse
in = (−)[n/2]γγ1 ∧γγ2 ∧ ...∧γγn . (14.17)
The quantity [n/2] in equation (14.17) signifies the largest integer less than or equal to n/2. The square of
the pseudoscalar is
i2n = (−)[n/2] =
1 if n = (0 or 1) modulo 4
−1 if n = (2 or 3) modulo 4 .(14.18)
The pseudoscalar anticommutes (commutes) with vectors a, that is, with multivectors of grade 1, if n is
even (odd):
ina = −ain if n is even
ina = ain if n is odd .(14.19)
This implies that the pseudoscalar in commutes with all even grade elements of the geometric algebra, and
that it anticommutes (commutes) with all odd elements of the algebra if n is even (odd).
Exercise 14.1 Prove that the only multivectors that commute with all elements of the algebra are linear
combinations of the scalar 1 and, if n is odd, the pseudoscalar in. ⋄
206 ∗The geometric algebra
The Hodge dual ∗a of a multivector a in n dimensions is defined by premultiplication by the pseudoscalar
in,
∗a ≡ ina . (14.20)
In 3 dimensions, the Hodge duals of the basis vectors γγi are the bivectors
i3γγ1 = γγ2 ∧γγ3 , i3γγ2 = γγ3 ∧γγ1 , i3γγ3 = γγ1 ∧γγ2 . (14.21)
Thus in 3 dimensions the bivector a∧ b is seen to be the pseudovector Hodge dual to the familiar vector
product a× b:
a∧ b = i3 a× b . (14.22)
14.5 Reflection
Multiplying a vector (a multivector of grade 1) by a vector shifts the grade (dimension) of the vector by
±1. Thus, if one wants to transform a vector into another vector (with the same grade, one), at least two
multiplications by a vector are required.
The simplest non-trivial transformation of a vector a is
n : a→ nan (14.23)
in which the vector a is multiplied on both left and right with a unit vector n. If a is resolved into components
a‖ and a⊥ respectively parallel and perpendicular to n, then the transformation (14.23) is
n : a‖ + a⊥ → a‖ − a⊥ (14.24)
which represents a reflection of the vector a through the axis n, a reversal of all components of the vector
n
a
−nan
nan
Figure 14.2 Reflection of a vector a through axis n.
14.6 Rotation 207
perpendicular to n, as illustrated by Figure 14.2. Note that −nan is the reflection of a through the
hypersurface normal to n, a reversal of the component of the vector parallel to n.
The operation of left- and right-multiplying by a unit vector n reflects not only vectors, but multivectors
a in general:
n : a→ nan . (14.25)
For example, the product ab of two vectors transforms as
n : ab→ n(ab)n = (nan)(nbn) (14.26)
which works because n2 = 1.
A reflection leaves any scalar λ unchanged, n : λ→ nλn = λn2 = λ. Geometrically, a reflection preserves
the lengths of, and angles between, all vectors.
14.6 Rotation
Two successive reflections yield a rotation. Consider reflecting a vector a (a multivector of grade 1) first
through the unit vector m, then through the unit vector n:
mn : a→ nmamn . (14.27)
Any component a⊥ of a simultaneously orthogonal to both m and n (i.e. m ·a⊥ = n ·a⊥ = 0) is unchanged
by the transformation (14.27), since each reflection flips the sign of a⊥:
mn : a⊥ → nma⊥mn = −na⊥n = a⊥ . (14.28)
Rotations inherit from reflections the property of preserving the lengths of, and angles between, all vectors.
Thus the transformation (14.27) must represent a rotation of those components a‖ of a lying in the 2-dim-
ensional plane spanned by m and n, as illustrated by Figure 14.3. To determine the angle by which the
plane is rotated, it suffices to consider the case where the vector a‖ is equal to m (or n, as a check). It is
not too hard to figure out that, if the angle from m to n is θ/2, then the rotation angle is θ in the same
sense, from m to n.
mamam
n
nmamnθ
Figure 14.3 Rotation of a vector a by the bivector mn. Baffled? Hey, draw your own picture.
208 ∗The geometric algebra
For example, if m and n are parallel, so that m = ±n, then the angle between m and n is θ/2 = 0 or π,
and the transformation (14.27) rotates the vector a‖ by θ = 0 or 2π, that is, it leaves a‖ unchanged. This
makes sense: two reflections through the same plane leave everything unchanged. If on the other hand m
and n are orthogonal, then the angle between them is θ/2 = ±π/2, and the transformation (14.27) rotates
a‖ by θ = ±π, that is, it maps a‖ to −a‖.
The rotation (14.27) can be abbreviated
R : a→ RaR (14.29)
where R = mn is called a rotor, and R = nm is its reverse. Rotors are unimodular, satisfying RR =
RR = 1. According to the discussion above, the transformation (14.29) corresponds to a rotation by angle
θ in the m–n plane if the angle from m to n is θ/2. Then m · n = cos θ/2 and m∧n = (γγ1 ∧γγ2) sin θ/2,
where γγ1 and γγ2 are two orthonormal vectors spanning the m–n plane, oriented so that the angle from γγ1
to γγ2 is positive π/2 (i.e. γγ1 is the x-axis and γγ2 the y-axis). Note that the outer product γγ1 ∧γγ2 is invariant
under rotations in the m–n plane, hence independent of the choice of orthonormal basis vectors γγ1 and γγ2.
It follows that the rotor R corresponding to a right-handed rotation by θ in the γγ1–γγ2 plane is given by
R = cosθ
2+ (γγ1 ∧γγ2) sin
θ
2. (14.30)
It is straightforward to check that the rotor (14.30) rotates the basis vectors γγi as
R : γγ1 → Rγγ1R = γγ1 cos θ + γγ2 sin θ , (14.31a)
R : γγ2 → Rγγ2R = γγ2 cos θ − γγ1 sin θ , (14.31b)
R : γγi → RγγiR = γγi (i 6= 1, 2) , (14.31c)
which indeed corresponds to a right-handed rotation by angle θ in the γγ1–γγ2 plane. The inverse rotation is
R : a→ RaR (14.32)
with
R = cosθ
2− (γγ1 ∧γγ2) sin
θ
2. (14.33)
A rotation of the form (14.30), a rotation in a single plane, is called a simple rotation.
A rotation first by R and then by S transforms a vector a as
RS : a→ SRaRS = RS aRS . (14.34)
Thus the composition of two rotations, first R and then S, is given by their geometric product RS. In three
dimensions or less, all rotations are simple, but in four dimensions or higher, compositions of simple rotations
can yield rotations that are not simple. For example, a rotation in the γγ1–γγ2 plane followed by a rotation in
the γγ3–γγ4 plane is not equivalent to any simple rotation. However, it will be seen in §14.17 that bivectors
in the 4D spacetime algebra have a natural complex structure, which allows 4D spacetime rotations to take
14.7 A rotor is a spin- 12 object 209
a simple form similar to (14.30), but with complex angle θ and two orthogonal planes of rotation combined
into a complex pair of planes.
Simple rotors are both even and unimodular, and composition preserves those properties. A rotor R is
defined in general to be any even, unimodular (RR = 1) element of the geometric algebra. The set of rotors
defines a group the rotor group, also referred to here as the rotation group. A rotor R rotates not only
vectors, but multivectors a in general:
R : a→ RaR . (14.35)
For example, the product ab of two vectors transforms as
R : ab→ R(ab)R = (RaR)(RbR) (14.36)
which works because RR = 1.
Concept question 14.2 If vectors rotate twice as fast as rotors, do bivectors rotate twice as fast as
vectors? What happens to a bivector when you rotate it by π radians? Construct a mental picture of a
rotating bivector. ⋄
To summarize, the characterization of rotations by rotors has considerable advantages. Firstly, the trans-
formation (14.35) applies to multivectors a of arbitrary grade in arbitrarily many dimensions. Secondly,
the composition law is particularly simple, the composition of two rotations being given by their geometric
product. A third advantage is that rotors rotate not only vectors and multivectors, but also spin- 12 objects
— indeed rotors are themselves spin- 12 objects — as might be suspected from the intriguing factor of 1
2 in
front of the angle θ in equation (14.30).
14.7 A rotor is a spin-12
object
A rotor was defined in the previous section, §14.6, as an even, unimodular element of the geometric algebra.
As a multivector, a rotor R would transform under a rotation by the rotor S as R → SRS. As a rotor,
however, the rotor R transforms under a rotation by the rotor S as
S : R→ RS , (14.37)
according to the transformation law (14.34). That is, composition in the rotor group is defined by the
transformation (14.37): R rotated by S is RS.
The expression (14.30) for a simple rotation in the γγ1–γγ2 plane shows that the rotor corresponding to a
rotation by 2π is −1. Thus under a rotation (14.37) by 2π, a rotor R changes sign:
2π : R→ −R . (14.38)
A rotation by 4π is necessary to bring the rotor R back to its original value:
4π : R→ R . (14.39)
210 ∗The geometric algebra
Thus a rotor R behaves like a spin- 12 object, requiring 2 full rotations to restore it to its original state.
The two different transformation laws for a rotor — as a multivector, and as a rotor — describe two
different physical situations. The transformation of a rotor as a multivector answers the question, what is
the form of a rotor R rotated into another, primed, frame? In the unprimed frame, the rotor R transforms
a multivector a to RaR. In the primed frame rotated by rotor S from the unprimed frame, a′ = SaS, the
transformed rotor is SRS, since
a′ = SaS → SRaRS = SRSa′SRS . (14.40)
By contrast, the transformation (14.37) of a rotor as a rotor answers the question, what is the rotor corre-
sponding to a rotation R followed by a rotation S?
14.8 A multivector rotation is an active rotation
In most of the rest of this book, indices indicate how an object transforms, so that the notation
amγγm (14.41)
indicates a scalar, an object that is unchanged by a transformation, because the transformation of the con-
travariant vector am cancels against the corresponding transformation of the covariant vector γγm. However,
the transformation (14.35) of a multivector is to be understood as an active transformation that rotates the
basis vectors γγα while keeping the coefficients aα fixed, as opposed to a passive transformation that rotates
the tetrad while keeping the thing itself unchanged. Thus a multivector a ≡ aαγγα (implicit summation over
α ⊆ 1, ..., n) is not a scalar under the transformation (14.35), but rather transforms to the multivector
a′ ≡ aαγγ′α given by
R : aαγγα → aαRγγαR = aαγγ′α . (14.42)
An explicit example is the transformation (14.31) of the tetrad axes γγi under a right-handed rotation by
angle θ.
14.9 2D rotations and complex numbers
Section 14.6 identified the rotation group in n dimensions with the geometric subalgebra of even, unimodular
multivectors. In two dimensions, the even grade multivectors are linear combinations of the basis set
1 ,
1 scalar
i2 ,
1 bivector (pseudoscalar)(14.43)
forming a linear space of dimension 2. The sole bivector is the pseudoscalar i2 ≡ γγ1 ∧γγ2, equation (14.16),
the highest grade element in 2 dimensions. The rotor R that generates a right-handed rotation by angle θ
14.9 2D rotations and complex numbers 211
is, according to equation (14.30),
R = eθ/2 = ei2 θ/2 = cosθ
2+ i2 sin
θ
2, (14.44)
where θ = i2 θ is the bivector of magnitude θ.
Since the square of the pseudoscalar i2 is minus one, the pseudoscalar resembles the pure imaginary i, the
square root of −1. Sure enough, the mapping
i2 ↔ i (14.45)
defines an isomorphism between the algebra of even grade multivectors in 2 dimensions and the field of
complex numbers
a+ i2b↔ a+ i b . (14.46)
With the isomorphism (14.46), the rotor R that generates a right-handed rotation by angle θ is equivalent
to the complex number
R = eiθ/2 . (14.47)
The associated reverse rotor R is
R = e−iθ/2 , (14.48)
the complex conjugate of R. The group of 2D rotors is isomorphic to the group of complex numbers of unit
magnitude, the unitary group U(1),
2D rotors↔ U(1) . (14.49)
Let z denote an even multivector, equivalent to some complex number by the isomorphism (14.46). Ac-
cording to the transformation formula (14.35), under the rotationR = eiθ/2, the even multivector, or complex
number, z transforms as
R : z → e−iθ/2z eiθ/2 = e−iθ/2eiθ/2z = z (14.50)
which is true because even multivectors in 2 dimensions commute, as complex numbers should. Equa-
tion (14.50) shows that the even multivector, or complex number, z is unchanged by a rotation. This might
seem strange: shouldn’t the rotation rotate the complex number z by θ in the Argand plane? The an-
swer is that the rotation R : a → RaR rotates vectors γγ1 and γγ2 (Exercise 14.3), as already seen in the
transformation (14.31). The same rotation leaves the scalar 1 and the bivector i2 ≡ γγ1 ∧γγ2 unchanged. If
temporarily you permit yourself to think in 3 dimensions, you see that the bivector γγ1 ∧γγ2 is Hodge dual to
the pseudovector γγ1 × γγ2, which is the axis of rotation and is itself unchanged by the rotation, even though
the individual vectors γγ1 and γγ2 are rotated.
Exercise 14.3 Confirm that a right-handed rotation by angle θ rotates the axes γγi by
R : γγ1 → e−iθ/2γγ1 eiθ/2 = γγ1 cos θ + γγ2 sin θ , (14.51a)
R : γγ2 → e−iθ/2γγ2 eiθ/2 = γγ2 cos θ − γγ1 sin θ , (14.51b)
212 ∗The geometric algebra
in agreement with (14.31). The important thing to notice is that the pseudoscalar i2, hence i, anticommutes
with the vectors γγi. ⋄
14.10 Quaternions
A quaternion can be regarded as a kind of souped-up complex number
q = a+ ıb1 + b2 + kb3 , (14.52)
where a and bi (i = 1, 2, 3) are real numbers, and the three imaginary numbers ı, , k, also denoted ı1, ı2, ı3here for convenience and brevity, are defined to satisfy1
ı2 = 2 = k2 = −ık = −1 . (14.53)
Remark the dotless ı, to distinguish these quaternionic imaginaries from other possible imaginaries. A
consequence of equations (14.53) is that each pair of imaginary numbers anticommutes:
ı = −ı = −k , k = −k = −ı , kı = −ık = − . (14.54)
A quaternion (14.52) can be expressed compactly as a sum of its scalar, a, and vector (actually pseudovector,
as will become apparent below from the isomorphism (14.68)), ı · b, parts
q = a+ ı · b , (14.55)
where ı is shorthand for the triple of quaternionic imaginaries,
ı ≡ ı, , k ≡ ı1, ı2, ı3 , (14.56)
and where b ≡ b1, b2, b3, and ı · b ≡ ıibi (implicit summation over i = 1, 2, 3) is the usual Euclidean dot
product. A fundamentally useful formula, which follows from the defining equations (14.53), is
(ı · a)(ı · b) = −a · b− ı · (a× b) (14.57)
where a × b is the usual 3D vector product. The product of two quaternions p ≡ a+ ı · b and q ≡ c+ ı · dcan thus be written
pq = (a+ ı · b)(c + ı · d) = ac− b · d + ı · (ad + cb− b× d) . (14.58)
The quaternionic conjugate q of a quaternion q ≡ a + ı · b is (the overbar symbol ¯ for quaternionic
conjugation distinguishes it from the asterisk symbol ∗ for complex conjugation)
q = a− ı · b . (14.59)
1 The choice ık = 1 in the definition (14.53) is the opposite of the conventional definition ijk = −1 famously carved by W.R. Hamilton in the stone of Brougham Bridge while walking with his wife along the Royal Canal to Dublin on 16 October1843 (S. O’Donnell, 1983, William Rowan Hamilton: Portrait of a Prodigy, Boole Press, Dublin). To map to Hamilton’sdefinition, you can take ı = −i, = −j, k = −k, or alternatively ı = i, = −j, k = k, or ı = k, = j, k = i. The adoptedchoice ık = 1 has the merit that it avoids a treacherous minus sign in the isomorphism (14.68) between 3-dimensionalpseudovectors and quaternions. The present choice also conforms to the convention used by OpenGL and other computergraphics programs.
14.11 3D rotations and quaternions 213
The quaternionic conjugate of a product is the reversed product of quaternionic conjugates
pq = qp (14.60)
just like reversion in the geometric algebra, equation (14.14) (the choice of the same symbol, an overbar,
to represent both reversion and quaternionic conjugation is not coincidental). The magnitude |q| of the
quaternion q ≡ a+ ı · b is
|q| = (qq)1/2 = (qq)1/2 = (a2 + b · b)1/2 = (a2 + b21 + b22 + b23)1/2 . (14.61)
The inverse q−1 of the quaternion, satisfying qq−1 = q−1q = 1, is
q−1 = q/(qq) = (a− ı · b)/(a2 + b · b) = (a− ı1b1 − ı2b2 − ı3b3)/(a2 + b21 + b22 + b23) . (14.62)
14.11 3D rotations and quaternions
As before, the rotation group is the group of even, unimodular multivectors of the geometric algebra. In
three dimensions, the even grade multivectors are linear combinations of the basis set
1 ,
1 scalar
i3γγ1 , i3γγ2 , i3γγ3 ,
3 bivectors (pseudovectors)(14.63)
forming a linear space of dimension 4. The three bivectors are pseudovectors, equation (14.21). The squares
of the pseudovector basis elements are all minus one,
(i3γγ1)2 = (i3γγ2)
2 = (i3γγ3)2 = −1 , (14.64)
and they anticommute with each other,
(i3γγ1)(i3γγ2) = −(i3γγ2)(i3γγ1) = −i3γγ3 ,
(i3γγ2)(i3γγ3) = −(i3γγ3)(i3γγ2) = −i3γγ1 , (14.65)
(i3γγ3)(i3γγ1) = −(i3γγ1)(i3γγ3) = −i3γγ2 .
The rotor R that generates a rotation by angle θ right-handedly about unit vector n in 3 dimensions is,
according to equation (14.30),
R = eθ/2 = ei3 n θ/2 = cosθ
2+ i3 n sin
θ
2. (14.66)
where θ is the bivector
θ ≡ i3 n θ (14.67)
of magnitude θ and unit vector direction n ≡ γγini.
Comparison of equations (14.64) and (14.65) to equations (14.53) and (14.54), shows that the mapping
i3γγi ↔ ıi (i = 1, 2, 3) (14.68)
214 ∗The geometric algebra
defines an isomorphism between the space of even multivectors in 3 dimensions and the non-commutative
division algebra of quaternions
a+ i3γγibi ↔ a+ ıibi . (14.69)
With the equivalence (14.69), the rotor R that generates a rotation by angle θ right-handedly about unit
vector n in 3 dimensions is equivalent to the quaternion
R = eθ/2 = eı·n θ/2 = cosθ
2+ ı · n sin
θ
2, (14.70)
where θ is the pseudovector quaternion
θ ≡ ı · n θ ≡ (ı1n1 + ı2n2 + ı3n3) θ (14.71)
whose magnitude is |θ| = θ and whose unit direction is θ ≡ θ/θ = ı · n. The associated reverse rotor R is
R = e−θ/2 = e−ı·n θ/2 = cosθ
2− ı · n sin
θ
2, (14.72)
the quaternionic conjugate of R.
The group of rotors is isomorphic to the group of unit quaternions, quaternions q = a+ ı1b1 + ı2b2 + ı3b3satisfying qq = a2 + b21 + b22 + b23 = 1. Unit quaternions evidently define a unit 3-sphere in the 4-dimensional
space of coordinates a, b1, b2, b3. From this it is apparent that the rotor group in 3 dimensions has the
geometry of a 3-sphere S3.
Exercise 14.4 This exercise is a precursor to Exercise 14.15. Let b ≡ γγibi be a 3D vector, a multivector of
grade 1 in the 3D geometric algebra. Use the quaternionic composition rule (14.57) to show that the vector
b transforms under a right-handed rotation by angle θ about unit direction n = γγini as
R : b→ R bR = b + 2 sinθ
2n×
(
cosθ
2b + sin
θ
2n× b
)
. (14.73)
Here the cross-product n × b denotes the usual vector product, which is dual to the bivector product
n∧ b, equation (14.22). Suppose that the quaternionic components of the rotor R are w, x, y, z, that is,
R = eı·n θ/2 = w+ı1x+ı2y+ı3z. Show that the transformation (14.73) is (note that the 3×3 rotation matrix
is written to the right of the vector, in accordance with the computer graphics convention that rotations
accumulate to the right — opposite to the physics convention; to recover the physics convention, take the
transpose):
R :(
b1 b2 b2)
→(
b1 b2 b3)
w2+x2−y2−z2 2(xy+wz) 2(zx−wy)2(xy−wz) w2−x2+y2−z2 2(yz+wx)
2(zx+wy) 2(yz−wx) w2−x2−y2+z2
. (14.74)
Confirm that the 3× 3 rotation matrix on the right hand side of the transformation (14.74) is an orthogonal
matrix (its inverse is its transverse) provided that the rotor is unimodular, RR = 1, so that w2+x2+y2+z2 =
14.12 Pauli matrices 215
1. As a simple example, show that the transformation (14.74) in the case of a right-handed rotation by angle
θ about the 3-axis (the 1–2 plane), where w = cos θ2 and z = sin θ
2 , is
R :(
b1 b2 b2)
→(
b1 b2 b3)
cos θ sin θ 0
− sin θ cos θ 0
0 0 1
. (14.75)
⋄
14.12 Pauli matrices
The Pauli matrices σ ≡ σi ≡ σ1, σ2, σ3 form a vector of 2 × 2 complex (with respect to a quantum-
mechanical imaginary i) matrices whose three components are each traceless (Tr σi = 0), Hermitian (σ†i = σi),
and unitary (σ†i σi = 1, no implicit summation):
σ1 ≡(
0 1
1 0
)
, σ2 ≡(
0 −ii 0
)
, σ3 ≡(
1 0
0 −1
)
. (14.76)
The Pauli matrices anticommute with each other
σ1σ2 = −σ1σ2 = iσ3 , σ2σ3 = −σ3σ2 = iσ1 , σ3σ1 = −σ1σ3 = iσ2 . (14.77)
The particular choice (14.76) of Pauli matrices is conventional but not unique: any three traceless, Hermitian,
unitary, anticommuting 2× 2 complex matrices will do. The product of the 3 Pauli matrices is i times the
unit matrix,
σ1σ2σ3 = i
(
1 0
0 1
)
. (14.78)
The multiplication rules of the Pauli matrices σi are identical to those of the basis vectors γγi of the 3D
geometric algebra. If the scalar 1 in the geometric algebra is identified with the unit 2× 2 matrix, and the
pseudoscalar i3 is identified with the imaginary i times the unit matrix, then the 3D geometric algebra is
isomorphic to the algebra generated by the Pauli matrices, the Pauli algebra, through the mapping
1↔(
1 0
0 1
)
, γγi ↔ σi , i3 ↔ i
(
1 0
0 1
)
. (14.79)
The 3D pseudoscalar i3 commutes with all elements of the 3D geometric algebra.
Concept question 14.5 The Pauli matrices are traceless, Hermitian, unitary, and anticommuting. What
do these properties correspond to in the geometric algebra? Are all these properties necessary for the Pauli
algebra to be isomorphic to the 3D geometric algebra? Are the properties sufficient?
The rotation group is the group of even, unimodular multivectors of the geometric algebra. The isomor-
phism (14.79) establishes that the rotation group is isomorphic to the group of complex 2 × 2 matrices of
216 ∗The geometric algebra
the form
a+ iσ · b , (14.80)
with a, bi (i = 1, 3) real, and with the unimodular condition requiring that a2+b ·b = 1. It is straightforward
to check (Exercise 14.6) that the group of such matrices constitutes the group of unitary complex 2 × 2
matrices of unit determinant, the special unitary group SU(2). The isomorphisms
a+ i3γγibi ↔ a+ ıibi ↔ a+ iσibi (14.81)
have thus established isomorphisms between the group of 3D rotors, the group of unit quaternions, and the
special unitary group of complex 2× 2 matrices
3D rotors↔ unit quaternions↔ SU(2) . (14.82)
An isomorphism that maps a group into a set of matrices, such that group multiplication corresponds to
ordinary matrix multiplication, is called a representation of the group. The representation of the rotation
group as 2×2 complex matrices may be termed the Pauli representation. The Pauli representation is the
lowest dimensional representation of the 3D rotation group.
Exercise 14.6 Translate a rotor into an element of SU(2). Show that the rotor R = ei3 n θ corre-
sponding to a right-handed rotation by angle θ about unit axis n ≡ n1, n2, n3 is equivalent to the special
unitary 2× 2 matrix
R↔(
cos θ2 + in3 sin θ
2 (n2 + in1) sin θ2
(−n2 + in1) sin θ2 cos θ
2 − in3 sin θ2
)
. (14.83)
Show that the reverse rotor R is equivalent to the Hermitian conjugate R† of the corresponding 2×2 matrix.
Show that the determinant of the matrix equals RR, which is 1. ⋄
14.13 Pauli spinors
In the Pauli representation, spin- 12 objects ϕ are Pauli spinors, 2-dimensional complex (with respect to i)
vectors
ϕ =
(
ϕ↑
ϕ↓
)
(14.84)
that are rotated by pre-multiplying by elements of the special unitary group SU(2). According to the
equivalence (14.83), a rotation by 2π is represented by minus the unit matrix,( −1 0
0 −1
)
. (14.85)
Consequently a rotation by 2π changes the sign of a Pauli spinor ϕ. A rotation by 4π is required to rotate
a Pauli spinor back to its original value. Thus a Pauli spinor indeed behaves like a spin- 12 object.
14.13 Pauli spinors 217
In quantum mechanics, the Pauli matrices σi, equations (14.76), provide a representation of the spin
operator s ≡ si ≡ s1, s2, s3 (equation (14.86) is in units ~ = 1; in standard units, s = ~
2σ)
s = 12 σ . (14.86)
The eigenvectors of the spin operator s · ζ projected along any axis ζ define objects of definite spin ± 12
measured along that axis. Each Pauli matrix σi has two eigenvalues ±1, and thus each spin operator
component si has two eigenvalues ± 12 . The eigenvectors of si are spin-up (eigenvalue + 1
2 ) and spin-down
(eigenvalue − 12 ), as measured along the i-axis. In particular, the normalized eigenvectors of σ3 are
↑ ≡(
1
0
)
, ↓ ≡(
0
1
)
, (14.87)
satisfying
s3 ↑ = 12 ↑ , s3 ↓ = − 1
2 ↓ . (14.88)
The Pauli spinor ϕ, equation (14.84), can thus be expressed
ϕ = ϕ↑ ↑ + ϕ↓ ↓ (14.89)
where ϕ↑ and ϕ↓ are the complex amplitudes along the up and down directions of the 3-axis (the z-axis).
Essential to quantum mechanics is the existence of an inner product. The inner product of two Pauli
spinors ϕ and ψ is the product ϕ†ψ of the Hermitian conjugate of ϕ with ψ. The Hermitian conjugate ϕ† of
the Pauli spinor ϕ (14.84) is
ϕ† =(
ϕ∗↑ ϕ∗
↓
)
. (14.90)
The magnitude squared of the spinor ϕ is the real number
|ϕ|2 = ϕ†ϕ = |ϕ↑|2 + |ϕ↓|2 . (14.91)
In quantum mechanics, the magnitude squared of the spinor |ϕ|2 is interpreted as the total probability (or
probability density) of the particle. The two parts |ϕ↑|2 and |ϕ↓|2 are the probabilities of the particle being
in the up and down states. The probabilities in the up and down states depend on the direction along which
the spin is measured, but the total probability |ϕ|2 is independent on the choice of direction.
Exercise 14.7 Orthonormal eigenvectors of the spin operator. Show that the orthonormal eigen-
vectors ↑ζ and ↓ζ of the spin operator s · ζ projected along the unit direction ζ ≡ ζ1, ζ2, ζ3 are
↑ζ =1
√
2(1 + ζ3)
(
1 + ζ3ζ1 + iζ2
)
, ↓ζ =1
√
2(1− ζ3)
( −1 + ζ3ζ1 + iζ2
)
. (14.92)
⋄
218 ∗The geometric algebra
14.14 Pauli spinors as scaled 3D rotors, or quaternions
Rotors and Pauli spinors both behave like spin- 12 objects, requiring a rotation of 4π to bring them full circle.
In fact a Pauli spinor is equivalent (14.96) to a (reverse) 3D rotor R scaled by a positive real scalar λ.
Consequently a Pauli spinor is equivalent (14.97) to a quaternion. The 2 complex degrees of freedom of the
Pauli spinor are equivalent to the 4 real degrees of freedom of the quaternion.
Start with the eigenequation (14.88) for the unit spin-up eigenvector ↑ in the 3-direction (z-direction),
s3 ↑ = 12 ↑ . (14.93)
If the spin operator s3 is rotated by rotor R, and the spin-up eigenvector ↑ is pre-multiplied by R, then the
eigenequation (14.93) transforms into an eigenequation for the unit Pauli spinor R ↑:
(Rs3R) (R ↑) = 12 (R ↑) (14.94)
(notice that the isomorphism (14.79) between the geometric algebra and the Pauli algebra guarantees that
the rule s3 → Rs3R for rotating the spin operator s3 is valid regardless of whether s3 and R are considered
as elements of the geometric algebra or as 2 × 2 complex matrices in the Pauli algebra). The Pauli spinor
↑, equation (14.87), is normalized to unit magnitude, and the rotated Pauli spinor R ↑ in equation (14.94)
is likewise of unit magnitude. A general Pauli spinor ϕ is the product of a real scalar λ and a rotated unit
spinor R ↑,ϕ = λR ↑ . (14.95)
The real scalar λ can be taken without loss of generality to be positive, since any minus sign can be absorbed
into a rotation by 2π of the rotor R. It is straightforward to check (Exercise 14.8) that any Pauli spinor ϕ
can be expressed in the form (14.95). Equation (14.95) establishes an equivalence between Pauli spinors and
reversed 3D rotors R scaled by a positive real scalar λ
ϕ↔ λR . (14.96)
Given the equivalence (14.82) between 3D rotors and unit quaternions, it follows that Pauli spinors are
equivalent to reverse quaternions (see Exercises 14.8 and 14.9 for the precise translation)
Pauli spinors↔ reverse quaternions . (14.97)
The equivalence means that there is a one-to-one correspondence between Pauli spinors and reverse quater-
nions, and that they transform in the same way under 3D rotations.
The Hermitian conjugate ϕ† of the Pauli spinor is
ϕ† = ↑†λR , (14.98)
where ↑†=(
1 0)
is the Hermitian conjugate of the spin-up eigenvector ↑. The squared amplitude of the
Pauli spinor
ϕ†ϕ = λ2 (14.99)
is the probability (or probability density) of the particle, which is unchanged by a rotation.
14.15 Spacetime algebra 219
One is used to thinking of a Pauli spinor as an instrinsically quantum mechanical object. The equiva-
lence (14.96) between Pauli spinors and scaled reverse rotors shows that Pauli spinors also have a classical
interpretation: they encode a real amplitude λ, and a rotation R. This provides a mathematical basis for
the idea that, through their spin, fundamental particles “know” about the rotational structure of space.
The spin axis ζ of a Pauli spinor χ is the direction along which the Pauli spinor is pure up. For example,
the spin axis of the of the spin-up eigenvector ↑ is the positive 3-axis (the z-axis), while the spin axis of the
of the spin-down eigenvector ↓ is the negative 3-axis. In general, the spin axis of a Pauli spinor (14.95) is
the unit direction ζ of the rotated 3-axis,
Rs3R = s · ζ . (14.100)
The rotor S corresponding to a right-handed rotation by angle ζ about the spin axis ζ is S = ei3ζ/2 where
ζ = ζζ. Such a rotation transforms R → SR, hence transforms the Pauli spinor (14.95) as ϕ → Sϕ.
Rotating the Pauli spinor ϕ right-handedly by angle ζ about its spin axis leaves the spin axis unchanged,
but multiplies the spinor by a phase e−iζ/2,
e−i3ζ/2 ϕ = e−iζ/2 ϕ . (14.101)
Exercise 14.8 Translate a Pauli spinor into a quaternion. Given any Pauli spinor ϕ ≡(
ϕ↑
ϕ↓
)
,
show that the corresponding scaled reverse rotor λR in the Pauli representation (14.76) is the unitary 2× 2
matrix
λR =
(
ϕ↑ −ϕ∗↓
ϕ↓ ϕ∗↑
)
. (14.102)
Show that the corresponding real quaternion is
q = λR = Reϕ↑, Imϕ↓,−Reϕ↓, Imϕ↑ . (14.103)
⋄
Exercise 14.9 Translate a quaternion into a Pauli spinor. Show that if the rotor R corresponds
to a right-handed rotation by angle θ about unit axis n ≡ n1, n2, n3, then the corresponding scaled Pauli
spinor ϕ ≡ λR ↑ is, from (14.83),
ϕ ≡ λR ↑ = λ
(
cos θ2 − in3 sin θ
2
(n2 − in1) sin θ2
)
. (14.104)
⋄
14.15 Spacetime algebra
So far this chapter has concerned itself with ordinary n-dimensional Euclidean space, in which the length
squared of a vector is the sum of the squares of its components. In special relativity, however, the scalar
220 ∗The geometric algebra
length s of a spacetime interval t, x, y, x is given by s2 = −t2 + x2 + y2 + z2. Happily, all the results of
previous sections hold with scarcely a change of stride.
Let γγm (m = 0, 1, 2, 3) denote an orthonormal basis of spacetime, with γγ0 representing the time axis, and
γγi (i = 1, 2, 3) the spatial axes. Geometric multiplication in the spacetime algebra is defined by
γγmγγn = γγm · γγn + γγm ∧γγn (14.105)
in the usual way. The key difference between the spacetime basis γγm and Euclidean bases is that scalar
products of the basis vectors γγm form the Minkowski metric ηmn,
γγm · γγn = ηmn (14.106)
whereas scalar products of Euclidean basis elements γγi formed the unit matrix, γγi ·γγj = δij , equation (14.6).
In less abbreviated form, equations (14.105) state that the geometric product of each basis element with
itself is
−γγ20 = γγ2
1 = γγ22 = γγ2
3 = 1 , (14.107)
while geometric products of different basis elements γγm anticommute
γγmγγn = −γγnγγm = γγm ∧γγn (m 6= n) . (14.108)
In the Dirac theory of relativistic spin- 12 particles, the Dirac γ-matrices are required to satisfy
γγm,γγn = 2 ηmn (14.109)
where denotes the anticommutator, γγm,γγn ≡ γγmγγn + γγnγγm. The multiplication rules (14.109) for
the Dirac γ-matrices are the same as those for geometric multiplication in the spacetime algebra, equa-
tions (14.107) and (14.108).
A 4-vector a, a multivector of grade 1 in the geometric algebra of spacetime, is
a = γγmam = γγ0a
0 + γγ1a1 + γγ2a
2 + γγ3a3 . (14.110)
Such a 4-vector a would be denoted 6a in the Dirac slash notation. The product of two 4-vectors a and b is
ab = a · b + a∧ b = ambnγγm · γγn + ambnγγm ∧γγn = ambnηmn + 12a
mbn[γγm,γγn] . (14.111)
It is convenient to denote three of the six bivectors of the spacetime algebra by σi,
σi ≡ γγ0γγi (i = 1, 2, 3) . (14.112)
The symbol σi is used because the algebra of bivectors σi is isomorphic to the algebra of Pauli matrices σi.
The triple of bivectors σi will often be denoted shorthandedly by the symbol σ
σ ≡ σ1,σ2,σ3 . (14.113)
The pseudoscalar, the highest grade basis element of the spacetime algebra, is denoted I
γγ0γγ1γγ2γγ3 = σ1σ2σ3 = I . (14.114)
14.16 Complex quaternions 221
The pseudoscalar I satisfies
I2 = −1 , Iγγm = −γγmI , Iσi = σiI . (14.115)
The basis elements of the 4-dimensional spacetime algebra are then
1 ,
1 scalar
γγm ,
4 vectors
σi , Iσi ,
6 bivectors
Iγγm ,
4 pseudovectors
I ,
1 pseudoscalar(14.116)
forming a linear space of dimension 1 + 4 + 6 + 4 + 1 = 16 = 24. The reverse is defined in the usual
way, equation (14.12), leaving unchanged multivectors of grade 0 or 1, modulo 4, and changing the sign of
multivectors of grade 2 or 3, modulo 4:
1 = 1 , γγm = γγm , σi = −σi , Iσi = −Iσi , Iγγm = −Iγγm , I = I . (14.117)
The mapping
γγ(3)i ↔ σi (i = 1, 2, 3) (14.118)
(the superscript (3) distinguishes the 3D basis vectors from the 4D spacetime basis vectors) defines an iso-
morphism between the 8-dimensional geometric algebra (14.3) of 3 spatial dimensions and the 8-dimensional
even spacetime subalgebra. Among other things, the isomorphism (14.118) implies the equivalence of the
3D spatial pseudoscalar i3 and the 4D spacetime pseudoscalar I
i3 ↔ I (14.119)
since i3 = γγ1γγ2γγ3 and I = σ1σ2σ3.
14.16 Complex quaternions
A complex quaternion (also called a biquaternion by W. R. Hamilton) is a quaternion
q = a+ ıibi = a+ ı · b (14.120)
in which the four coefficients a, bi (i = 1, 2, 3) are each complex numbers
a = aR + IaI , bi = bi,R + Ibi,I . (14.121)
The imaginary I is taken to commute with each of the quaternionic imaginaries ıi. The choice of symbol
I is deliberate: in the isomorphism (14.133) between the even spacetime algebra and complex quaternions,
the commuting imaginary I is isomorphic to the spacetime pseudoscalar I.
All of the equations in §14.10 on real quaternions remain valid without change, including the multiplication,
conjugation, and inversion formulae (14.57)–(14.62). In the quaternionic conjugate q of a complex quaternion
q ≡ a+ ı · b,
q = a− ı · b , (14.122)
222 ∗The geometric algebra
the complex coefficients a and b are not conjugated with respect to the complex imaginary I. The magnitude
|q| of a complex quaternion q ≡ a+ ı · b,
|q| = (qq)1/2 = (qq)1/2 = (a2 + b · b)1/2 = (a2 + b21 + b22 + b23)1/2 , (14.123)
is a complex number, not a real number. The complex conjugate q∗ of the complex quaternion is
q∗ = a∗ + ı · b∗ , (14.124)
in which the complex coefficients a and b are conjugated with respect to the imaginary I, but the quaternionic
imaginaries ı are not conjugated.
A non-zero complex quaternion can have zero magnitude (unlike a real quaternion), in which case it is
null. The null condition qq = a2 + b21 + b22 + b23 = 0 is a complex condition. The product of two null complex
quaternions is a null quaternion. Under multiplication, null quaternions form a 6-dimensional subsemigroup
(not a subgroup, because null quaternions do not have inverses) of the 8-dimensional semigroup of complex
quaternions.
Exercise 14.10 Show that any non-trivial null complex quaternion q can be written uniquely in the form
q = (1 + Iı · n)p , (14.125)
where p is a real quaternion, and n is a unit real 3-vector. Equivalently,
q = p(1 + Iı · n′) , (14.126)
where n′ is the unit real 3-vector
n′ =pnp
|p|2 . (14.127)
Solution. Write the null quaternion q as
q = p+ Ir (14.128)
where p and r are real quaternions, both of which must be non-zero if q is non-trivial. Then equation (14.125)
is true with
ı · n =rp
|p|2. (14.129)
The null condition is qq = 0. The vanishing of the real part, Re (qq) = pp− rr = 0, shows that |p|2 = |r|2.The vanishing of the imaginary (I) part, Im (qq) = rp+ pr = rp+ rp = 0 shows that the rp must be a pure
quaternionic imaginary, since the quaternionic conjugate of rp is minus itself, so rp/ |p|2 must be of the form
ı ·n. Its squared magnitude ı · n ı · n = rp pr/ |p|4 = |r|2/ |p|2 = 1 is unity, so n must be a unit 3-vector. It
follows immediately from the manner of construction that the expression (14.125) is unique, as long as q is
non-trivial. ⋄
14.17 Lorentz transformations and complex quaternions 223
14.17 Lorentz transformations and complex quaternions
Lorentz transformations are rotations of spacetime. Such rotations correspond, in the usual way, to even,
unimodular elements of the geometric algebra of spacetime. The basis elements of the even spacetime algebra
are
1 ,
1 scalar
σi , Iσi ,
6 bivectors
I ,
1 pseudoscalar(14.130)
forming a linear space of dimension 1 + 6 + 1 = 8 over the real numbers. However, it is more elegant to
treat the even spacetime algebra as a linear space of dimension 8 ÷ 2 = 4 over complex scalars of the form
λ = λR + IλI . The pseudoscalar I qualifies as a scalar because it commutes with all elements of the even
spacetime algebra, and it qualifies as an imaginary because I2 = −1. It is convenient to take the basis
elements of the even spacetime algebra over the complex numbers to be
1 ,
1 scalar
Iσi ,
3 bivectors(14.131)
forming a linear space of dimension 1+3 = 4. The reason for choosing Iσi rather than σi as the elements of
the basis (14.131) is that the basis is 1, Iσi is equivalent to the basis (14.63) of the even algebra of 3-dim-
ensional Euclidean space through the isomorphism (14.118) and (14.119). This basis in turn is equivalent to
the quaternionic basis 1, ıi through the isomorphism (14.68):
Iσi ↔ i3γγ(3)i ↔ ıi (i = 1, 2, 3) . (14.132)
In other words, the even spacetime algebra is isomorphic to the algebra of quaternions with complex coeffi-
cients:
a+ Iσ · b↔ a+ ı · b (14.133)
where a = aR + IaI is a complex number, b = bR + IbI , is a triple of complex numbers, σ is the triple of
bivectors σi, and ı is the triple of quaternionic imaginaries.
The isomorphism (14.133) between even elements of the spacetime algebra and complex quaternions implies
that the group of Lorentz rotors, which are unimodular elements of the even spacetime algebra, is isomorphic
to the group of unimodular complex quaternions
spacetime rotors↔ unit complex quaternions . (14.134)
In §14.11 it was found that the group of 3D spatial rotors is isomorphic to the group of unimodular real
quaternions. Thus Lorentz transformations are mathematically equivalent to complexified spatial rotations.
A Lorentz rotor can be written as a complex quaternion in what looks like the same form as the expres-
sion (14.70) for a 3D spatial rotor, with the difference that the rotation angle θ is complex, and the axis n
of rotation is likewise complex. Thus
R = eθ/2 = eı·n θ/2 = cosθ
2+ ı · n sin
θ
2(14.135)
224 ∗The geometric algebra
where θ is the bivector complex quaternion
θ ≡ ı · n θ ≡ (ı1n1 + ı2n2 + ı3n3) θ (14.136)
whose complex magnitude is |θ| ≡ (θθ)1/2 = θ and whose complex unit direction is θ ≡ θ/θ ≡ ı·n. The angle
θ = θR + IθI is a complex angle, and n = nR + InI is a complex-valued unit 3-vector, satisfying n · n = 1.
The condition n ·n = 1 of unit normalization is equivalent to the two conditions nR ·nR −nI ·nI = 1 and
2nR ·nI = 0 on the real and imaginary parts of n ·n. The complex angle θ has 2 degrees of freedom, while
the complex unit vector has 4 degrees of freedom, so the Lorentz rotor R has 6 degrees of freedom, which
is the correct number of degrees of freedom of the group of Lorentz transformations. The associated reverse
rotor R is
R = e−θ/2 = e−ı·n θ/2 = cosθ
2− ı · n sin
θ
2(14.137)
the quaternionic conjugate of R. Note that θ and n in equation (14.137) are not conjugated with respect to
the imaginary I. The sine and cosine of the complex angle θ appearing in equations (14.135) and (14.137)
are related to its real and imaginary parts in the usual way,
cosθ
2= cos
θR
2cosh
θI
2− I sin
θR
2sinh
θI
2, sin
θ
2= sin
θR
2cosh
θI
2+ I cos
θR
2sinh
θI
2. (14.138)
In the case of a pure spatial rotation, the angle θ = θR and axis n = nR in the rotor (14.135) are both
real. The rotor corresponding to a pure spatial rotation by angle θR right-handedly about unit real axis nR
is
R = eı·nR θR/2 = cosθR
2+ ı · nR sin
θR
2. (14.139)
A Lorentz boost is a change of velocity in some direction, without any spatial rotation, and represents
a rotation of spacetime about some time-space plane. For example, a Lorentz boost along the 1-axis (the
x-axis) is a rotation of spacetime in the 0–1 plane (the t–x plane). In the case of a pure Lorentz boost, the
angle θ = IθI is pure imaginary, but the axis n = nR remains pure real. The rotor corresponding to a boost
by velocity v = tanh θI in unit real direction nR is
R = eı·nR IθI/2 = coshθI
2+ ı · nR I sinh
θI
2. (14.140)
Exercise 14.11 Factor a general Lorentz rotor R = eı·n θ/2 into the product UL of a pure spatial rotation
U followed by a pure Lorentz boost L. Do the two factors commute? ⋄Exercise 14.12 Show that the geometry of the group of Lorentz rotors is the product of the geometries
of the spatial rotation group and the boost group, which is a 3-sphere times Euclidean 3-space, S3 × R3. ⋄
14.18 Spatial Inversion (P ) and Time Inversion (T )
Spatial inversion, or P for parity, is the operation of reflecting all spatial coordinates while keeping the time
coordinate unchanged. Spatial inversion may be accomplished by reflecting the spatial vector basis elements
14.19 Electromagnetic field bivector 225
γγi → −γγi, while keeping the time vector basis element γγ0 unchanged. This results in σ → −σ and I → −I.The equivalence Iσ ↔ ı means that the quaternionic imaginary ı is unchanged. Thus, if multivectors in
the geometric spacetime algebra are written as linear combinations of products of γγ0, ı, and I, then spatial
inversion P corresponds to the transformation
P : γγ0 → γγ0 , ı→ ı , I → −I . (14.141)
In other words spatial inversion may be accomplished by the rule, take the complex conjugate (with respect
to I) of a multivector.
Time inversion, or T , is the operation of reversing time while keeping all spatial coordinates unchanged.
Time inversion may be accomplished by reflecting the time vector basis element γγ0 → −γγ0, while keeping the
spatial vector basis elements γγi unchanged. As with spatial inversion, this results in σ → −σ and I → −I,which keeps Iσ hence ı unchanged. If multivectors in the geometric spacetime algebra are written as linear
combinations of products of γγ0, ı, and I, then time inversion T corresponds to the transformation
T : γγ0 → −γγ0 , ı→ ı , I → −I . (14.142)
For any multivector, time inversion corresponds to the instruction to flip γγ0 and take the complex conjugate
(with respect to I).
The combined operation PT of inverting both space and time corresponds to
PT : γγ0 → −γγ0 , ı→ ı , I → I . (14.143)
For any multivector, spacetime inversion corresponds to the instruction to flip γγ0, while keeping ı and I
unchanged.
14.19 Electromagnetic field bivector
The electromagnetic field tensor Fmn can be expressed as the bivector
F = 12F
mnγγm ∧γγn , (14.144)
the factor of 12 compensating for the double-counting over indices m and n (the 1
2 could be omitted if the
counting were over distinct bivector indices only). In terms of the electric and magnetic fields E and B, and
the bivector basis elements σ ≡ σ1,σ2,σ3 defined by equation (14.112), the electromagnetic field bivector
F is
F = −σ · (E + IB) . (14.145)
14.20 How to implement Lorentz transformations on a computer
The advantages of quaternions for implementing spatial rotations are well-known to 3D game programmers.
Compared to standard rotation matrices, quaternions offer increased speed and require less storage, and
226 ∗The geometric algebra
their algebraic properties simplify interpolation and splining. Complex quaternions retain similar advantages
for implementing Lorentz transformations. They are fast, compact, and straightforward to interpolate or
spline (Exercises 14.13 and 14.14). Moreover, since complex quaternions contain real quaternions, Lorentz
transformations can be implemented simply as an extension of spatial rotations in 3D programs that use
quaternions to implement spatial rotations.
Lorentz rotors, 4-vectors, spacetime bivectors, and spinors (spin- 12 objects) can all be implemented as
complex quaternions. A complex quaternion
q = w + ı1x+ ı2y + ı3z (14.146)
with complex coefficients w, x, y, z can be stored as the 8-component object
q =
wR xR yR zR
wI xI yI zI
. (14.147)
Actually, OpenGL and other computer software store the scalar (w) component of a quaternion in the last
(fourth) place, but here the scalar components are put in the zeroth position to conform to standard physics
convention. The quaternion conjugate q of the quaternion (14.147) is
q =
wR −xR −yR −zR
wI −xI −yI −zI
, (14.148)
while its complex conjugate q∗ is
q∗ =
wR xR yR zR
−wI −xI −yI −zI
. (14.149)
A Lorentz rotor R corresponds to a complex quaternion of unit modulus. The unimodular condition RR =
1, a complex condition, removes 2 degrees of freedom from the 8 degrees of freedom of complex quaternions,
leaving the Lorentz group with 6 degrees of freedom, which is as it should be. Spatial rotations correspond
to real unimodular quaternions, and account for 3 of the 6 degrees of freedom of Lorentz transformations.
A spatial rotation by angle θ right-handedly about the 1-axis (the x-axis) is the real Lorentz rotor
R = cos(θ/2) + ı1 sin(θ/2) , (14.150)
or, stored as a complex quaternion,
R =
cos(θ/2) sin(θ/2) 0 0
0 0 0 0
. (14.151)
Lorentz boosts account for the remaining 3 of the 6 degrees of freedom of Lorentz transformations. A Lorentz
boost by velocity v, or equivalently by boost angle θ = atanh(v), along the 1-axis (the x-axis) is the complex
Lorentz rotor
R = cosh(θ/2) + Iı1 sinh(θ/2) , (14.152)
14.20 How to implement Lorentz transformations on a computer 227
or, stored as a complex quaternion,
R =
cosh(θ/2) 0 0 0
0 sinh(θ/2) 0 0
. (14.153)
The rule for composing Lorentz transformations is simple: a Lorentz transformation R followed by a Lorentz
transformation S is just the product RS of the corresponding complex quaternions.
The inverse of a Lorentz rotor R is its quaternionic conjugate R.
Any even multivector q is equivalent to a complex quaternion by the isomorphism (14.133). According to
the usual transformation law (14.35) for multivectors, the rule for Lorentz transforming an even multivector
q is
R : q → RqR (even multivector) . (14.154)
The transformation (14.154) instructs to multiply three complex quaternions R, q, and R, a one-line expres-
sion in a c++ program.
As an example of an even multivector, the electromagnetic field F , equation (14.145), is a bivector in the
spacetime algebra. In view of the isomorphism (14.133), the electromagnetic field bivector F can be written
as the complex quaternion
F =
0 −B1 −B2 −B3
0 E1 E2 E3
. (14.155)
Under the parity transformation P (14.141), the electric field E changes sign, whereas the magnetic field B
does not, which is as it should be:
P : E → −E , B → B . (14.156)
According to the rule (14.154), the electromagnetic field bivector F Lorentz transforms as F → RFR, which
is a powerful and elegant way to Lorentz transform the electromagnetic field.
A 4-vector a ≡ γγmam is a multivector of grade 1 in the spacetime algebra. A general odd multivector in
the spacetime algebra is the sum of a vector (grade 1) part a and a pseudovector (grade 3) part Ib = Iγγmbm.
The odd multivector can be written as the product of the time basis vector γγ0 and an even multivector q
a + Ib = γγ0q = γγ0
(
a0 + Iıiai − Ib0 + ıib
i)
. (14.157)
By the isomorphism (14.133), the even multivector q is equivalent to the complex quaternion
q =
a0 b1 b2 b3
−b0 a1 a2 a3
. (14.158)
According to the usual transformation law (14.35) for multivectors, the rule for Lorentz transforming the
odd multivector γγ0q is
R : γγ0q → Rγγ0qR = γγ0R∗qR . (14.159)
In the last expression of (14.159), the factor γγ0 has been brought to the left, to be consistent with the
convention (14.157) that an odd multivector is γγ0 on the left times an even multivector on the right. Notice
228 ∗The geometric algebra
that commuting γγ0 through R converts the latter to its complex conjugate (with respect to I) R∗, which
is true because γγ0 commutes with the quaternionic imaginary ı, but anticommutes with the pseudoscalar
I. Thus if the components of an odd multivector are stored as a complex quaternion (14.158), then that
complex quaternion q Lorentz transforms as
R : q → R∗qR (odd multivector) . (14.160)
The rule (14.160) again instructs to multiply three complex quaternions R∗, q, and R, a one-line expression in
a c++ program. The transformation rule (14.160) for an odd multivector encoded as a complex quaternion
differs from that (14.154) for an even multivector in that the first factor R is complex conjugated (with
respect to I).
A vector a differs from a pseudovector Ib in that the vector a changes sign under a parity transformation
P whereas the pseudovector Ib does not. However, the behaviour of a pseudovector under a normal Lorentz
transformation (which preserves parity) is identical to that of a vector. Thus in practical situations two
4-vectors a and b can be encoded into a single complex quaternion (14.158), and Lorentz transformed
simultaneously, enabling two transformations to be done for the price of one.
Finally, a Dirac spinor is equivalent to a complex quaternion q (§14.23). It Lorentz transforms as
R : q → Rq (spinor) . (14.161)
Exercise 14.13 Interpolate a Lorentz transformation. Argue that the interpolating Lorentz rotor
R(x) that corresponds to uniform rotation and acceleration between initial and final Lorentz rotors R0 and
R1 as the parameter x varies uniformly from 0 to 1 is
R(x) = R0 exp [x ln(R1/R0)] . (14.162)
What are the exponential and logarithm of a complex quaternion in terms of its components? Address the
issue of the multi-valued character of the logarithm. ⋄
Exercise 14.14 Spline a Lorentz transformation. A spline is a polynomial that interpolates between
two points with given values and derivatives at the two points. Confirm that the cubic spline of a real
function f(x) with given initial and final values f0 and f1 and given initial and final derivatives f ′0 and f ′
1
at x = 0 and x = 1 is
f(x) = f0 + f ′0x+ [3(f1 − f0)− 2f ′
0 − f ′1]x
2 + [2(f0 − f1) + f ′0 + f ′
1]x3 . (14.163)
The case in which the derivatives at the endpoints are set to zero, f ′0 = f ′
1 = 0, is called the “natural” spline.
Argue that a Lorentz rotor can be splined by splining the quaternionic components of the logarithm of the
Lorentz rotor. ⋄
Exercise 14.15 The wrong way to implement a Lorentz transformation. The purpose of this
exercise is to persuade you that Lorentz transforming a 4-vector by the rule (14.160) is a much better idea
than Lorentz transforming by multiplying by an explicit 4× 4 matrix. Suppose that the Lorentz rotor R is
14.21 Dirac matrices 229
the complex quaternion
R =
wR xR yR zR
wI xI yI zI
. (14.164)
Show that the Lorentz transformation (14.160) transforms the 4-vector components am = a0, a1, a2, a3 as
(note that the 4× 4 rotation matrix is written to the right of the 4-vector in accordance with the computer
graphics convention that rotations accumulate to the right — opposite to the physics convention; to recover
the physics convention, take the transpose):
R :(
a0 a1 a2 a3)
→(
a0 a1 a2 a3)
|w|2 + |x|2 + |y|2 + |z|2 2 (wRxI − wIxR + yRzI − yIzR)
2 (wRxI − wIxR − yRzI + yIzR) |w|2 + |x|2 − |y|2 − |z|22 (wRyI − wIyR − zRxI + zIxR) 2 (xRyR + xIyI − wRzR − wIzI)
2 (wRzI − wIzR − xRyI + xIyR) 2 (zRxR + zIxI + wRyR + wIyI)
2 (wRyI − wIyR + zRxI − zIxR) 2 (wRzI − wIzR + xRyI − xIyR)
2 (xRyR + xIyI + wRzR + wIzI) 2 (zRxR + zIxI − wRyR − wIyI)
|w|2 − |x|2 + |y|2 − |z|2 2 (yRzR + yIzI + wRxR + wIxI)
2 (yRzR + yIzI − wRxR − wIxI) |w|2 − |x|2 − |y|2 + |z|2
, (14.165)
where | | signifies the absolute value of a complex number, as in |w|2 = w2R +w2
I . As a simple example, show
that the transformation (14.165) in the case of a Lorentz boost by velocity v along the 1-axis, where the
rotor R takes the form (14.153), is
R :(
a0 a1 a2 a3)
→(
a0 a1 a2 a3)
γ γv 0 0
γv γ 0 0
0 0 1 0
0 0 0 1
, (14.166)
with γ the familiar Lorentz gamma factor
γ = cosh θ =1
(1 − v2)1/2, γv = sinh θ =
v
(1− v2)1/2. (14.167)
⋄
14.21 Dirac matrices
The multiplication rules (14.105) for the basis vectors γγm of the spacetime algebra are identical to the
rules (14.109) governing the Clifford algebra of the Dirac γ-matrices used in the Dirac theory of relativistic
spin- 12 particles.
The Dirac γ-matrices are conventionally represented by 4 × 4 complex matrices. To ensure consistency
230 ∗The geometric algebra
between the relativistic and quantum mechanical ways of taking the scalar (inner) product, it is desirable to
require that taking the Hermitian conjugate of any of the basis vectors γγm be equivalent to raising its index,
γγ†m = γγm . (14.168)
Given the requirement (14.168), the matrices representing the basis vectors γγm must be traceless (because
a trace is a scalar, and the basis vectors cannot contain any scalar part), Hermitian or anti-Hermitian as
the self-product of the matrix is ±1 (so that γγ†m = γγm = ηmnγγn), and unitary and anticommuting (so that
γγ†m · γγn = γγm · γγn = δm
n ). The precise choice of matrices is not fundamental: any set of 4 matrices satisfying
these conditions will do.
The high-energy physics community conventionally adopts the +−−− metric signature, which is opposite
to the convention adopted here. With the high-energy +−−− signature, the standard convention for the
Dirac γ-matrices is
γγ0 =
(
1 0
0 −1
)
, γγi =
(
0 σi
−σi 0
)
, (14.169)
where 1 denotes the unit 2 × 2 matrix, and σi denote the three 2× 2 Pauli matrices (14.76). The choice of
γγ0 as a diagonal matrix is motivated by Dirac’s discovery that eigenvectors of the time basis vector γγ0 with
eigenvalues of opposite sign define particles and antiparticles in their rest frames (see §14.22). To convert to
the −+++ metric signature adopted here while retaining the conventional set of eigenvectors, an additional
factor of i must be inserted into the γ-matrices:
γγ0 = i
(
1 0
0 −1
)
, γγi = i
(
0 σi
−σi 0
)
, (14.170)
In the representation of equations (14.169) or (14.170), the bivectors σi and Iσi and the pseudoscalar I of
the spacetime algebra are
σi =
(
0 −σi
−σi 0
)
, Iσi = i
(
σi 0
0 σi
)
, I = −i(
0 1
1 0
)
, (14.171)
whose representation as matrices is the same for either signature−+++ or +−−−. The Hermitian conjugates
of the bivector and pseudoscalar basis elements are
σ†i = σi , (Iσi)
† = −Iσi , I† = −I . (14.172)
The conventional chiral matrix γ5 of Dirac theory is defined by
γ5 ≡ iγγ0γγ1γγ2γγ3 = iI =
(
0 1
1 0
)
, (14.173)
whose representation is again the same for either signature −+++ or +−−−. The chiral matrix is Hermitian
γ†5 = γ5 . (14.174)
14.22 Dirac spinors 231
14.22 Dirac spinors
In the Dirac theory of relativistic spin- 12 particles, a Dirac spinor ϕ is represented as a 2-component column
vector of Pauli spinors ϕ⇑ and ϕ⇓, comprising 4 complex (with respect to i) components and hence 8 degrees
of freedom,
ϕ =
(
ϕ⇑
ϕ⇓
)
=
ϕ⇑↑
ϕ⇑↓
ϕ⇓↑
ϕ⇓↓
. (14.175)
The Dirac γ-matrices operate by pre-multiplication on Dirac spinors ϕ, yielding other Dirac spinors.
In the Dirac representation (14.170), the four unit Dirac spinors
⇑↑ =
1
0
0
0
, ⇑↓ =
0
1
0
0
, ⇓↑ =
0
0
1
0
, ⇓↓ =
0
0
0
1
, (14.176)
are eigenvectors of the time basis vector γγ0 and of the bivector Iσ3, with ⇑ and ⇓ denoting eigenvectors of
γγ0, and ↑ and ↓ eigenvectors of Iσ3,
γγ0 ⇑ = i⇑ , γγ0 ⇓ = −i⇓ , Iσ3 ↑ = i ↑ , Iσ3 ↓ = −i ↓ . (14.177)
The bivector Iσ3 is the generator of a spatial rotation about the 3-axis (z-axis), equation (14.132). The four
eigenvectors (14.177) form an orthonormal basis
A pure spin-up state ↑ can be rotated into a pure spin-down state ↓, or vice versa, by a spatial rotation
about the 1-axis or 2-axis. By contrast, a pure time-up state ⇑ cannot be rotated into a pure time-down
state ⇓, or vice versa, by any Lorentz transformation. Consider for example trying to rotate the pure time-up
spin-up ⇑↑ state into any combination of pure time-down ⇓ states. According to the expression (14.192), the
Dirac spinor ϕ obtained by Lorentz transforming the ⇑↑ state is pure ⇓ only if the corresponding complex
quaternion q is pure imaginary. But a pure imaginary quaternion has negative squared magnitude qq, so
cannot be equivalent to any rotor of unit magnitude.
Thus the pure time-up and pure time-down states ⇑ and ⇓ are distinct states that cannot be transformed
into each other by any Lorentz transformation. The two states represent distinct species, particles and
antiparticles.
Although a pure time-up state cannot be transformed into a pure time-down state or vice versa by any
Lorentz transformation, the time-up and time-down eigenstates ⇑ and ⇓ do mix under Lorentz transforma-
tions. The manner in which Dirac spinors transform is described in §14.23.
The choice of time-axis γγ0 and spin-axis γγ3 with respect to which the eigenvectors are defined can of course
be adjusted arbitrarily by a Lorentz boost and a spatial rotation. The eigenvectors of a particular time-axis
γγ0 correspond to particles and antiparticles that are at rest in that frame. The eigenvectors associated with
a particular spin-axis γγ3 correspond to particles or antiparticles that are pure spin-up or pure spin-down in
that frame.
232 ∗The geometric algebra
14.23 Dirac spinors as complex quaternions
In §14.14 it was found that a spin- 12 object in 3D space, a Pauli spinor, is isomorphic to a scaled 3D reverse
rotor, or real quaternion. In the relativistic theory, the corresponding spin- 12 object, a Dirac spinor ϕ, is
isomorphic (14.181) to a complex quaternion. The 4 complex degrees of freedom of the Dirac spinor ϕ are
equivalent to the 8 degrees of freedom of a complex quaternion. A physically interesting complication arises
in the relativistic case because a non-trivial Dirac spinor can be null, with zero magnitude, whereas any non-
trivial Pauli spinor is necessarily non-null. The case of non-null (massive) and null (massless) Dirac spinors
are considered respectively in §14.24 and §14.25. The present section establishes an isomorphism (14.181)
between Dirac spinors and complex quaternions that is valid in general, regardless of whether the Dirac
spinor is null or not.
If a is a spacetime multivector, equivalent to an element of the Clifford algebra of Dirac γ-matrices, then
under rotation by Lorentz rotor R, the multivector a operating on the Dirac spinor ϕ transforms as
R : aϕ→ (RaR)(Rϕ) = Raϕ . (14.178)
This shows that a Dirac spinor ϕ Lorentz transforms, by construction, as
R : ϕ→ Rϕ . (14.179)
The rule (14.179) is precisely the transformation rule for reverse spacetime rotors under Lorentz trans-
formations: under a rotation by rotor R, a reverse rotor S transforms as S → RS. More generally, the
transformation law (14.179) holds for any linear combination of Dirac spinors ϕ. The isomorphism (14.134)
between spacetime rotors and unit quaternions shows that unit Dirac spinors are isomorphic to unit (reverse)
complex quaternions. The algebra of linear combinations of unit complex quaternions is just the algebra of
complex quaternions. Thus the algebra of Dirac spinors is isomorphic to the algebra of (reverse) complex
quaternions. Specifically, any Dirac spinor ϕ can be expressed uniquely in the form of a 4× 4 matrix q, the
Dirac representation of a reverse complex quaternion q, acting on the time-up spin-up eigenvector ⇑↑ (the
precise translation between Dirac spinors and complex quaternions is left as Exercises 14.16 and 14.17):
ϕ = q ⇑↑ . (14.180)
In this section (including the Exercises) the 4× 4 matrix q is written in boldface to distinguish it from the
quaternion q that it represents; but the distinction is not fundamental, so the temporary boldface notation
is dropped in subsequent sections. The equivalence (14.180) establishes that Dirac spinors are isomorphic to
reverse complex quaternions
ϕ↔ q . (14.181)
The isomorphism means that there is a one-to-one correspondence between Dirac spinors ϕ and reverse
complex quaternions q, and that they transform in the same way under Lorentz transformations.
Notwithstanding the isomorphism (14.181), Dirac spinors differ from complex quaternions in that they
have an additional structure that is essential to quantum mechanics, an inner product ϕ†1ϕ2 of two Dirac
14.23 Dirac spinors as complex quaternions 233
spinors ϕ1 and ϕ2. The inner product ϕ†1ϕ2 is a complex (with respect to i) number. The Hermitian
conjugate ϕ† of a Dirac spinor ϕ (14.180) is defined to be
ϕ† = (⇑↑)†q† , (14.182)
where (⇑↑)† =(
1 0 0 0)
is the Hermitian conjugate of the time-up spin-up eigenvector ⇑↑, and q† is
the Hermitian conjugate of the matrix q. A related spinor is the the reverse, or adjoint, spinor ϕ, defined to
be
ϕ = (⇑↑)†q , (14.183)
where q is the reverse of the matrix q. The Hermitian conjugate spinor ϕ† is related to the adjoint spinor
ϕ by (Exercise 14.18)
ϕ† = iϕγγ0 . (14.184)
The product ϕϕ of the adjoint spinor ϕ with ϕ is a Lorentz-invariant scalar, as follows from the Lorentz
invariance of qq. On the other hand, the product ϕ†ϕ of the Hermitian conjugate spinor ϕ† with ϕ is the
time component of a 4-vector
iϕγγmϕ . (14.185)
Exercise 14.16 Translate a Dirac spinor into a complex quaternion. Given any Dirac spinor
ϕ =
ϕ⇑↑
ϕ⇑↓
ϕ⇓↑
ϕ⇓↓
, (14.186)
show that the corresponding reverse complex quaternion q, and the equivalent 4× 4 matrix q in the Dirac
representation (14.170), such that ϕ = q⇑↑, are (the complex conjugates ϕ∗a of the components ϕa of the
spinor are with respect to the quantum mechanical imaginary i)
q =
Reϕ⇑↑ Imϕ⇑↓ −Reϕ⇑↓ Imϕ⇑↑
−Imϕ⇓↑ Reϕ⇓↓ Imϕ⇓↓ Reϕ⇓↑
↔ q =
ϕ⇑↑ −ϕ∗⇑↓ ϕ⇓↑ ϕ∗
⇓↓
ϕ⇑↓ ϕ∗⇑↑ ϕ⇓↓ −ϕ∗
⇓↑
ϕ⇓↑ ϕ∗⇓↓ ϕ⇑↑ −ϕ∗
⇑↓
ϕ⇓↓ −ϕ∗⇓↑ ϕ⇑↓ ϕ∗
⇑↑
. (14.187)
Show that the complex quaternion q (the reverse of q), and the equivalent 4× 4 matrix q (the reverse of q)
in the Dirac representation (14.170), are
q =
Reϕ⇑↑ −Imϕ⇑↓ Reϕ⇑↓ −Imϕ⇑↑
−Imϕ⇓↑ −Reϕ⇓↓ −Imϕ⇓↓ −Reϕ⇓↑
↔ q =
ϕ∗⇑↑ ϕ∗
⇑↓ −ϕ∗⇓↑ −ϕ∗
⇓↓
−ϕ⇑↓ ϕ⇑↑ −ϕ⇓↓ ϕ⇓↑
−ϕ∗⇓↑ −ϕ∗
⇓↓ ϕ∗⇑↑ ϕ∗
⇑↓
−ϕ⇓↓ ϕ⇓↑ −ϕ⇑↓ ϕ⇑↑
. (14.188)
Conclude that the reverse spinor ϕ ≡ (⇑↑)†q is
ϕ ≡ (⇑↑)†q =(
ϕ∗⇑↑ ϕ∗
⇑↓ −ϕ∗⇓↑ −ϕ∗
⇓↓
)
. (14.189)
234 ∗The geometric algebra
⋄
Exercise 14.17 Translate a complex quaternion into a Dirac spinor. Show that the complex
quaternion q ≡ w + ıx+ y + kz is equivalent in the Dirac representation (14.170) to the 4× 4 matrix q
q =
wR xR yR zR
wI xI yI zI
↔ q =
wR + izR ixR + yR −iwI + zI xI − iyI
ixR − yR wR − izR xI + iyI −iwI − zI
−iwI + zI xI − iyI wR + izR ixR + yR
xI + iyI −iwI − zI ixR − yR wR − izR
. (14.190)
Show that the reverse quaternion q, the complex conjugate (with respect to I) quaternion q∗, and the reverse
complex conjugate (with respect to I) quaternion q∗ are respectively equivalent to the 4× 4 matrices
q ↔ q ≡ −γγ0q†γγ0 , (14.191a)
q∗ ↔ q† = −γγ0qγγ0 , (14.191b)
q∗ ↔ q† = −γγ0qγγ0 . (14.191c)
Conclude that the Dirac spinor ϕ ≡ q ⇑↑ corresponding to the reverse complex quaternion q is
ϕ ≡ q ⇑↑ =
wR − izR
−ixR + yR
−iwI − zI
−xI − iyI
, (14.192)
that the reverse spinor ϕ ≡ (⇑↑)†q is
ϕ ≡ (⇑↑)†q =(
wR + izR ixR + yR −iwI + zI xI − iyI
)
, (14.193)
and that the Hermitian conjugate spinor ϕ† ≡ (⇑↑)†q† is
ϕ† ≡ (⇑↑)†q† =(
wR + izR ixR + yR iwI − zI −xI + iyI
)
. (14.194)
Hence conclude that ϕϕ and ϕ†ϕ are respectively the real part of, and the absolute value of, the complex
magnitude squared qq ≡ λ2 of the complex quaternion q,
ϕϕ = λ2R − λ2
I , (14.195a)
ϕ†ϕ = λ2R + λ2
I , (14.195b)
with
λ2R = w2
R + x2R + y2
R + z2R , (14.196a)
λ2I = w2
I + x2I + y2
I + z2I , (14.196b)
⋄
14.24 Non-null Dirac spinor — particle and antiparticle 235
Exercise 14.18 Relation between ϕ† and ϕ. Confirm equation (14.184) by showing from equa-
tion (14.191b) that
ϕ† = i(⇑↑)†qγγ0 . (14.197)
⋄
14.24 Non-null Dirac spinor — particle and antiparticle
A non-null, or massive, Dirac spinor ϕ, one for which ϕϕ 6= 0, is isomorphic (14.181) to a non-null reverse
complex quaternion q, which can be factored as a non-zero complex scalar times a unit quaternion, a rotor.
Thus a non-null Dirac spinor can be expressed as the product of a complex scalar λ = λR +IλI and a reverse
Lorentz rotor R, acting on the time-up spin-up eigenvector ⇑↑,
ϕ = λR⇑↑ . (14.198)
The complex scalar λ can be taken without loss of generality to lie in the right hemisphere of the complex
plane (positive real part), since a minus sign can be absorbed into a spatial rotation by 2π of the rotor
R. There is no further ambiguity in the decomposition (14.198) into scalar and rotor, because the squared
magnitude λRλR = λ2 of the scaled rotor λR is the same for any decomposition.
The fact that a non-null Dirac spinor ϕ encodes a Lorentz rotor shows that a non-null Dirac spinor in
some sense “knows” about the Lorentz structure of spacetime. It is intriguing that the Lorentz structure of
spacetime is built in to a non-null Dirac particle.
As discussed in §14.22, a pure time-up eigenvector ⇑ represents a particle in its own rest frame, while a pure
time-down eigenvector ⇓ represents an antiparticle in its own rest frame. The time-up spin-up eigenvector ⇑↑is by definition (14.198) equivalent to the unit scaled rotor, λR = 1, so in this case the scalar λ is pure real.
Lorentz transforming the eigenvector multiplies it by a rotor, but leaves the scalar λ unchanged, therefore
pure real. Conversely, if the time-up spin-up eigenvector ⇑↑ is multiplied by the imaginary I, then according
to the expression (14.192) the resulting spinor can be Lorentz transformed into a pure ⇓ state, corresponding
to a pure antiparticle. Thus one may conclude that the real and imaginary parts (with respect to I)
of the complex scalar λ = λR + IλI correspond respectively to particles and antiparticles. The
magnitude squared ϕ†ϕ of the Dirac spinor, equation (14.195b),
ϕ†ϕ = |λ|2 = λ2R + λ2
I , (14.199)
is the sum of the probabilities λ2R of particles and λ2
I of antiparticles.
Among other things, the decomposition of a Dirac spinor into its particle and antiparticle parts shows
that multiplying a non-null Dirac spinor by the pseudoscalar I converts a particle to an antiparticle, and
vice versa.
236 ∗The geometric algebra
14.25 Null Dirac Spinor
A null spinor is a spinor ϕ whose magnitude is zero,
ϕϕ = 0 . (14.200)
Such a spinor is equal to a null complex quaternion q acting on the time-up spin-up eigenvector ⇑↑,
ϕ = q ⇑↑ . (14.201)
Physically, a null spinor represents a spin- 12 particle moving at the speed of light. A non-trivial null spinor
must be moving at the speed of light because if it were not, then there would be a rest frame where the rotor
part of the spinor ϕ = λR⇑↑ would be unity, R = 1, and the spinor, being non-trivial, λ 6= 0, would not be
null. The null condition (14.200) is a complex constraint, which eliminates 2 of the 8 degrees of freedom of
a complex quaternion, so that a null spinor has 6 degrees of freedom.
Any non-trivial null complex quaternion q can be written uniquely as the product of a null factor
(1 + Iı ·n)/√
2 and a real quaternion λU (Exercise 14.10):
q =(1 + Iı · n)√
2λU . (14.202)
Here n is a unit real 3-vector, λ is a positive real scalar, and U is a purely spatial (i.e. real, with no I part)
rotor. The factor of 1/√
2 is inserted for normalization purposes. Physically, equation (14.202) contains the
instruction to boost to light speed in the direction n, then scale by the real scalar λ and rotate spatially by
U . The 2 + 1 + 3 = 6 degrees of freedom from the real unit vector n, the real scalar λ, and the spatial rotor
U in the expression (14.202) are precisely the number needed to specify a null quaternion. One might have
thought that the boost factor 1 + Iı · n in equation (14.202) would change under a Lorentz transformation,
but in fact it is Lorentz-invariant. For if the boost factor 1+Iı ·n is transformed (multiplied) by any complex
quaternion p+ Ir, then the result
(1 + Iı · n)(p+ Ir) = (1 + Iı · n)(p− ı · n r) (14.203)
is the same unchanged boost factor 1+Iı·n multiplied by a purely spatial transformation, the real quaternion
p−ı·n r. Equation (14.203) is true because (ı·n)2 = −1. Since the boost factor 1+Iı·n is Lorentz-invariant,
Lorentz transforming the null quaternion q (14.202) probes only 4 of the 6 degrees of freedom of the group
of null quaternions.
The null Dirac spinor ϕ corresponding to the reverse q of the null complex quaternion q, equation (14.202),
is
ϕ ≡ q ⇑↑ = λU(1 − Iı · n)√
2⇑↑ . (14.204)
It is natural to choose basis vectors of the representation to be eigenvectors of the Lorentz-invariant boost
factor 1− Iı ·n. In the Dirac representation (14.170), the basis spinors (14.176) are eigenvectors of Iσ3, and
it natural to choose the 3-direction to be in either the positive or negative n direction, in which case
(1− Iı · n)⇑↑ = (1∓ Iı3)⇑↑ = (1± σ3)⇑↑ = (⇑ ∓ ⇓) ↑ . (14.205)
14.26 Chiral decomposition of a Dirac spinor 237
The null basis vectors are left- and right-handed chiral eigenvectors, eigenvectors of the chiral operator γ5
with eigenvalues ∓ respectively,
γ5(⇑∓⇓) ↑√
2= ∓ (⇑∓⇓) ↑√
2. (14.206)
The 1/√
2 factor ensures unit normalization
[
(⇑∓⇓) ↑√2
]†(⇑∓⇓) ↑√
2= 1 . (14.207)
A general null Dirac spinor ϕ, equation (14.204), is the real scaled spatial reverse rotor λU acting on one of
the two null chiral basis spinors, either left-handed (⇑−⇓) ↑/√
2, or right-handed (⇑+⇓) ↑/√
2,
ϕL = λU(⇑−⇓) ↑√
2, ϕR = λU
(⇑+⇓) ↑√2
. (14.208)
A left- or right-handed null Dirac spinor is called a Weyl spinor.
Concept question 14.19 The null boost factor (1 + Iı · n) in a null quaternion, equation (14.202), is
Lorentz-invariant, as shown by equation (14.203) (which you should confirm). Consequently a null Dirac
spinor has a Lorentz-invariant boost axis n. Does a null 4-vector have a Lorentz-invariant axis? What does
it mean physically that a null Dirac spinor has a Lorentz-invariant boost axis n?
14.26 Chiral decomposition of a Dirac spinor
A general (non-null or null) Dirac spinor ϕ can be decomposed into a sum of left- and right-handed chiral
components
ϕ = ϕL + ϕR , (14.209)
that are eigenvectors of the chiral operator γ5,
γ5ϕLR
= ∓ϕLR. (14.210)
The left- and right-handed chiral components can be projected out by applying the chiral projection operators12 (1∓ γ5) (which are projection operators because their squares are themselves):
1
2(1∓ γ5)ϕ = ϕL
R. (14.211)
The decomposition into chiral components is Lorentz-invariant because the pseudoscalar I, hence the chiral
operator γ5 ≡ iI, is Lorentz-invariant, which is true because the pseudoscalar I commutes with any Lorentz
rotor. The chiral projection operators are null, because the reverse of the chiral operator is γ5 = iI = −iI =
−γ5, and its square is one, γ25 = 1, so
(
1 + γ5
)
(1 + γ5) = (1− γ5)(
1− γ5
)
= (1− γ5) (1 + γ5) = 0 . (14.212)
238 ∗The geometric algebra
Consequently each of the chiral components is null, hence massless,
ϕLϕL = ϕRϕR = 0 . (14.213)
Since a pure left- or right-handed spinor must be null, a non-null Dirac particle cannot be purely left- or
right-handed.
14.27 Dirac equation
The Dirac equation is the relativistic quantum mechanical wave equation for spin- 12 particles. By itself, the
Dirac equation does not provide a consistent theory of relativistic quantum mechanics, because in relativistic
quantum mechanics there is no such thing as a single particle that evolves in isolation. Rather, a “funda-
mental” particle such as an electron is dressed in a sea of particle-antiparticle pairs polarized out of the
vacuum by the presence of the electron. Nevertheless, the Dirac equation is a fundamental building block
for the quantum field theory of spin- 12 particles.
The Dirac theory starts with the momentum 4-vector in the form p = γγmpm, where γγm are not only the
basis vectors of an orthonormal tetrad, but also the basis vectors of the spacetime (Clifford) algebra. In the
Dirac slash notation p =6p, but the notation is superfluous here. For a particle of rest mass m, the geometric
square of the momentum is
pp = γγmγγnpmpn = pnpn = −m2 . (14.214)
The vanishing sum pp +m2 factors as
pp +m2 = (p + im)(p− im) = 0 . (14.215)
The factorization provides the motivation for the Dirac wave equation for a free relativistic spin- 12 particle
or antiparticle,
(p− im)ϕ = 0 (particle) , (14.216a)
(p + im)ϕ = 0 (antiparticle) , (14.216b)
in which ϕ is a Dirac spinor, and the momentum operator p ≡ γγmpm should be replaced, according to the
usual rules of quantum mechanics, by minus i times the gradient operator ∂ = γγm∂m,
p = −i∂ . (14.217)
With respect to locally inertial coordinates xm ≡ t, xi,
p0 = −p0 = i∂
∂t, pi = pi = −i ∂
∂xi. (14.218)
With the replacement (14.217), the Dirac equations (14.216) are
(∂ +m)ϕ = 0 (particle) , (14.219a)
(∂ −m)ϕ = 0 (antiparticle) . (14.219b)
14.28 Antiparticles are negative mass particles moving backwards in time 239
The reason for choosing opposite signs for the mass m in the particle and antiparticle equations is discussed
further in §14.28. For now, two comments can be made about the different choice of sign. Firstly, as seen in
§14.24, particles and antiparticles belong to distinct representations that do not mix under Lorentz transfor-
mations, so it is consistent to allow them to satisfy different equations. Secondly, the factorization (14.215)
involves factors with both signs of m, so it is reasonable — one could say demanded — that both equations
would occur.
In flat (Minkowski) space, the Dirac wave equations (14.219) for a free particle or antiparticle are most
easily solved by Fourier transforming with respect to space and time. The differential wave equations (14.219)
then revert to being algebraic equations (14.216), with p being the momentum of the corresponding Fourier
mode. If the Dirac spinor ϕ is a particle as opposed to an antiparticle, so that ϕ is (up to an irrelevant real
scale factor) R⇑ where ⇑ is any rest-frame particle eigenvector, then the following calculation
pϕ = (Rmγγ0R)(R⇑) = mRγγ0⇑ = imR⇑ = imϕ (particle) , (14.220)
confirms that the particle Dirac equation (14.216a) recovers the expected momentum p = Rmγγ0R, which
is the rest frame momentum p ≡ pmγγm = mγγ0 Lorentz transformed into the tetrad frame. Likewise, if the
Dirac spinor ϕ is an antiparticle as opposed to a particle, so that ϕ is (up to an irrelevant real scale factor)
R⇓ where ⇓ is any rest-frame antiparticle eigenvector, then
pϕ = (Rmγγ0R)(R⇓) = mRγγ0⇓ = −imR⇓ = −imϕ (antiparticle) (14.221)
confirms that the antiparticle Dirac equation (14.216b) again recovers the expected momentum p = Rmγγ0R.
The sign flip between the particle and antiparticle equations (14.220) and (14.221) occurs because the
time basis vector γγ0 yields opposite signs when acting on a rest-frame particle eigenvector ⇑ versus rest-
frame antiparticle eigenvector ⇓, equation (14.177). The same sign flip occurs if the rest-frame antiparticle
eigenvector is taken to be I⇑ (per §14.24) instead of ⇓, since γγ0 anticommutes with I.
Fourier-transformed back into real space, the free Dirac spinor wavefunctions ϕ in flat space are, for either
particles or antiparticles, FREQUENCY HAS WRONG SIGN
ϕ = ϕ0 eipmxm
, (14.222)
where pm is the momentum, satisfying p = γγmpm, and ϕ0 is the value of the Dirac spinor at the origin
xm = 0. Regarded as an element of the spacetime algebra, the exponential factor eipmxm
is a scalar, so it
commutes with ϕ0, so it does not matter on which side of ϕ0 the exponential is placed.
14.28 Antiparticles are negative mass particles moving backwards in time
The original Dirac treatment took the particle Dirac equation (14.219a) as describing both particles and
antiparticles. This led to solutions in which the free-wave factor contained not only a positive frequency
component, as in equation (14.222), but also a negative frequency component e−ipmxm
. These negative
frequency components were interpreted as indicating an antiparticle, with negative mass m.
In the original Dirac theory, the prediction of particles with negative rest mass was problematic: where are
240 ∗The geometric algebra
such particles? And if they existed, why wouldn’t pairs of positive and negative mass particles spontaneously
pop out of the vacuum, causing a catastrophic breakdown of the vacuum? To solve the problem, Dirac
proposed that all negative energy states of the vacuum are already occupied, and that antiparticles correspond
to holes in the negative energy sea.
Dirac’s conundrum was eventually solved by Feynman, who realised that anti-particles are equivalent
to negative mass particles moving backwards in time. A negative mass particle moving backwards in
time looks like a positive mass particle moving forwards in time. Feynman’s solution obviates the need for
any negative energy sea of antiparticles. Feynman’s dictum corresponds mathematically to choosing particle
and antiparticle spinors not only to yield opposite signs when acted on by the time basis vector γγ0 in their
rest frames, as in Dirac theory, but also to have the opposite sign of mass m in the Dirac equations (14.219a).
This is the approach adopted in the previous section, §14.27.
14.29 Dirac equation with electromagnetism
The Dirac equation for a spin- 12 particle of charge e moving in an external electromagnetic field with potential
A ≡ Amγγm is given by the same Dirac equations (14.216), but now the momentum p is rewritten in terms
of the canonical momentum π,
p = π + eA , (14.223)
and it is the canonical momentum π that is replaced by minus i times the gradient operator ∂ ≡ γγm∂m,
π = −i∂ . (14.224)
Consequently the Dirac equation for charged particle or antiparticle is
(∂ + ieA +m)ϕ = 0 (particle) , (14.225a)
(∂ + ieA−m)ϕ = 0 (antiparticle) . (14.225b)
Equation (14.225b) appears to describe an antiparticle as having mass −m opposite to that of a particle,
and charge e the same as that of a particle. If the antiparticle is interpreted as having negative mass
moving backwards in time, then the antiparticle has positive mass m moving forwards in time, and charge
−e opposite to that of a particle.
Equations (14.225) are not easy to solve in general. Analytic solutions exist in some cases, such as when
the electromagnetic field consists of a uniform magnetic field B.
14.30 CPT
It was seen in §14.24 that multiplying a non-null Dirac spinor ϕ by the pseudoscalar I converts a particle
spinor into an antiparticle spinor, and vice versa. This operation is conventionally called CPT ,
CPT : ϕ→ Iϕ . (14.226)
14.31 Charge conjugation C 241
The operation is called CPT because it is conventionally parsed into 3 distinct discrete transformations C,
P , and T , discussed in turn below. In the spacetime algebra, the CPT operation flips the sign of all the
spacetime axes γγm,
CPT : γγm → IγγmI−1 = −γγm , (14.227)
which is true because the pseudoscalar I anti-commutes with each of the basis vectors γγm. The CPT
operation leaves all even multivectors unchanged since I commutes with all even multivectors. In particular,
CPT is Lorentz invariant, since I commutes with Lorentz rotors.
The CPT operation converts the particle Dirac equation (14.225a) into the antiparticle Dirac equa-
tion (14.225b):
CPT : −I(∂ + ieA +m)ϕ = (∂ + ieA−m)Iϕ . (14.228)
Equation (14.228) shows that the spinor Iϕ satisfies the Dirac equation (14.225b) for an antiparticle, con-
sistent with conclusion of §14.24 that multiplying a non-null Dirac spinor by I converts a particle into an
antiparticle.
14.31 Charge conjugation C
A Dirac spinor, non-null or null, contains two distinct components that remain separate under Lorentz
transformations. For a non-null spinor the two components are particles and antiparticles. For a null spinor
the two components are the left- and right-handed chiralities. For a non-null spinor, the CPT operation
of multiplying the spinor by I converts a particle into an antiparticle and versa. But for a null spinor,
multiplying by I = −iγ5 leaves left- and right-handed particles as they are: it does not transform opposite
chiralities into each other.
A charge conjugation operation C can be defined with the property that it converts particles of one type
into the opposite type for both non-null and null spinors: particles into antiparticles and vice versa, and
left-handed into right-handed chiralities and vice versa.
In quantum mechanics, it is natural to regard particles and antiparticles as belonging to complex (with
respect to i) conjugate representations. The charge conjugation operator C is defined by the requirement
that it transforms the spacetime basis vectors γγm to their complex conjugates (with respect to i)
C : γγm → CγγmC−1 = γγ∗
m . (14.229)
In the Dirac representation (14.170), the condition (14.229) requires that C commute with γγ2, but anticom-
mute with γγ0, γγ1, and γγ3. A suitable matrix is γγ2 itself,
C = γγ2 . (14.230)
Charge conjugation of a Dirac spinor ϕ is accomplished by taking the complex conjugate (with respect to i)
of the spinor, and multiplying by the charge conjugation operator C:
C : ϕ→ Cϕ∗ . (14.231)
242 ∗The geometric algebra
The charge-conjugates Cϕ∗ of the non-null basis eigenvectors (14.176) in the Dirac representation are
C(⇑↑)∗ = ⇓↓ , C(⇑↓)∗ = −⇓↑ , C(⇓↑)∗ = −⇑↓ , C(⇓↓)∗ = ⇑↑ , (14.232)
while the charge-conjugates of the null chiral basis eigenvectors (14.206) are
C
[
(⇑ ∓ ⇓) ↑√2
]∗= ± (⇑ ± ⇓) ↓√
2. (14.233)
Equations (14.232) and (14.233) show that charge conjugation not only converts rest-frame particle eigen-
vectors ⇑ into rest-frame antiparticle eigenvectors ⇓, and vice versa, but also converts a left-handed null
eigenvector into a right-handed null eigenvector and vice versa. Since the operation of complex conjuga-
tion commutes with Lorentz transformation (because Lorentz rotors R are independent of the quantum
mechanical imaginary i), it is true in general that charge conjugation switches particles into antiparticles,
and left-handed into right-handed chiralities.
The complex conjugate (with respect to i) of the charged particle Dirac equation (14.225a) is
(∂∗ − ieA∗ +m)ϕ∗ = 0 . (14.234)
Complex conjugation leaves the components ∂m and Am of the gradient operator and electromagnetic po-
tential unchanged, but conjugates the basis vectors γγm. Left-multiplying equation (14.234) by C, and
commuting C through the wave operator, yields
(∂ − ieA +m)Cϕ∗ = 0 . (14.235)
Thus the charge-conjugated spinor Cϕ∗ satisfies the Dirac equation (14.225a) for a particle with the same
mass m but opposite charge −e.
14.32 Parity reversal P
The parity operation P is the operation of reversing all the spatial axes, while keeping the time axis un-
changed,
P : γγm → PγγmP−1 =
γγm m = 0 ,
−γγm m = 1, 2, 3 .(14.236)
A suitable matrix is the time axis γγ0
P = γγ0 . (14.237)
Parity reversal transforms a Dirac spinor ϕ as
P : ϕ→ Pϕ . (14.238)
Parity reversal commutes with spatial rotations, but not with Lorentz boosts. Parity reversal transforms the
left- and right-handed chiral components of a Dirac spinor into each other:
P : ϕL ↔ ϕR , (14.239)
14.33 Time reversal T 243
which is true because parity reversal flips the sign of the pseudoscalar, PIP−1 = −I, hence also the sign of
the chiral operator γ5 ≡ iI.
14.33 Time reversal T
The Wigner time reversal operation T is conventionally defined so that CPT is the product of the three
operations C, P , and T . This requires that time-reversal transforms a Dirac spinor ϕ as
T : ϕ→ Tϕ∗ (14.240)
with
T = γγ1γγ3 , (14.241)
so that
CPT = γγ2γγ0γγ1γγ3 = I . (14.242)
14.34 Majorana spinor
The charge conjugation operation considered in §14.31 switched left- and right-handed null spinors into each
other. However, it is possible for a null spinor to be its own antiparticle. In this case the complex (with
respect to i) conjugate representation is itself, rather than being a distinct representation. A null spinor
which is its own antiparticle is a Majorana spinor.
For a Majorana spinor, complex conjugation should leave the representation unchanged; that is, the γ-
matrices should be real, in contrast to the Dirac representation (14.170), where the γ-matrices are complex.
A suitable real representation of the γ-matrices is
γγ0 =
(
0 iσ2
iσ2 0
)
, γγ1 =
(
σ1 0
0 σ1
)
, γγ2 =
(
0 −iσ2
iσ2 0
)
, γγ3 =
(
σ3 0
0 σ3
)
. (14.243)
The chiral matrix γ5 ≡ iI is
γ5 =
(
σ2 0
0 σ2
)
. (14.244)
14.35 Covariant derivatives revisited
Under a Lorentz transformation by rotor R, any multivector a transforms as a → RaR. The covariant
derivative D ≡ γγmDm must transform likewise as
R : D → RDR . (14.245)
244 ∗The geometric algebra
14.36 General relativistic Dirac equation
To convert the Dirac equations (14.219) or (14.225) into general relativistic equations, the derivative ∂ ≡γγm∂m must be converted to a covariant derivative.
14.37 3D Vectors as rank-2 spinors
THIS NEEDS TO BE CONVERTED FROM 3D TO SPACETIME.
Concept question 14.20 One is used to thinking of a spin- 12 particle as in some sense the square root of
a spin-1 particle, a vector. How is this concept compatible with the idea that a spin- 12 object is a (scaled)
rotor?
A Pauli spinor ϕ contains two complex components ϕa, where the index a runs over the two indices ↑ and
↓. The Pauli spinor is a spinor of rank 1, having one spinor index. Under a rotation by rotor R, the spinor
ϕ transforms as ϕ→ Rϕ. Under a rotation, the spinor components ϕa of the Pauli spinor transform as
ϕa → Rbaϕb (14.246)
where Rba is the special unitary 2 × 2 matrix representing the reverse rotor R. The rotor R itself has
components Rab. Note the placement of indices: for the rotor R, the first index is down and the second up,
while for the reverse rotor R, the first index is up and the second down. The rotor and its reverse are inverse
to each other, satisfying RR = RcaRc
b = δba.
The Hermitian conjugate spinor ϕ† transforms in the opposite way ϕ† → ϕ†R, and can therefore be written
as a spinor with contravariant (raised) index, ϕ† = ϕa. The contravariant spinor components ϕa transform
as
ϕa → ϕbRba . (14.247)
The product of a Hermitial conjugate spinor χ† with another spinor ϕ defines their scalar product, which is
unchanged by a rotation,
χ†ϕ→ χ†RRϕ = χ†ϕ = χaϕa . (14.248)
Explicitly, the scalar product χaϕa is the complex number
χaϕa = χ∗↑ϕ↑ + χ∗
↓ϕ↓ . (14.249)
An element a of the 3D geometric algebra transforms under rotation as a→ RaR. The behaviour under
rotation shows that a 3D multivector is a rank-2 spinor a = aab, a complex 2×2 matrix, with the first index
being covariant (lowered), and the second being contravariant (raised). The spinor components aab of the
multivector transform as
aab → Rc
a acdRd
b . (14.250)
14.37 3D Vectors as rank-2 spinors 245
A 3D multivector a contains 8 real components, comprising a scalar, a vector, a pseudovector, and a pseu-
doscalar. All 8 complex components are represented by the complex 2 × 2 spinor aab. The scalar and
pseudoscalar components are contained in the complex trace aaa of the spinor. The rotational transforma-
tion a → RaR preserves the grade of the multivector, so scalars transform to scalars, vectors to vectors,
pseudovectors to pseudovectors, and pseudoscalars to pseudoscalars.
A reversed multivector a transforms in the same way a→ RaR as the parent multivector a. Consquently,
like the multivector a, the reversed multivector a transforms like a rank-2 spinor with one covariant (lowered)
and one contravariant (raised) index. It is consistent to write the reversed spinor as the original spinor with
the first index raised and the second lowered, a = aab.
The above has shown that a 3D multivector a is represented naturally as a rank-2 spinor aab with one
covariant and one contravariant index. This might seem strange: one might have expected that a vector — a
spin-1 particle — might be represented as a rank-2 spinor aab with two covariant indices — a sum of tensor
products ϕ⊗ χ = ϕaχb of two spin- 12 particles. Under a rotation, the components aab of a covariant rank-2
spinor transform as
aab → RcaR
db acd . (14.251)
One aspect of this transformation is straightforward: the following combination of tensor products of spinors
is invariant under rotation, and is therefore a scalar:
↑ ⊗ ↓ − ↓ ⊗ ↑ . (14.252)
The remaining tensor combinations ↑ ⊗ ↑, ↓ ⊗ ↓, and ↑ ⊗ ↓ + ↓ ⊗ ↑ provide a basis for vectors, but under
rotation they transform into complex, not real, linear combinations of each other. This contrasts with the
representation of multivectors a by spinors aab with one covariant and one contravariant index, where the
grade-preserving property of the rotational transformation a → RaR ensures that vectors rotate into real,
not complex, linear combinations of each other.
PART SIX
BLACK HOLE INTERIORS
Concept Questions
1. Explain how the equation for the Gullstrand-Painleve metric (15.16) encodes not merely a metric but a
full vierbein.
2. In what sense does the Gullstrand-Painleve metric (15.16) depict a flow of space? [Are the coordinates
moving? If not, then what is moving?]
3. If space has no substance, what does it mean that space falls into a black hole?
4. Would there be any gravitational field in a spacetime where space fell at constant velocity instead of
accelerating?
5. In spherically symmetric spacetimes, what is the most important Einstein equation, the one that causes
Reissner-Nordstrom black holes to be repulsive in their interiors, and causes mass inflation in non-empty
(non Reissner-Nordstrom) charged black holes?
What’s important?
1. The tetrad formalism provides a firm mathematical foundation for the concept that space falls faster than
light inside a black hole.
2. Whereas the Kerr-Newman geometry of an ideal rotating black hole contains inside its horizon wormhole
and white hole connections to other universes, real black holes are subject to the mass inflation stability
discovered by Eric Poisson & Werner Israel (1990, “Internal structure of black holes,” Phys. Rev. D 41,
1796-1809).
15
Black hole waterfalls
15.1 Tetrads move through coordinates
As already discussed in §11.3, the way in which metrics are commonly written, as a (weighted) sum of squares
of differentials,
ds2 = γmn em
µ en
ν dxµdxν , (15.1)
encodes not only a metric gµν = γmn em
µ en
ν , but also an inverse vierbein emµ, and consequently a vierbein
emµ, and associated tetrad γγm. Most commonly the tetrad metric is orthonormal (Minkowski), γmn = ηmn,
but other tetrad metrics, such as Newman-Penrose, occur. Usually it is self-evident from the form of the
line-element what the tetrad metric γmn is in any particular case.
If the tetrad is orthonormal, γmn = ηmn, then the 4-velocity um of an object at rest in the tetrad, or
equivalently the 4-velocity of the tetrad rest frame itself, is
um = 1, 0, 0, 0 . (15.2)
The tetrad-frame 4-velocity (15.2) of the tetrad rest frame is transformed to a coordinate-frame 4-velocity
uµ in the usual way, by applying the vierbein,
dxµ
dτ≡ uµ = em
µum = e0µ . (15.3)
Equation (15.3) says that the tetrad rest frame moves through the coordinates at coordinate 4-velocity given
by the zero’th row of the vierbein, dxµ/dτ = e0µ. The coordinate 4-velocity uµ is related to the lapse α and
shift βα in the ADM formalism by uµ = 1, βα/α, equation (13.10).
The idea that locally inertial frames move through the coordinates provides the simplest way to conceptu-
alize black holes. The motion of locally inertial frames through coordinates is what is meant by the “dragging
of inertial frames” around rotating masses.
252 Black hole waterfalls
Figure 15.1 The fish upstream can make way against the current, but the fish downstream is swept to thebottom of the waterfall. Art by Wildrose Hamilton.
15.2 Gullstrand-Painleve waterfall
The Gullstrand-Painleve metric is a version of the metric for a spherical (Schwarzschild or Reissner-Nordstrom)
black hole discovered in 1921 independently by Allvar Gullstrand (1922, “Allgemeine Losung des statischen
Einkorperproblems in der Einsteinschen Gravitationstheorie,” translated by http://babelfish.altavista.
com/tr as “General solution of the static body problem in Einstein’s Gravitation Theory,” Arkiv. Mat. As-
tron. Fys. 16(8), 1–15) and Paul Painleve (1921, “La mecanique classique et la theorie de la relativite”, C.
R. Acad. Sci. (Paris) 173, 677–680). Although Gullstrand’s paper was published in 1922, after Painleve’s, it
appears that Gullstrand’s work has priority. Gullstrand’s paper was dated 25 May 1921, whereas Painleve’s
is a write up of a presentation to the Academie des Sciences in Paris on 24 October 1921. Moreover, Gull-
strand seems to have had a better grasp of what he had discovered than Painleve, for Gullstrand recognized
that observables such as the redshift of light from the Sun are unaffected by the choice of coordinates in
the Schwarzschild geometry, whereas Painleve, noting that the spatial metric was flat at constant free-fall
time, dtff = 0, concluded in his final sentence that, as regards the redshift of light and such, “c’est pure
imagination de pretendre tirer du ds2 des consequences de cette nature”.
Although neither Gullstrand nor Painleve understood it, their metric paints a picture of space falling like
a river, or waterfall, into a spherical black hole, Figure 15.1. The river has two key features: first, the river
flows in Galilean fashion through a flat Galilean background, equation (15.19); and second, as a freely-falling
fishy swims through the river, its 4-velocity, or more generally any 4-vector attached to it, evolves by a
series of infinitesimal Lorentz boosts induced by the change in the velocity of the river from place to place,
equation (15.24). Because the river moves in Galilean fashion, it can, and inside the horizon does, move
15.2 Gullstrand-Painleve waterfall 253
faster than light through the background coordinates. However, objects moving in the river move according
to the rules of special relativity, and so cannot move faster than light through the river.
15.2.1 Gullstrand-Painleve tetrad
The Gullstrand-Painleve metric (7.26) is
ds2 = − dt2ff + (dr − β dtff)2 + r2(dθ2 + sin2θ dφ2) , (15.4)
where β is defined to be the radial velocity of a person who free-falls radially from rest at infinity,
β =dr
dτ=
dr
dtff, (15.5)
and tff is the free-fall time, the proper time experienced by a person who free-falls from rest at infinity. The
radial velocity β is the (apparently) Newtonian escape velocity
β = ∓√
2M(r)
r, (15.6)
where M(r) is the interior mass within radius r, and the sign is − (infalling) for a black hole, + (outfalling)
for a white hole. For the Schwarzschild or Reissner-Nordstrom geometry the interior mass M(r) is the mass
M at infinity minus the mass Q2/2r in the electric field outside r,
M(r) = M − Q2
2r. (15.7)
Figure 15.2 illustrates the velocity fields in Schwarzschild and Reissner-Nordstrom black holes. Horizons
occur where the radial velocity β equals the speed of light
β = ∓1 , (15.8)
with − for black hole solutions, + for white hole solutions. The phenomenology of Schwarzschild and
Reissner-Nordstrom black holes has already been explored in Chapters 7 and 8.
Exercise 15.1 Schwarzschild to Gullstrand-Painleve. Show that the Schwarzschild metric transforms
into the Gullstrand-Painleve metric under the coordinate transformation of the time coordinate
dtff = dt− β
1− β2dr . (15.9)
Exercise 15.2 Radial free-fall from rest. Confirm that β given by equation (15.6) is indeed the
velocity (15.5) of a person who free-falls radially from rest at infinity in the Reissner-Nordstrom geometry.
254 Black hole waterfalls
Horizon
Inner horizon
Outer horizon
Turnaround
Figure 15.2 Radial velocity β in (upper panel) a Schwarzschild black hole, and (lower panel) a Reissner-Nordstrom black hole with electric charge Q = 0.96.
The Gullstrand-Painleve line-element (15.4) encodes an inverse vierbein with an orthonormal tetrad metric
γmn = ηmn through
e0µ dxµ = dtff , (15.10a)
e1µ dxµ = dr − β dtff , (15.10b)
e2µ dxµ = r dθ , (15.10c)
e3µ dxµ = r sin θ dφ . (15.10d)
Explicitly, the inverse vierbein emµ of the Gullstrand-Painleve line-element (15.4), and the corresponding
15.2 Gullstrand-Painleve waterfall 255
vierbein emµ, are the matrices
emµ =
1 0 0 0
−β 1 0 0
0 0 r 0
0 0 0 r sin θ
, emµ =
1 β 0 0
0 1 0 0
0 0 1/r 0
0 0 0 1/(r sin θ)
. (15.11)
According to equation (15.3), the coordinate 4-velocity of the tetrad frame through the coordinates is
dtffdτ
,dr
dτ,dθ
dτ,dφ
dτ
= uµ = e0µ = 1, β, 0, 0 , (15.12)
consistent with the claim (15.5) that β represents a radial velocity, while tff coincides with the proper time
in the tetrad frame.
The tetrad and coordinate axes γγm and gµ are related to each other by the vierbein and inverse vierbein
in the usual way, γγm = emµ gµ and gµ = em
µ γγm. The Gullstrand-Painleve orthonormal tetrad axes γγm are
thus related to the coordinate axes gµ by
γγ0 = gtff + βgr , γγ1 = gr , γγ2 = gθ/r , γγ3 = gφ/(r sin θ) . (15.13)
Physically, the Gullstrand-Painleve-Cartesian tetrad (15.13) are the axes of locally inertial orthonormal
frames (with spatial axes γγi oriented in the polar directions r, θ, φ) attached to observers who free-fall
radially, without rotating, starting from zero velocity and zero angular momentum at infinity. The fact
that the tetrad axes γγm are parallel-transported, without precessing, along the worldlines of the radially
free-falling observers can be confirmed by checking that the tetrad connections Γnm0 with final index 0 all
vanish, which implies that
dγγm
dτ= ∂0γγm ≡ Γn
m0γγn = 0 . (15.14)
That the proper time derivative d/dτ in equation (15.14) of a person at rest in the tetrad frame, with
4-velocity (15.2), is equal to the directed time derivative ∂0 follows from
d
dτ= uµ ∂
∂xµ= um∂m = ∂0 . (15.15)
15.2.2 Gullstrand-Painleve-Cartesian tetrad
The manner in which the Gullstrand-Painleve line-element depicts a flow of space into a black hole is eluci-
dated further if the line-element is written in Cartesian rather than spherical polar coordinates. Introduce a
Cartesian coordinate system xµ ≡ tff , xi ≡ tff , x, y, z. The Gullstrand-Painleve metric in these Cartesian
coordinates is
ds2 = − dt2ff + δij(dxi − βidtff)(dxj − βjdtff) , (15.16)
256 Black hole waterfalls
with implicit summation over spatial indices i, j = 1, 2, 3. The βi in the metric (15.16) are the components
of the radial velocity expressed in Cartesian coordinates
βi = βx
r,y
r,z
r
. (15.17)
The inverse vierbein emµ and vierbein em
µ encoded in the Gullstrand-Painleve-Cartesian line-element (15.16)
are
emµ =
1 0 0 0
−β1 1 0 0
−β2 0 1 0
−β3 0 0 1
, emµ =
1 β1 β2 β3
0 1 0 0
0 0 1 0
0 0 0 1
. (15.18)
The tetrad axes γγm of the Gullstrand-Painleve-Cartesian line-element (15.16) are related to the coordinate
tangent axes gµ by
γγ0 = gtff + βigi , γγi = gi , (15.19)
and conversely the coordinate tangent axes gµ are related to the tetrad axes γγm by
gtff = γγ0 − βiγγi , gi = γγi . (15.20)
Note that the tetrad-frame contravariant components βi of the radial velocity coincide with the coordinate-
frame contravariant components βi; for clarification of this point see the more general equation (15.48)
for a rotating black hole. The Gullstrand-Painleve-Cartesian tetrad axes (15.19) are the same as the tetrad
axes (15.13), but rotated to point in Cartesian directions x, y, z rather than in polar directions r, θ, φ. Like the
polar tetrad, the Cartesian tetrad axes γγm are parallel-transported, without precessing, along the worldlines
of radially free-falling observers, as can be confirmed by checking once again that the tetrad connections
Γnm0 with final index 0 all vanish.
Remarkably, the transformation (15.19) from coordinate to tetrad axes is just a Galilean transformation
of space and time, which shifts the time axis by velocity β along the direction of motion, but which leaves
unchanged both the time component of the time axis and all the spatial axes. In other words, the black
hole behaves as if it were a river of space that flows radially inward through Galilean space and time at the
Newtonian escape velocity.
15.2.3 Gullstrand-Painleve fishies
The Gullstrand-Painleve line element paints a picture of locally inertial frames falling like a river of space into
a spherical black hole. What happens to fishies swimming in that river? Of course general relativity supplies
a mathematical answer in the form of the geodesic equation of motion (15.21). Does that mathematical
answer lead to further conceptual insight?
Consider a fishy swimming in the Gullstrand-Painleve river, with some arbitrary tetrad-frame 4-velocity
15.2 Gullstrand-Painleve waterfall 257
um, and consider a tetrad-frame 4-vector pk attached to the fishy. If the fishy is in free-fall, then the geodesic
equation of motion for pk is as usual
dpk
dτ+ Γk
mnunpm = 0 . (15.21)
As remarked in §11.13, for a constant (for example Minkowski) tetrad metric, as here, the tetrad connections
Γkmn constitute a set of four generators of Lorentz transformations, one in each of the directions n. In
particular Γkmnu
n is the generator of a Lorentz transformation along the path of a fishy moving with 4-
velocity un. In a small (infinitesimal) time δτ , the fishy moves a proper distance δξn ≡ unδτ relative to the
infalling river. This proper distance δξn = enνδx
ν = δnν (δxν − βνδtff) = δxn − βnδτ equals the distance
δxn moved relative to the background Gullstrand-Painleve-Cartesian coordinates, minus the distance βnδτ
moved by the river. The geodesic equation (15.21) says that the change δpk in the tetrad 4-vector pk in the
time δτ is
δpk = −Γkmnδξ
npm . (15.22)
Equation (15.22) describes an infinitesimal Lorentz transformation −Γkmnδξ
n of the 4-vector pk.
Equation (15.22) is quite general in general relativity: it says that as a 4-vector pk free-falls through a
system of locally inertial tetrads, it finds itself Lorentz-transformed relative those tetrads. What is special
about the Gullstrand-Painleve-Cartesian tetrad is that the tetrad-frame connections, computed by the usual
formula (11.41), are given by the coordinate gradient of the radial velocity (the following equation is valid
component-by-component despite the non-matching up-down placement of indices)
Γ0ij = Γi
0j =∂βi
∂xj(i, j = 1, 2, 3) . (15.23)
The same property, that the tetrad connections are a pure coordinate gradient, holds also for the Doran-
Cartesian tetrad for a rotating black hole, equation (15.51). With the connections (15.23), the change
δpk (15.22) in the tetrad 4-vector is
δp0 = − δβi pi , δpi = − δβi p0 , (15.24)
where δβi is the change in the velocity of the river as seen in the tetrad frame,
δβi = δξj ∂βi
∂xj. (15.25)
But equation (15.24) is nothing more than an infinitesimal Lorentz boost by a velocity change δβi. This
shows that a fishy swimming in the river follows the rules of special relativity, being Lorentz boosted by tidal
changes δβi in the river velocity from place to place.
Is it correct to interpret equation (15.25) as giving the change δβi in the river velocity seen by a fishy? Of
course general relativity demands that equation (15.25) be mathematically correct; the issue is merely one
of interpretation. Shouldn’t the change in the river velocity really be
δβi ?= δxν ∂β
i
∂xν, (15.26)
258 Black hole waterfalls
where δxν is the full change in the coordinate position of the fishy? No. Part of the change (15.26) in the
river velocity can be attributed to the change in the velocity of the river itself over the time δτ , which is
δxνriver∂β
i/∂xν with δxνriver = βνδτ = βνδtff . The change in the velocity relative to the flowing river is
δβi = (δxν − δxνriver)
∂βi
∂xν= (δxν − βνδtff)
∂βi
∂xν(15.27)
which reproduces the earlier expression (15.25). Indeed, in the picture of fishies being carried by the river,
it is essential to subtract the change in velocity of the river itself, as in equation (15.27), because otherwise
fishies at rest in the river (going with the flow) would not continue to remain at rest in the river.
15.3 Boyer-Lindquist tetrad
The Boyer-Lindquist metric for an ideal rotating black hole was explored already in Chapter 9. With the
tetrad formalism in hand, the advantages of the Boyer-Lindquist tetrad for portraying the Kerr-Newman
geometry become manifest. With respect to the orthonormal Boyer-Lindquist tetrad, the electromagnetic
field is purely radial, and the energy-momentum and Weyl tensors are diagonal. The Boyer-Lindquist tetrad
is aligned with the principal (ingoing or outgoing) null congruences.
The Boyer-Lindquist orthonormal tetrad is encoded in the Boyer-Lindquist metric
ds2 = − ∆
ρ2
(
dt− a sin2θ dφ)2
+ρ2
∆dr2 + ρ2dθ2 +
R4 sin2θ
ρ2
(
dφ− a
R2dt)2
(15.28)
where
R ≡√
r2 + a2 , ρ ≡√
r2 + a2 cos2θ , ∆ ≡ R2 − 2Mr +Q2 = R2(1− β2) . (15.29)
Explicitly, the vierbein emµ of the Boyer-Lindquist orthonormal tetrad is
emµ =
1
ρ
R2/√
∆ 0 0 a/√
∆
0√
∆ 0 0
0 0 1 0
a sin θ 0 0 1/ sin θ
, (15.30)
with inverse vierbein emµ
emµ =
√∆/ρ 0 0 − a sin2θ
√∆/ρ
0 ρ/√
∆ 0 0
0 0 ρ 0
− a sin θ/ρ 0 0 R2 sin θ/ρ
. (15.31)
With respect to the Boyer-Lindquist tetrad, only the time component At of the electromagnetic potential
Am is non-vanishing,
Am =
Qr
ρ√
∆, 0, 0, 0
. (15.32)
15.4 Doran waterfall 259
Only the radial components Er and Br of the electric and magnetic fields are non-vanishing, and they are
given by the complex combination
Er + i Br =Q
(r − ia cos θ)2, (15.33)
or explicitly
Er =Q(
r2−a2 cos2θ)
ρ4, Br =
2Qar cos θ
ρ4. (15.34)
The electrogmagnetic field (15.33) satisfies Maxwell’s equations (11.64) and (11.65) with zero electric charge
and current, jn = 0, except at the singularity ρ = 0.
The non-vanishing components of the tetrad-frame Einstein tensor Gmn are
Gmn =Q2
ρ4
1 0 0 0
0 −1 0 0
0 0 1 0
0 0 0 1
, (15.35)
which is the energy-momentum tensor of the electromagnetic field. The non-vanishing components of the
tetrad-frame Weyl tensor Cklmn are
− 12 Ctrtr = 1
2 Cθφθφ = Ctθtθ = Ctφtφ = −Crθrθ = −Crφrφ = ReC , (15.36a)
12 Ctrθφ = Ctθrφ = −Ctφrθ = ImC , (15.36b)
where C is the complex Weyl scalar
C = − 1
(r − ia cos θ)3
(
M − Q2
r + ia cos θ
)
. (15.37)
In the Boyer-Lindquist tetrad, the photon 4-velocity vm on the principal null congruences is radial,
vt = ± ρ√∆, vr = ± ρ√
∆, vθ = 0 , vφ = 0 . (15.38)
Exercise 15.3 Dragging of inertial frames around a Kerr-Newman black hole. What is the
coordinate-frame 4-velocity uµ of the Boyer-Lindquist tetrad through the Boyer-Lindquist coordinates?
15.4 Doran waterfall
The picture of space falling into a black hole like a river or waterfall works also for rotating black holes. For
Kerr-Newman rotating black holes, the counterpart of the Gullstrand-Painleve metric is the Doran (2000)
metric.
260 Black hole waterfalls
The space river that falls into a rotating black hole has a twist. One might have expected that the rotation
of the black hole would be manifested by a velocity that spirals inward, but this is not the case. Instead,
the river is characterized not merely by a velocity but also by a twist. The velocity and the twist together
comprise a 6-dimensional river bivector ωkm, equation (15.52) below, whose electric part is the velocity,
and whose magnetic part is the twist. Recall that the 6-dimensional group of Lorentz transformations is
generated by a combination of 3-dimensional Lorentz boosts and 3-dimensional spatial rotations. A fishy
that swims through the river is Lorentz boosted by tidal changes in the velocity, and rotated by tidal changes
in the twist, equation (15.61).
Thanks to the twist, unlike the Gullstrand-Painleve metric, the Doran metric is not spatially flat at
constant free-fall time tff . Rather, the spatial metric is sheared in the azimuthal direction. Just as the
velocity produces a Lorentz boost that makes the metric non-flat with respect to the time components, so
also the twist produces a rotation that makes the metric non-flat with respect to the spatial components.
15.4.1 Doran-Cartesian coordinates
In place of the polar coordinates r, θ, φff of the Doran metric, introduce corresponding Doran-Cartesian
coordinates x, y, z with z taken along the rotation axis of the black hole (the black hole rotates right-
handedly about z, for positive spin parameter a)
x ≡ R sin θ cosφff , y ≡ R sin θ sinφff , z ≡ r cos θ . (15.39)
The metric in Doran-Cartesian coordinates xµ ≡ tff , xi ≡ tff , x, y, z, is
ds2 = − dt2ff + δij(
dxi − βiακdxκ) (
dxj − βjαλdxλ)
(15.40)
where αµ is the rotational velocity vector
αµ =
1,ay
R2, − ax
R2, 0
, (15.41)
and βµ is the velocity vector
βµ =βR
ρ
0,xr
Rρ,yr
Rρ,zR
rρ
. (15.42)
The rotational velocity and radial velocity vectors are orthogonal
αµβµ = 0 . (15.43)
For the Kerr-Newman metric, the radial velocity β is
β = ∓√
2Mr −Q2
R(15.44)
with − for black hole (infalling), + for white hole (outfalling) solutions. Horizons occur where
β = ∓1 , (15.45)
15.4 Doran waterfall 261
with β = −1 for black hole horizons, and β = 1 for white hole horizons. Note that the squared magnitude
βµβµ of the velocity vector is not β2, but rather differs from β2 by a factor of R2/ρ2:
βµβµ = βmβ
m =β2R2
ρ2. (15.46)
The point of the convention adopted here is that β(r) is any and only a function of r, rather than depending
also on θ through ρ. Moreover, with the convention here, β is ∓1 at horizons, equation (15.45). Finally, the
4-velocity βµ is simply related to β by βµ = (β/r) ∂r/∂xµ.
The Doran-Cartesian metric (15.40) encodes a vierbein emµ and inverse vierbein em
µ
emµ = δµ
m + αmβµ , em
µ = δmµ − αµβ
m . (15.47)
Here the tetrad-frame components αm of the rotational velocity vector and βm of the radial velocity vector
are
αm = emµαµ = δµ
mαµ , βm = emµβ
µ = δmµ β
µ , (15.48)
which works thanks to the orthogonality (15.43) of αµ and βµ. Equation (15.48) says that the covariant
tetrad-frame components of the rotational velocity vector α are the same as its covariant coordinate-frame
components in the Doran-Cartesian coordinate system, αm = αµ, and likewise the contravariant tetrad-frame
components of the radial velocity vector β are the same as its contravariant coordinate-frame components,
βm = βµ.
15.4.2 Doran-Cartesian tetrad
Like the Gullstrand-Painleve tetrad, the Doran-Cartesian tetrad γγm ≡ γγ0,γγ1,γγ2,γγ3 is aligned with the
Cartesian rest frame gµ ≡ gtff , gx, gy, gz at infinity, and is parallel-transported, without precessing, by
observers who free-fall from zero velocity and zero angular momentum at infinity, as can be confirmed by
checking that the tetrad connections with final index 0 all vanish, Γnm0 = 0, equation (15.14).
Let ‖ and ⊥ subscripts denote horizontal radial and azimuthal directions respectively, so that
γγ‖ ≡ cosφff γγ1 + sinφff γγ2 , γγ⊥ ≡ − sinφff γγ1 + cosφff γγ2 ,
g‖ ≡ cosφff gx + sinφff gy , g⊥ ≡ − sinφff gx + cosφff gy .(15.49)
Then the relation between Doran-Cartesian tetrad axes γγm and the tangent axes gµ of the Doran-Cartesian
metric (15.40) is
γγ0 = gtff + βigi , (15.50a)
γγ‖ = g‖ , (15.50b)
γγ⊥ = g⊥ −a sin θ
Rβigi , (15.50c)
γγ3 = gz . (15.50d)
The relations (15.50) resemble those (15.19) of the Gullstrand-Painleve tetrad, except that the azimuthal
262 Black hole waterfalls
tetrad axis γγ⊥ is shifted radially relative to the azimuthal tangent axis g⊥. This shift reflects the fact that,
unlike the Gullstrand-Painleve metric, the Doran metric is not spatially flat at constant free-fall time, but
rather is sheared azimuthally.
15.4.3 Doran fishies
The tetrad-frame connections equal the ordinary coordinate partial derivatives in Doran-Cartesian coordi-
nates of a bivector (antisymmetric tensor) ωkm
Γkmn = − ∂ωkm
∂xn, (15.51)
which I call the river field because it encapsulates all the properties of the infalling river of space. The
bivector river field ωkm is
ωkm = αkβm − αmβk − ε0kmi ζi , (15.52)
where βm = ηmnβm, the totally antisymmetric tensor εklmn is normalized so that ε0123 = −1, and the vector
ζi points vertically upward along the rotation axis of the black hole
ζi ≡ 0, 0, 0, ζ , ζ ≡ a∫ r
∞
β dr
R2. (15.53)
The electric part of ωkm, where one of the indices is time 0, constitutes the velocity vector βi
ω0i = βi (15.54)
while the magnetic part of ωkm, where both indices are spatial, constitutes the twist vector µi defined by
µi ≡ 12 ε
0ikmωkm = ε0ikmαkβm + ζi . (15.55)
The sense of the twist is that induces a right-handed rotation about an axis equal to the direction of µi by
an angle equal to the magnitude of µi. In 3-vector notation, with µ ≡ µi, α ≡ αi, β ≡ βi, ζ ≡ ζi,
µ ≡ α× β + ζ . (15.56)
In terms of the velocity and twist vectors, the river field ωkm is
ωkm =
0 βx βy βz
−βx 0 µz −µy
−βy −µz 0 µx
−βz µy −µx 0
. (15.57)
Note that the sign of the magnetic part β of ωkm is opposite to the sign of the analogous magnetic field
B associated with an electromagnetic field Fkm; but the adopted signs are natural in that the river field
15.4 Doran waterfall 263
Rot
atio
nax
is
Inner horizon
Outer horizon
360° Rot
atio
nax
is
Inner horizon
Outer horizon
Figure 15.3 (Upper panel) velocity βi and (lower panel) twist µi vector fields for a Kerr black hole with spinparameter a = 0.96. Both vectors lie, as shown, in the plane of constant free-fall azimuthal angle φff . Thevertical bar in the lower panel shows the length of a twist vector corresponding to a full rotation of 360.
induces boosts in the direction of the velocity βi, and right-handed rotations about the twist µi. Like a
static electric field, the velocity vector βi is the gradient of a potential
βi =∂
∂xi
∫ r
β dr , (15.58)
but unlike a magnetic field the twist vector µi is not pure curl: rather, it is µi + ζi that is pure curl.
Figure 15.3 illustrates the velocity and twist fields in a Kerr black hole.
With the tetrad connection coefficients given by equation (15.51), the equation of motion (15.21) for a
264 Black hole waterfalls
4-vector pk attached to a fishy following a geodesic in the Doran river translates to
dpk
dτ=∂ωk
m
∂xnunpm . (15.59)
In a proper time δτ , the fishy moves a proper distance δξm ≡ umδτ relative to the background Doran-
Cartesian coordinates. As a result, the fishy sees a tidal change δωkm in the river field
δωkm = δξn ∂ω
km
∂xn. (15.60)
Consequently the 4-vector pk is changed by
pk → pk + δωkm pm . (15.61)
But equation (15.61) corresponds to an infinitesimal Lorentz transformation by δωkm, equivalent to a Lorentz
boost by δβi and a rotation by δµi.
As discussed previously with regard to the Gullstrand-Painleve river, §15.2.3, the tidal change δωkm,
equation (15.60), in the river field seen by a fishy is not the full change δxν ∂ωkm/∂x
ν relative to the
background coordinates, but rather the change relative to the river
δωkm = (δxν − δxν
river)∂ωk
m
∂xν=[
δxν − βν(δtff − a sin2θ δφff)] ∂ωk
m
∂xν, (15.62)
with the change in the velocity and twist of the river itself subtracted off.
That there exists a tetrad (the Doran-Cartesian tetrad) where the tetrad-frame connections are a coor-
dinate gradient of a bivector, equation (15.51), is a peculiar feature of ideal black holes. It is an intriguing
thought that perhaps the 6 physical degrees of freedom of a general spacetime might always be encoded in
the 6 degrees of freedom of a bivector, but I suspect that that is not true.
Exercise 15.4 River model of the Friedmann-Robertson-Walker metric. Show that the flat FRW
line-element
ds2 = − dt2 + a2(dx2 + x2do2) (15.63)
can be re-expressed as
ds2 = − dt2 + (dr −Hr dt)2 + r2do2 , (15.64)
where r ≡ ax is the proper radial distance, and H ≡ a/a is the Hubble parameter. Interpret the line-
element (15.64). What is the generalization to a non-flat FRW universe?
Exercise 15.5 Program geodesics in a rotating black hole. Write a graphics (Java?) program
that uses the prescription (15.60) to draw geodesics of test particles in an ideal (Kerr-Newman) black
hole, expressed in Doran-Cartesian coordinates. Attach 3D bodies to your test particles, and use the same
prescription (15.60) to rotate the bodies. Implement an option to translate to Boyer-Lindquist coordinates.
16
General spherically symmetric spacetime
16.1 Spherical spacetime
The most important equations in this chapter are the two Einstein equations (16.52). Spherical spacetimes
have 2 physical degrees of freedom. Spherical symmetry eliminates any angular degrees of freedom, leaving
4 adjustable metric coefficients gtt, gtr, grr, and gθθ. But coordinate transformations of the time t and radial
r coordinates remove 2 degrees of freedom, leaving a spherical spacetime with a net 2 physical degrees of
freedom. Spherical spacetimes have 4 distinct Einstein equations (16.30). But 2 of the Einstein equations
serve to enforce energy-momentum conservation, so the evolution of the spacetime is governed by 2 Einstein
equations, in agreement with the number of physical degrees of freedom of spherical spacetime.
The 2 degrees of freedom mean that spherical spacetimes in general relativity have a richer structure than
in Newtonian gravity, which has only degree of freedom, the Newtonian potential Φ. The richer structure
is most striking in the case of the mass inflation instability, Chapter 17, which is an intrinsically general
relativistic instability, with no Newtonian analogue.
16.1.1 Spherical line-element
The spherical line-element adopted in this chapter is, in spherical polar coordinates xµ ≡ t, r, θ, φ,
ds2 = − dt2
α2+
1
β2r
(
dr − βtdt
α
)2
+ r2do2 . (16.1)
Here r is the circumferential radius, defined such that the circumference around any great circle is 2πr. The
line-element (16.1) is somewhat unconventional in that it is not diagonal: gtr does not vanish. There are
two good reasons to consider a non-diagonal metric. Firstly, as discussed in §16.1.11, Einstein’s equations
take a more insightful form when expressed in a non-diagonal frame where βt does not vanish. Secondly, if a
horizon is present, as in the case of black holes, and if the radial coordinate is taken to be the circumferential
radius r, then a diagonal metric will have a coordinate singularity at the horizon, which is not ideal.
266 General spherically symmetric spacetime
The vierbein emµ and inverse vierbein em
µ corresponding to the spherical line-element (16.1) are
emµ =
α βt 0 0
0 βr 0 0
0 0 1/r 0
0 0 0 1/(r sin θ)
, emµ =
1/α 0 0 0
− βt/(αβr) 1/βr 0 0
0 0 r 0
0 0 0 r sin θ
. (16.2)
As in the ADM formalism, Chapter 13, the tetrad time axis γγt is chosen to be orthogonal to hypersurfaces of
constant time t. However, the convention here for the vierbein coefficients differs from the ADM convention:
here 1/α is the ADM lapse, while βt/α is the ADM shift. The directed derivatives ∂t and ∂r along the time
and radial tetrad axes γγt and γγr are
∂t = etµ ∂
∂xµ= α
∂
∂t+ βt
∂
∂r, ∂r = er
µ ∂
∂xµ= βr
∂
∂r. (16.3)
The tetrad-frame 4-velocity um of a person at rest in the tetrad frame is by definition um = 1, 0, 0, 0. It
follows that the coordinate 4-velocity uµ of such a person is
uµ = emµum = et
µ = α, βt, 0, 0 . (16.4)
A person instantaneously at rest in the tetrad frame satisfies dr/dt = βt/α according to equation (16.4),
so it follows from the line-element (16.1) that the proper time τ of a person at rest in the tetrad frame is
related to the coordinate time t by
dτ =dt
αin tetrad rest frame . (16.5)
The directed time derivative ∂t is just the proper time derivative along the worldline of a person continuously
at rest in the tetrad frame (and who is therefore not in free-fall, but accelerating with the tetrad frame),
which follows from
d
dτ=dxµ
dτ
∂
∂xµ= uµ ∂
∂xµ= um∂m = ∂t . (16.6)
By contrast, the proper time derivative measured by a person who is instantaneously at rest in the tetrad
frame, but is in free-fall, is the covariant time derivative
D
Dτ=dxµ
dτDµ = uµDµ = umDm = Dt . (16.7)
Since the coordinate radius r has been defined to be the circumferential radius, a gauge-invariant definition,
it follows that the tetrad-frame gradient ∂m of the coordinate radius r is a tetrad-frame 4-vector (a coordinate
gauge-invariant object)
∂mr = emµ ∂r
∂xµ= em
r = βm = βt, βr, 0, 0 is a tetrad 4-vector . (16.8)
This accounts for the notation βt and βr introduced above. Since βm is a tetrad 4-vector, its scalar product
16.1 Spherical spacetime 267
with itself must be a scalar. This scalar defines the interior mass M(t, r), also called the Misner-Sharp
mass, by
1− 2M
r≡ − βmβ
m = − β2t + β2
r is a coordinate and tetrad scalar . (16.9)
The interpretation of M as the interior mass will become evident below, §16.1.8.
16.1.2 Rest diagonal line-element
Although this is not the choice adopted here, the line-element (16.1) can always be brought to diagonal form
by a coordinate transformation t→ t× (subscripted × for diagonal) of the time coordinate. The t–r part of
the metric is
gtt dt2 + 2 gtr dt dr + grr dr
2 =1
gtt
[
(gtt dt+ gtr dr)2 + (gttgrr − g2
tr) dr2]
. (16.10)
This can be diagonalized by choosing the time coordinate t× such that
f dt× = gtt dt+ gtr dr (16.11)
for some integrating factor f(t, r). Equation (16.11) can be solved by choosing t× to be constant along
integral curves
dr
dt= − gtt
gtr. (16.12)
The resulting diagonal rest line-element is
ds2 = − dt2×α2×
+dr2
1− 2M/r+ r2do2 . (16.13)
The line-element (16.13) corresponds physically to the case where the tetrad frame is taken to be at rest in
the spatial coordinates, βt = 0, as can be seen by comparing it to the earlier line-element (16.1). In changing
the tetrad frame from one moving at dr/dt = βt/α to one that is at rest (at constant circumferential radius
r), a tetrad transformation has in effect been done at the same time as the coordinate transformation (16.11),
the tetrad transformation being precisely that needed to make the line-element (16.13) diagonal. The metric
coefficient grr in the line-element (16.13) follows from the fact that β2r = 1 − 2M/r when βt = 0, equa-
tion (16.9). The transformed time coordinate t× is unspecified up to a transformation t× → f(t×). If the
spacetime is asymptotically flat at infinity, then a natural way to fix the transformation is to choose t× to
be the proper time at rest at infinity.
268 General spherically symmetric spacetime
16.1.3 Comoving diagonal line-element
Although once again this is not the path followed here, the line-element (16.1) can also be brought to diagonal
form by a coordinate transformation r → r×, where, analogously to equation (16.11), r× is chosen to satisfy
f dr× = gtr dt+ grr dr ≡1
βr
(
dr − βtdt
α
)
(16.14)
for some integrating factor f(t, r). The new coordinate r× is constant along the worldline of an object at
rest in the tetrad frame, with dr/dt = βt/α, equation (16.4), so r× can be regarded as a comoving radial
coordinate. The comoving radial coordinate r× could for example be chosen to equal the circumferential
radius r at some fixed instant of coordinate time t (say t = 0). The diagonal comoving line-element in
this comoving coordinate system takes the form
ds2 = − dt2
α2+dr2×λ2
+ r2do2 , (16.15)
where the circumferential radius r(t, r×) is considered to be an implicit function of time t and the comoving
radial coordinate r×. Whereas in the rest line-element (16.13) the tetrad was changed from one that was
moving at dr/dt = βt/α to one that was at rest, here the transformation keeps the tetrad unchanged. In
both the rest and comoving diagonal line-elements (16.13) and (16.15) the tetrad is at rest relative to the
respective radial coordinate r or r×; but whereas in the rest line-element (16.13) the radial coordinate was
fixed to be the circumferential radius r, in the comoving line-element (16.15) the comoving radial coordinate
r× is a label that follows the tetrad. Because the tetrad is unchanged by the transformation to the comoving
radial coordinate r×, the directed time and radial derivatives are unchanged:
∂t = α∂
∂t
∣
∣
∣
∣
r×
= α∂
∂t
∣
∣
∣
∣
r
+ βt∂
∂r
∣
∣
∣
∣
t
, ∂r = λ∂
∂r×
∣
∣
∣
∣
t
= βr∂
∂r
∣
∣
∣
∣
t
. (16.16)
16.1.4 Tetrad connections
Now turn the handle to proceed towards the Einstein equations. The tetrad connections coefficients Γkmn
corresponding to the spherical line-element (16.1) are
Γrtt = ht , (16.17a)
Γrtr = hr , (16.17b)
Γθtθ = Γφtφ =βt
r, (16.17c)
Γθrθ = Γφrφ =βr
r, (16.17d)
Γφθφ =cot θ
r, (16.17e)
16.1 Spherical spacetime 269
where ht is the proper radial acceleration (minus the gravitational force) experienced by a person at rest in
the tetrad frame
ht ≡ − ∂r lnα , (16.18)
and hr is the “Hubble parameter” of the radial flow, as measured in the tetrad rest frame, defined by
hr ≡ − βt∂ lnα
∂r+∂βt
∂r− ∂t lnβr . (16.19)
The interpretation of ht as a proper acceleration and hr as a radial Hubble parameter goes as follows. The
tetrad-frame 4-velocity um of a person at rest in the tetrad frame is by definition um = 1, 0, 0, 0. If the
person at rest were in free fall, then the proper acceleration would be zero, but because this is a general
spherical spacetime, the tetrad frame is not necessarily in free fall. The proper acceleration experienced by
a person continuously at rest in the tetrad frame is the proper time derivative Dum/Dτ of the 4-velocity,
which is
Dum
Dτ= Dtu
m = ∂tum + Γm
ttut = Γm
tt = 0,Γrtt, 0, 0 = 0, ht, 0, 0 , (16.20)
the first step of which follows from equation (16.7). Similarly, a person at rest in the tetrad frame will
measure the 4-velocity of an adjacent person at rest in the tetrad frame a small proper radial distance δξr
away to differ by δξrDrum. The Hubble parameter of the radial flow is thus the covariant radial derivative
Drum, which is
Drum = ∂ru
m + Γmtru
t = Γmtr = 0,Γr
tr, 0, 0 = 0, hr, 0, 0 . (16.21)
Confined to the t–r-plane (that is, considering only Lorentz transformations in the t–r-plane, which is to
say radial Lorentz boosts), the acceleration ht and Hubble parameter hr constitute the components of a
tetrad-frame 2-vector hn = ht, hr:
hn = Γrtn . (16.22)
The Riemann tensor, equations (16.24) below, involves covariant derivatives Dmhn of hn, which coincide
with the covariant derivatives D(2)hn confined to t–r-plane.
Since hr is a kind of radial Hubble parameter, it can be useful to define a corresponding radial scale factor
λ by
hr ≡ −∂t lnλ . (16.23)
The scale factor λ is the same as the λ in the comoving line-element of equation (16.15). This is true because
hr is a tetrad connection and therefore coordinate gauge-invariant, and the line-element (16.15) is related
to the line-element (16.1) being considered by a coordinate transformation r → r× that leaves the tetrad
unchanged.
270 General spherically symmetric spacetime
16.1.5 Riemann, Einstein, and Weyl tensors
The non-vanishing components of the tetrad-frame Riemann tensor Rklmn corresponding to the spherical
line-element (16.1) are
Rtrtr = Drht −Dthr , (16.24a)
Rtθtθ = Rtφtφ = − 1
rDtβt , (16.24b)
Rrθrθ = Rrφrφ = − 1
rDrβr , (16.24c)
Rtθrθ = Rtφrφ = − 1
rDtβr = − 1
rDrβt , (16.24d)
Rθφθφ =2M
r3, (16.24e)
where Dm denotes the covariant derivative as usual.
The non-vanishing components of the tetrad-frame Einstein tensor Gkm are
Gtt = 2Rrθrθ +Rθφθφ , (16.25a)
Grr = 2Rtθtθ −Rθφθφ , (16.25b)
Gtr = − 2Rtθrθ , (16.25c)
Gθθ = Gφφ = Rtrtr +Rtθtθ −Rrθrθ , (16.25d)
whence
Gtt =2
r
(
−Drβr +M
r2
)
, (16.26a)
Grr =2
r
(
−Dtβt −M
r2
)
, (16.26b)
Gtr =2
rDtβr =
2
rDrβt , (16.26c)
Gθθ = Gφφ = Drht −Dthr +1
r(Drβr −Dtβt) . (16.26d)
The non-vanishing components of the tetrad-frame Weyl tensor Cklmn are
12 Ctrtr = −Ctθtθ = −Ctφtφ = Crθrθ = Crφrφ = − 1
2 Cθφθφ = C , (16.27)
where C is the Weyl scalar (the spin-0 component of the Weyl tensor),
C ≡ 1
6(Rtrtr −Rtθtθ +Rrθrθ −Rθφθφ) =
1
6
(
Gtt −Grr +Gθθ)
− M
r3. (16.28)
16.1.6 Einstein equations
The tetrad-frame Einstein equations
Gkm = 8πT km (16.29)
16.1 Spherical spacetime 271
imply that
Gtt Gtr 0 0
Gtr Grr 0 0
0 0 Gθθ 0
0 0 0 Gφφ
= 8πT km = 8π
ρ f 0 0
f p 0 0
0 0 p⊥ 0
0 0 0 p⊥
(16.30)
where ρ ≡ T tt is the proper energy density, f ≡ T tr is the proper radial energy flux, p ≡ T rr is the proper
radial pressure, and p⊥ ≡ T θθ = T φφ is the proper transverse pressure.
16.1.7 Choose your frame
So far the radial motion of the tetrad frame has been left unspecified. Any arbitrary choice can be made.
For example, the tetrad frame could be chosen to be at rest,
βt = 0 , (16.31)
as in the Schwarzschild or Reissner-Nordstrom line-elements. Alternatively, the tetrad frame could be chosen
to be in free-fall,
ht = 0 , (16.32)
as in the Gullstrand-Painleve line-element. For situations where the spacetime contains matter, one natural
choice is the center-of-mass frame, defined to be the frame in which the energy flux f is zero
Gtr = 8πf = 0 . (16.33)
Whatever the choice of radial tetrad frame, tetrad-frame quantities in different radial tetrad frames are
related to each other by a radial Lorentz boost.
16.1.8 Interior mass
Equations (16.26a) with the middle expression of (16.26c), and (16.26b) with the final expression of (16.26c),
respectively, along with the definition (16.9) of the interior mass M , and the Einstein equations (16.30),
imply
p =1
βt
(
− 1
4πr2∂tM − βrf
)
, (16.34a)
ρ =1
βr
(
1
4πr2∂rM − βtf
)
. (16.34b)
In the center-of-mass frame, f = 0, these equations reduce to
∂tM = − 4πr2βt p , (16.35a)
∂rM = 4πr2βr ρ . (16.35b)
272 General spherically symmetric spacetime
Equations (16.35) amply justify the interpretation of M as the interior mass. The first equation (16.35a)
can be written
∂tM + p 4πr2∂tr = 0 , (16.36)
which can be recognized as an expression of the first law of thermodynamics,
dE + p dV = 0 , (16.37)
with mass-energy E equal to M . The second equation (16.35b) can be written, since ∂r = βr ∂/∂r, equa-
tion (16.3),
∂M
∂r= 4πr2ρ , (16.38)
which looks exactly like the Newtonian relation between interior mass M and density ρ. Actually, this
apparently Newtonian equation (16.38) is deceiving. The proper 3-volume element d3r in the center-of-mass
tetrad frame is given by
d3r γγr ∧ γγθ ∧ γγφ = gr dr ∧ gθ dθ ∧ gφ dφ =r2 sin θ dr dθ dφ
βrγγr ∧ γγθ ∧ γγφ , (16.39)
so that the proper 3-volume element dV ≡ d3r of a radial shell of width dr is
dV =4πr2dr
βr. (16.40)
Thus the “true” mass-energy dMm associated with the proper density ρ in a proper radial volume element
dV might be expected to be
dMm = ρ dV =4πr2dr
βr, (16.41)
whereas equation (16.38) indicates that the actual mass-energy is
dM = ρ 4πr2dr = βr ρ dV . (16.42)
A person in the center-of-mass frame might perhaps, although there is really no formal justification for doing
so, interpret the balance of the mass-energy as gravitational mass-energy Mg
dMg = (βr − 1)ρ dV . (16.43)
Whatever the case, the moral of this is that you should beware of interpreting the interior mass M too
literally as palpable mass-energy.
16.1.9 Energy-momentum conservation
Covariant conservation of the Einstein tensor DmGmn = 0 implies conservation of energy-momentum
DmTmn = 0. The two non-vanishing equations represent conservation of energy and of radial momentum,
16.1 Spherical spacetime 273
and are
DmTmt = ∂tρ+
2βt
r(ρ+ p⊥) + hr (ρ+ p) +
(
∂r +2βr
r+ 2 ht
)
f = 0 , (16.44a)
DmTmr = ∂rp+
2βr
r(p− p⊥) + ht (ρ+ p) +
(
∂t +2βt
r+ 2 hr
)
f = 0 . (16.44b)
In the center-of-mass frame, f = 0, these energy-momentum conservation equations reduce to
∂tρ+2βt
r(ρ+ p⊥) + hr (ρ+ p) = 0 , (16.45a)
∂rp+2βr
r(p− p⊥) + ht (ρ+ p) = 0 . (16.45b)
In a general situation where the mass-energy is the sum over several individual components a,
Tmn =∑
species a
Tmna , (16.46)
the individual mass-energy components a of the system each satisfy an energy-momentum conservation
equation of the form
DmTmna = Fn
a , (16.47)
where Fna is the flux of energy into component a. Einstein’s equations enforce energy-momentum conservation
of the system as a whole, so the sum of the energy fluxes must be zero∑
species a
Fna = 0 . (16.48)
16.1.10 First law of thermodynamics
For an individual species a, the energy conservation equation (16.44a) in the center-of-mass frame of the
species, fa = 0, can be written
DmTmta = ∂tρa + (ρa + p⊥a)∂t ln r2 + (ρa + pa)∂t lnλa = F t
a , (16.49)
where λa is the radial “scale factor,” equation (16.23), in the center-of-mass frame of the species (the scale
factor is different in different frames). Equation (16.49) can be recognized as an expression of the first law
of thermodynamics for a volume element V of species a, in the form
V −1[
∂t(ρaV ) + p⊥a Vr ∂tV⊥ + pa V⊥ ∂tVr
]
= F ta , (16.50)
with transverse volume (area) V⊥ ∝ r2, radial volume (width) Vr ∝ λa, and total volume V ∝ V⊥Vr . The
flux F ta on the right hand side is the heat per unit volume per unit time going into species a. If the pressure
of species a is isotropic, p⊥a = pa, then equation (16.50) simplifies to
V −1[
∂t(ρaV ) + pa ∂tV]
= F ta , (16.51)
with volume V ∝ r2λa.
274 General spherically symmetric spacetime
16.1.11 Structure of the Einstein equations
The spherically symmetric spacetime under consideration is described by 3 vierbein coefficients, α, βt, and βr.
However, some combination of the 3 coefficients represents a gauge freedom, since the spherically symmetric
spacetime has only two physical degrees of freedom. As commented in §16.1.7, various gauge-fixing choices
can be made, such as choosing to work in the center-of-mass frame, f = 0.
Equations (16.26) give 4 equations for the 4 non-vanishing components of the Einstein tensor. The two
expressions for Gtr are identical when expressed in terms of the vierbein and vierbein derivatives, so are not
distinct equations. Conservation of energy-momentum of the system as a whole is built in to the Einstein
equations, a consequence of the Bianchi identities, so 2 of the Einstein equations are effectively equivalent to
the energy-momentum conservation equations (16.44). In the general case where the matter contains multiple
components, it is usually a good idea to include the equations describing the conservation or exchange of
energy-momentum separately for each component, so that global conservation of energy-momentum is then
satisfied as a consequence of the matter equations.
This leaves 2 independent Einstein equations to describe the 2 physical degrees of freedom of the spacetime.
The 2 equations may be taken to be the evolution equations (16.26c) and (16.26b) for βt and βr,
Dtβt = −Mr2− 4πrp , (16.52a)
Dtβr = 4πrf , (16.52b)
which are valid for any choice of tetrad frame, not just the center-of-mass frame.
Equations (16.52) are the most important of the general relativistic equations governing spherically sym-
metric spacetimes. It is these equations that are responsible (to the extent that equations may be con-
sidered responsible) for the strange internal structure of Reissner-Nordstrom black holes, and for mass
inflation. The coefficient βt equals the coordinate radial 4-velocity dr/dτ = ∂tr = βt of the tetrad
frame, equation (16.4), and thus equation (16.52a) can be regarded as giving the proper radial acceleration
D2r/Dτ2 = Dβt/Dτ = Dtβt of the tetrad frame as measured by a person who is in free-fall and instanta-
neously at rest in the tetrad frame. If the acceleration is measured by an observer who is continuously at
rest in the tetrad frame (as opposed to being in free-fall), then the proper acceleration is ∂tβt = Dtβt +βrht.
The presence of the extra term βrht, proportional to the proper acceleration ht actually experienced by
the observer continuously at rest in the tetrad frame, reflects the principle of equivalence of gravity and
acceleration.
The right hand side of equation (16.52a) can be interpreted as the radial gravitational force, which consists
of two terms. The first term, −M/r2, looks like the familiar Newtonian gravitational force, which is attractive
(negative, inward) in the usual case of positive mass M . The second term, −4πrp, proportional to the radial
pressure p, is what makes spherical spacetimes in general relativity interesting. In a Reissner-Nordstrom black
hole, the negative radial pressure produced by the radial electric field produces a radial gravitational repulsion
(positive, outward), according to equation (16.52a), and this repulsion dominates the gravitational force at
small radii, producing an inner horizon. In mass inflation, the (positive) radial pressure of relativistically
16.1 Spherical spacetime 275
counter-streaming ingoing and outgoing streams just above the inner horizon dominates the gravitational
force (inward), and it is this that drives mass inflation.
Like the second half of a vaudeville act, the second Einstein equation (16.52b) also plays an indispensible
role. The quantity βr ≡ ∂rr on the left hand side is the proper radial gradient of the circumferential radius
r measured by a person at rest in the tetrad frame. The sign of βr determines which way an observer at
rest in the tetrad frame thinks is “outwards”, the direction of larger circumferential radius r. A positive
βr means that the observer thinks the outward direction points away from the black hole, while a negative
βr means that the observer thinks the outward direction points towards from the black hole. Outside the
outer horizon βr is necessarily positive, because βm must be spacelike there. But inside the horizon βr may
be either positive or negative. A tetrad frame can be defined as “ingoing” if the proper radial gradient βr
is positive, and “outgoing” if βr is negative. In the Reissner-Nordstrom geometry, ingoing geodesics have
positive energy, and outgoing geodesics have negative energy. However, the present definition of ingoing or
outgoing based on the sign of βr is general – there is no need for a timelike Killing vector such as would be
necessary to define the (conserved) energy of a geodesic.
Equation (16.52b) shows that the proper rate of change Dtβr in the radial gradient βr measured by an
observer who is in free-fall and instantaneously at rest in the tetrad frame is proportional to the radial energy
flux f in that frame. But ingoing observers tend to see energy flux pointing away from the black hole, while
outgoing observers tend to see energy flux pointing towards the black hole. Thus the change in βr tends to
be in the same direction as βr, amplifying βr whatever its sign.
Exercise 16.1 Birkhoff’s theorem. Prove Birkhoff’s theorem from equations (16.52).
16.1.12 Comment on the vierbein coefficient α
Whereas the Einstein equations (16.52) give evolution equations for the vierbein coefficients βt and βr, there
is no evolution equation for the vierbein coefficient α. Indeed, the Einstein equations involve the vierbein
coefficient α only in the combination ht ≡ −∂r lnα. This reflects the fact that, even after the tetrad frame
is fixed, there is still a coordinate freedom t → t′(t) in the choice of coordinate time t. Under such a gauge
transformation, α transforms as α→ α′ = f(t)α where f(t) = ∂t′/∂t is an arbitrary function of coordinate
time t. Only ht ≡ −∂r lnα is independent of this coordinate gauge freedom, and thus only ht, not α itself,
appears in the tetrad-frame Einstein equations.
Since α is needed to propagate the equations from one coordinate time to the next (because ∂t = α∂/∂t+
βt ∂/∂r), it is necessary to construct α by integrating ht ≡ −βr ∂ lnα/∂r along the radial direction r at each
time step. The arbitrary normalization of α at each step might be fixed by choosing α to be unity at infinity,
which corresponds to fixing the time coordinate t to equal the proper time at infinity.
In the particular case that the tetrad frame is taken to be in free-fall everywhere, ht = 0, as in the
Gullstrand-Painleve line-element, then α is constant at fixed t, and without loss of generality it can be fixed
equal to unity everywhere, α = 1. I like to think of a free-fall frame as being realized physically by tracer
276 General spherically symmetric spacetime
“dark matter” particles that free-fall radially (from zero velocity, typically) at infinity, and stream freely,
without interacting, through any actual matter that may be present.
16.2 Spherical electromagnetic field
The internal structure of a charged black hole resembles that of a rotating black hole because the negative
pressure (tension) of the radial electric field produces a gravitational repulsion analogous to the centrifugal
repulsion in a rotating black hole. Since it is much easier to deal with spherical than rotating black holes, it
is common to use charge as a surrogate for rotation in exploring black holes.
16.2.1 Electromagnetic field
The assumption of spherical symmetry means that any electromagnetic field can consist only of a radial elec-
tric field (in the absence of magnetic monopoles). The only non-vanishing components of the electromagnetic
field Fmn are then
Ftr = −Frt = E =Q
r2, (16.53)
where E is the radial electric field, and Q(t, r) is the interior electric charge. Equation (16.53) can be
regarded as defining what is meant by the electric charge Q interior to radius r at time t.
16.2.2 Maxwell’s equations
A radial electric field automatically satisfies two of Maxwell’s equations, the source-free ones (11.64). For
the radial electric field (16.53), the other two Maxwell’s equations, the sourced ones (11.65), are
∂rQ = 4πr2q , (16.54a)
∂tQ = −4πr2j , (16.54b)
where q ≡ jt is the proper electric charge density and j ≡ jr is the proper radial electric current density in
the tetrad frame.
16.2.3 Electromagnetic energy-momentum tensor
For the radial electric field (16.53), the electromagnetic energy-momentum tensor (11.70) in the tetrad frame
is the diagonal tensor
Tmne =
Q2
8πr4
1 0 0 0
0 −1 0 0
0 0 1 0
0 0 0 1
. (16.55)
16.3 General relativistic stellar structure 277
The radial electric energy-momentum tensor is independent of the radial motion of the tetrad frame, which
reflects the fact that the electric field is invariant under a radial Lorentz boost. The energy density ρe and
radial and transverse pressures pe and p⊥e of the electromagnetic field are the same as those from a spherical
charge distribution with interior electric charge Q in flat space
ρe = −pe = p⊥e =Q2
8πr4=E2
8π. (16.56)
The non-vanishing components of the covariant derivative DmTmne of the electromagnetic energy-mom-
entum (16.55) are
DmTmte = ∂tρe +
4βt
rρe =
Q
4πr4∂tQ = − jQ
r2= − jE , (16.57a)
DmTmre = ∂rpe +
4βr
rpe = − Q
4πr4∂rQ = − qQ
r2= − qE . (16.57b)
The first expression (16.57a), which gives the rate of energy transfer out of the electromagnetic field as the
current density j times the electric field E, is the same as in flat space. The second expression (16.57b),
which gives the rate of transfer of radial momentum out of the electromagnetic field as the charge density q
times the electric field E, is the Lorentz force on a charge density q, and again is the same as in flat space.
16.3 General relativistic stellar structure
A star can be well approximated as static as well as spherically symmetric. In this case all time derivatives
can be taken to vanish, ∂/∂t = 0, and, since the center-of-mass frame coincides with the rest frame, it is
natural to choose the tetrad frame to be at rest, βt = 0. The Einstein equation (16.52b) then vanishes
identically, while the Einstein equation (16.52a) becomes
βrht =M
r2+ 4πrp , (16.58)
which expresses the proper acceleration ht in the rest frame in terms of the familiar Newtonian gravitational
force M/r2 plus a term 4πrp proportional to the radial pressure. The radial pressure p, if positive as is the
usual case for a star, enhances the inward gravitational force, helping to destabilize the star. Because βt is
zero, the interior mass M given by equation (16.9) reduces to
1− 2M/r = β2r . (16.59)
When equations (16.58) and (16.59) are substituted into the momentum equation (16.44b), and if the pressure
is taken to be isotropic, so p⊥ = p, the result is the Oppenheimer-Volkov equation for general relativistic
hydrostatic equilibrium
∂p
∂r= − (ρ+ p)(M + 4πr3p)
r2(1− 2M/r). (16.60)
278 General spherically symmetric spacetime
In the Newtonian limit p≪ ρ and M ≪ r this goes over to (with units restored)
∂p
∂r= − ρ GM
r2, (16.61)
which is the usual Newtonian equation of spherically symmetric hydrostatic equilibrium.
16.4 Self-similar spherically symmetric spacetime
Even with the assumption of spherical symmetry, it is by no means easy to solve the system of partial
differential equations that comprise the Einstein equations coupled to mass-energy of various kinds. One
way to simplify the system of equations, transforming them into ordinary differential equations, is to consider
self-similar solutions.
16.4.1 Self-similarity
The assumption of self-similarity (also known as homothety, if you can pronounce it) is the assumption
that the system possesses conformal time translation invariance. This implies that there exists a conformal
time coordinate η such that the geometry at any one time is conformally related to the geometry at any
other time
ds2 = a(η)2[
g(c)ηη (x) dη2 + 2 g(c)
ηx (x) dη dx+ g(c)xx (x) dx2 + e2x do2
]
. (16.62)
Here the conformal metric coefficients g(c)µν (x) are functions only of conformal radius x, not of conformal time
η. The choice e2x of coefficient of do2 is a gauge choice of the conformal radius x, carefully chosen here so as
to bring the self-similar line-element into a form (16.66) below that resembles as far as possible the spherical
line-element (16.1). In place of the conformal factor a(η) it is convenient to work with the circumferential
radius r
r ≡ a(η)ex (16.63)
which is to be considered as a function r(η, x) of the coordinates η and x. The circumferential radius r has
a gauge-invariant meaning, whereas neither a(η) nor x are independently gauge-invariant. The conformal
factor r has the dimensions of length. In self-similar solutions, all quantities are proportional to some power of
r, and that power can be determined by dimensional analysis. Quantites that depend only on the conformal
radial coordinate x, independent of the circumferential radius r, are called dimensionless.
The fact that dimensionless quantities such as the conformal metric coefficients g(c)µν (x) are independent of
conformal time η implies that the tangent vector gη, which by definition satisfies
∂
∂η= gη · ∂ , (16.64)
is a conformal Killing vector, also known as the homothetic vector. The tetrad-frame components of the
16.4 Self-similar spherically symmetric spacetime 279
conformal Killing vector gη defines the tetrad-frame conformal Killing 4-vector ξm
∂
∂η≡ r ξm∂m , (16.65)
in which the factor r is introduced so as to make ξm dimensionless. The conformal Killing vector gη is
the generator of the conformal time translation symmetry, and as such it is gauge-invariant (up to a global
rescaling of conformal time, η → bη for some constant b). It follows that its dimensionless tetrad-frame
components ξm constitute a tetrad 4-vector (again, up to global rescaling of conformal time).
16.4.2 Self-similar line-element
The self-similar line-element can be taken to have the same form as the spherical line-element (16.1), but
with the dependence on the dimensionless conformal Killing vector ξm made manifest:
ds2 = − r2[
(ξη dη)2 +1
β2x
(dx + βx ξxdη)
2+ do2
]
. (16.66)
The vierbein emµ and inverse vierbein em
µ corresponding to the self-similar line-element (16.66) are
emµ =
1
r
1/ξη − βx ξx/ξη 0 0
0 βx 0 0
0 0 1 0
0 0 0 1/ sin θ
, emµ = r
ξη 0 0 0
ξx 1/βx 0 0
0 0 1 0
0 0 0 sin θ
. (16.67)
It is straightforward to see that the coordinate time components of the inverse vierbein must be emη = r ξm,
since ∂/∂η = emη ∂m equals r ξm∂m, equation (16.65).
16.4.3 Tetrad-frame scalars and vectors
Since the conformal factor r is gauge-invariant, the directed gradient ∂mr constitutes a tetrad-frame 4-vector
βm (which unlike ξm is independent of any global rescaling of conformal time)
βm ≡ ∂mr . (16.68)
It is straightforward to check that βx defined by equation (16.68) is consistent with its appearance in the
vierbein (16.67) provided that r ∝ ex as earlier assumed, equation (16.63).
With two distinct dimensionless tetrad 4-vectors in hand, βm and the conformal Killing vector ξm, three
gauge-invariant dimensionless scalars can be constructed, βmβm, ξmβm, and ξmξm,
1− 2M
r= − βmβm = − β2
η + β2x , (16.69)
v ≡ ξmβm =1
r
∂r
∂η=
1
a
∂a
∂η, (16.70)
280 General spherically symmetric spacetime
∆ ≡ − ξmξm = (ξη)2 − (ξx)2 . (16.71)
Equation (16.69) is essentially the same as equation (16.9). The dimensionless quantity v, equation (16.70),
may be interpreted as a measure of the expansion velocity of the self-similar spacetime. Equation (16.70)
shows that v is a function only of η (since a(η) is a function only of η), and it therefore follows that v must
be constant (since being dimensionless means that v must be a function only of x). Equation (16.70) then
also implies that the conformal factor a(η) must take the form
a(η) = evη . (16.72)
Because of the freedom of a global rescaling of conformal time, it is possible to set v = 1 without loss of
generality, but in practice it is convenient to keep v, because it is then transparent how to take the static
limit v → 0. Equation (16.72) along with (16.63) shows that the circumferential radius r is related to the
conformal coordinates η and x by
r = evη+x . (16.73)
The dimensionsless quantity ∆, equation (16.71), is the dimensionless horizon function: horizons occur where
the horizon function vanishes
∆ = 0 at horizons . (16.74)
16.4.4 Self-similar diagonal line-element
The self-similar line-element (16.66) can be brought to diagonal form by a coordinate transformation to
diagonal conformal coordinates η×, x× (subscripted × for diagonal)
η → η× = η + f(x) , x→ x× = x− vf(x) , (16.75)
which leaves unchanged the conformal factor r, equation (16.73). The resulting diagonal metric is
ds2 = r2(
−∆ dη2× +
dx2×
1− 2M/r + v2/∆+ do2
)
. (16.76)
The diagonal line-element (16.76) corresponds physically to the case where the tetrad frame is at rest in
the similarity frame, ξx = 0, as can be seen by comparing it to the line-element (16.66). The frame can be
called the similarity frame. The form of the metric coefficients follows from the line-element (16.66) and
the gauge-invariant scalars (16.69)–(16.71).
The conformal Killing vector in the similarity frame is ξm = ∆1/2, 0, 0, 0, and the 4-velocity of the
similarity frame in its own frame is um = 1, 0, 0, 0. Since both are tetrad 4-vectors, it follows that with
respect to a general tetrad frame
ξm = um ∆1/2 (16.77)
where um is the 4-velocity of the similarity frame with respect to the general frame. This shows that the
conformal Killing vector ξm in a general tetrad frame is proportional to the 4-velocity of the similarity frame
16.4 Self-similar spherically symmetric spacetime 281
through the tetrad frame. In particular, the proper 3-velocity of the similarity frame through the tetrad
frame is
proper 3-velocity of similarity frame through tetrad frame =ξx
ξη. (16.78)
16.4.5 Ray-tracing line-element
It proves useful to introduce a “ray-tracing” conformal radial coordinate X related to the coordinate x× of
the diagonal line-element (16.76) by
dX ≡ ∆ dx×
[(1− 2M/r)∆ + v2]1/2
. (16.79)
In terms of the ray-tracing coordinate X , the diagonal metric is
ds2 = r2(
−∆ dη2× +
dX2
∆+ do2
)
. (16.80)
For the Reissner-Nordstrom geometry, ∆ = (1 − 2M/r)/r2, η× = t, and X = −1/r.
16.4.6 Geodesics
Spherical symmetry and conformal time translation symmetry imply that geodesic motion in spherically
symmetric self-similar spacetimes is described by a complete set of integrals of motion.
The integral of motion associated with conformal time translation symmetry can be obtained from La-
grange’s equations of motion
d
dτ
∂L
∂uη=∂L
∂η, (16.81)
with effective Lagrangian L = gµνuµuν for a particle with 4-velocity uµ. The self-similar metric depends on
the conformal time η only through the overall conformal factor gµν ∝ a(η)2. The derivative of the conformal
factor is given by ∂ ln a/∂η = v, equation (16.70), so it follows that ∂L/∂η = 2 vL. For a massive particle,
for which conservation of rest mass implies gµνuµuν = −1, Lagrange’s equations (16.81) thus yield
duη
dτ= − v . (16.82)
In the limit of zero accretion rate, v → 0, equation (16.82) would integrate to give uη as a constant, the
energy per unit mass of the geodesic. But here there is conformal time translation symmetry in place of
time translation symmetry, and equation (16.82) integrates to
uη = − v τ , (16.83)
in which an arbitrary constant of integration has been absorbed into a shift in the zero point of the proper
time τ . Although the above derivation was for a massive particle, it holds also for a massless particle, with the
understanding that the proper time τ is constant along a null geodesic. The quantity uη in equation (16.83)
282 General spherically symmetric spacetime
is the covariant time component of the coordinate-frame 4-velocity uµ of the particle; it is related to the
covariant components um of the tetrad-frame 4-velocity of the particle by
uη = emη um = r ξmum . (16.84)
Without loss of generality, geodesic motion can be taken to lie in the equatorial plane θ = π/2 of the
spherical spacetime. The integrals of motion associated with conformal time translation symmetry, rotational
symmetry about the polar axis, and conservation of rest mass, are, for a massive particle
uη = − v τ , uφ = Lz , uµuµ = −1 , (16.85)
where Lz is the orbital angular momentum per unit rest mass of the particle. The coordinate 4-velocity
uµ ≡ dxµ/dτ that follows from equations (16.85) takes its simplest form in the conformal coordinates
η×, X, θ, φ of the ray-tracing metric (16.80)
uη× =v τ
r2∆, uX = ± 1
r2[
v2τ2 − (r2 + L2z)∆
]1/2, uφ =
Lz
r2. (16.86)
16.4.7 Null geodesics
The important case of a massless particle follows from taking the limit of a massive particle with infinite
energy and angular momentum, v τ →∞ and Lz → ∞. To obtain finite results, define an affine parameter
λ by CHECK dλ ≡ v τ dτ , and a 4-velocity in terms of it by vµ ≡ dxµ/dλ. The integrals of motion (16.85)
then become, for a null geodesic,
vη×= −1 , vφ = Jz , vµv
µ = 0 , (16.87)
where Jz ≡ Lz/(v τ) is the (dimensionless) conformal angular momentum of the particle. The 4-velocity vµ
along the null geodesic is then, in terms of the coordinates of the ray-tracing metric (16.80),
vη =1
r2∆, vX = ± 1
r2(
1− J2z ∆)1/2
, vφ =Jz
r2. (16.88)
Equations (16.88) yield the shape of a null geodesic by quadrature
φ =
∫
Jz dX
(1− J2z ∆)1/2
. (16.89)
Equation (16.89) shows that the shape of null geodesics in spherically symmetric self-similar spacetimes
hinges on the behavior of the dimensionless horizon function ∆(X) as a function of the dimensionless ray-
tracing variable X . Null geodesics go through periapsis or apoapsis in the self-similar frame where the
denominator of the integrand of (16.89) is zero, corresponding to vX = 0.
In the Reissner-Nordstrom geometry there is a radius, the photon sphere, where photons can orbit in
circles for ever. In non-stationary self-similar solutions there is no conformal radius where photons can orbit
for ever (to remain at fixed conformal radius, the photon angular momentum would have to increase in
proportion to the conformal factor r). There is however a separatrix between null geodesics that do or do
not fall into the black hole, and the conformal radius where this occurs can be called the photon sphere
16.4 Self-similar spherically symmetric spacetime 283
equivalent. The photon sphere equivalent occurs where the denominator of the integrand of equation (16.89)
not only vanishes, vX = 0, but is an extremum, which happens where the horizon function ∆ is an extremum,
d∆
dX= 0 at photon sphere equivalent . (16.90)
16.4.8 Dimensional analysis
Dimensional analysis shows that the conformal coordinates xµ ≡ η, x, θ, φ and the tetrad metric γmn are
dimensionless, while the coordinate metric gµν scales as r2,
xµ ∝ r0 , γmn ∝ r0 , gµν ∝ r2 . (16.91)
The vierbein emµ and inverse vierbein em
µ, equations (16.67), scale as
emµ ∝ r−1 , em
µ ∝ r . (16.92)
Coordinate derivatives ∂/∂xµ are dimensionless, while directed derivatives ∂m scale as 1/r,
∂
∂xµ∝ r0 , ∂m ∝ r−1 . (16.93)
The tetrad connections Γkmn and the tetrad-frame Riemann tensor Rklmn scale as
Γkmn ∝ r−1 , Rklmn ∝ r−2 . (16.94)
16.4.9 Variety of self-similar solutions
Self-similar solutions exist provided that the properties of the energy-momentum introduce no additional
dimensional parameters. For example, the pressure-to-density ratio w ≡ p/ρ of any species is dimensionless,
and since the ratio can depend only on the nature of the species itself, not for example on where it happens
to be located in the spacetime, it follows that the ratio w must be a constant. It is legitimate for the
pressure-to-density ratio to be different in the radial and transverse directions (as it is for a radial electric
field), but otherwise self-similarity requires that
w ≡ p/ρ , w⊥ ≡ p⊥/ρ , (16.95)
be constants for each species. For example, w = 1 for an ultrahard fluid (which can mimic the behaviour
of a massless scalar field: E. Babichev, S. Chernov, V. Dokuchaev, Yu. Eroshenko, 2008, “Ultra-hard fluid
and scalar field in the Kerr-Newman metric,” Phys. Rev. D 78, 104027, arXiv:0807.0449), w = 1/3 for a
relativistic fluid, w = 0 for pressureless cold dark matter, w = −1 for vacuum energy, and w = −1 with
w⊥ = 1 for a radial electric field.
Self-similarity allows that the energy-momentum may consist of several distinct components, such as
a relativistic fluid, plus dark matter, plus an electric field. The components may interact with each other
provided that the properties of the interaction introduce no additional dimensional parameters. For example,
the relativistic fluid (and the dark matter) may be charged, and if so then the charged fluid will experience
284 General spherically symmetric spacetime
a Lorentz force from the electric field, and will therefore exchange momentum with the electric field. If
the fluid is non-conducting, then there is no dissipation, and the interaction between the charged fluid and
electric field automatically introduces no additional dimensional parameters.
However, if the charged fluid is electrically conducting, then the electrical conductivity could potentially
introduce an additional dimensional parameter, and this must not be allowed if self-similarity is to be
maintained. In diffusive electrical conduction in a fluid of conductivity σ, an electric field E gives rise to a
current
j = σE , (16.96)
which is just Ohm’s law. Dimensional analysis shows that j ∝ r−2 and E ∝ r−1, so the conductivity must
scale as σ ∝ r−1. The conductivity can depend only on the intrinsic properties of the conducting fluid, and
the only intrinsic property available is its density, which scales as ρ ∝ r−2. If follows that the conductivity
must be proportional to the square root of the density ρ of the conducting fluid
σ = κ ρ1/2 , (16.97)
where κ is a dimensionless conductivity constant. The form (16.97) is required by self-similarity, and is
not necessarily realistic (although it is realistic that the conductivity increases with density). However, the
conductivity (16.97) is adequate for the purpose of exploring the consequences of dissipation in simple models
of black holes.
16.4.10 Tetrad connections
The expressions for the tetrad connections for the self-similar spacetime are the same as those (16.17) for
a general spherically symmetric spacetime, with just a relabelling of the time and radial coordinates into
conformal coordinates
t→ η , r → x . (16.98)
Specifically, equations (16.17) for the tetrad connections become become
Γηxη = hη , Γηxx = hx , Γηθθ = Γηφφ =βη
r, Γxθθ = Γxφφ =
βx
r, Γθφφ =
cot θ
r, (16.99)
in which hη and hx have the same physical interpretation discussed in §16.1.4 for the general spherically sym-
metric case: hη is the proper radial acceleration, and hx is the radial Hubble parameter. Expressions (16.18)
and (16.19) for hη and hx translate in the self-similar spacetime to
hη ≡ ∂x ln(r ξη) , hx ≡ ∂η ln(r ξx) . (16.100)
Comparing equations (16.100) to equations (16.18) and (16.23) shows that the vierbein coefficient α and
scale factor λ translate in the self-similar spacetime to
α =1
rξη, λ =
1
rξx. (16.101)
16.4 Self-similar spherically symmetric spacetime 285
16.4.11 Spherical equations carry over to the self-similar case
The tetrad-frame Riemann, Weyl, and Einstein tensors in the self-similar spacetime take the same form
as in the general spherical case, equations (16.24)–(16.26), with just a relabelling (16.98) into conformal
coordinates.
Likewise, the equations for the interior mass in §16.1.8, for energy-momentum conservation in §16.1.9, for
the first law in §16.1.10, and the various equations for the electromagnetic field in §16.2, all carry through
unchanged except for a relabelling (16.98) of coordinates.
16.4.12 From partial to ordinary differential equations
The central simplifying feature of self-similar solutions is that they turn a system of partial differential
equations into a system of ordinary differential equations.
By definition, a dimensionless quantity F (x) is independent of conformal time η. It follows that the partial
derivative of any dimensionless quantity F (x) with respect to conformal time η vanishes
0 =∂F (x)
∂η= ξm∂mF (x) = (ξη∂η + ξx∂x)F (x) . (16.102)
Consequently the directed radial derivative ∂xF of a dimensionless quantity F (x) is related to its directed
time derivative ∂ηF by
∂xF (x) = − ξx
ξη∂ηF (x) . (16.103)
Equation (16.103) allows radial derivatives to be converted to time derivatives.
16.4.13 Integrals of motion
As remarked above, equation (16.102), in self-similar solutions ξm∂mF (x) = 0 for any dimensionless function
F (x). If both the directed derivatives ∂ηF (x) and ∂xF (x) are known from the Einstein equations or elsewhere,
then the result will be an integral of motion.
The spherically symmetric, self-similar Einstein equations admit two integrals of motion
0 = r ξm∂mβη =r βx(ξηhη + ξxhx)− ξη
(
M
r+ 4πr2p
)
+ ξx4πr2f , (16.104a)
0 = r ξm∂mβx =r βη(ξηhη + ξxhx) + ξx
(
M
r− 4πr2ρ
)
+ ξη4πr2f . (16.104b)
Taking ξx times (16.104a) plus ξη times (16.104b), and then βη times (16.104a) minus βx times (16.104b),
gives
0 = r v (ξηhη + ξxhx)− 4πr2[
ξηξx(ρ+ p)−(
(ξη)2 + (ξx)2)
f]
, (16.105a)
0 = r ξm∂mM
r= − v
M
r+ 4πr2 [βxξ
xρ− βηξηp+ (βηξ
x − βxξη)f ] . (16.105b)
286 General spherically symmetric spacetime
The quantities in square brackets on the right hand sides of equations (16.105) are scalars for each species
a, so equations (16.105) can also be written
r v (ξηhη + ξxhx) = 4πr2∑
species a
ξηaξ
xa (ρa + pa) , (16.106a)
vM
r= 4πr2
∑
species a
(βa,xξxaρa − βa,ηξ
ηapa) , (16.106b)
where the sum is over all species a, and βa,m and ξma are the 4-vectors βm and ξm expressed in the rest frame
of species a.
For electrically charged solutions, a third integral of motion comes from
0 = r ξm∂mQ
r= − v
Q
r+ 4πr2 (ξxq − ξηj) (16.107)
which is valid in any radial tetrad frame, not just the center-of-mass frame.
For a fluid with equation of state p = wρ, a fourth integral comes from considering
0 = r ξm∂m(r2p) = r[
w ξη∂η(r2ρ) + ξx∂x(r2p)]
(16.108)
and simplifying using the energy conservation equation for ∂ηρ and the momentum conservation equation
for ∂xp.
16.4.14 Integration variable
It is desirable to choose an integration variable that varies monotonically. A natural choice is the proper
time τ in the tetrad frame, since this is guaranteed to increase monotonically. Since the 4-velocity at rest
in the tetrad frame is by definition um = 1, 0, 0, 0, the proper time derivative is related to the directed
conformal time derivative in the tetrad frame by d/dτ = um∂m = ∂η.
However, there is another choice of integration variable, the ray-tracing variable X defined by equa-
tion (16.79), that is not specifically tied to any tetrad frame, and that has a desirable (tetrad and coordinate)
gauge-invariant meaning. The proper time derivative of any dimensionless function F (x) in the tetrad frame
is related to its derivative dF/dX with respect to the ray-tracing variable X by
∂ηF = um∂mF = uX∂XF = − ξx
r
dF
dX. (16.109)
In the third expression, uX∂XF is um∂mF expressed in the similarity frame of §16.4.4, the time contribu-
tion uη×∂η×F vanishing in the similarity frame because it is proportional to the conformal time derivative
∂F/∂η× = 0. In the last expression of (16.109), uX has been replaced by −ux = −ξx/∆1/2 in view of equa-
tion (16.77), the minus sign coming from the fact that uX is the radial component of the tetrad 4-velocity of
the tetrad frame relative to the similarity frame, while ux in equation (16.77) is the radial component of the
tetrad 4-velocity of the similarity frame relative to the tetrad frame. Also in the last expression of (16.109),
the directed derivative ∂X with respect to the ray-tracing variable X has been translated into its coordinate
partial derivative, ∂X = (∆1/2/r) ∂/∂X , which follows from the metric (16.80).
16.4 Self-similar spherically symmetric spacetime 287
In summary, the chosen integration variable is the dimensionless ray-tracing variable −X (with a minus
because −X is monotonically increasing), the derivative with respect to which, acting on any dimensionless
function, is related to the proper time derivative in any tetrad frame (not just the baryonic frame) by
− d
dX=
r
ξx∂η . (16.110)
Equation (16.110) involves ξx, which is proportional to the proper velocity of the tetrad frame through the
similarity frame, equation (16.78), and which therefore, being initially positive, must always remain positive
in any tetrad frame attached to a fluid, as long as the fluid does not turn back on itself, as must be true for
the self-similar solution to be consistent.
16.4.15 Summary of equations for a single charged fluid
For reference, it is helpful to collect here the full set of equations governing self-similar spherically symmetric
evolution in the case of a single charged fluid with isotropic equation of state
p = p⊥ = w ρ , (16.111)
and conductivity
σ = κ ρ1/2 . (16.112)
In accordance with the arguments in §16.4.9, equations (16.95) and (16.97), self-similarity requires that the
pressure-to-density ratio w and the conductivity coefficient κ both be (dimensionless) constants.
It is natural to work in the center-of-mass frame of the fluid, which also coincides with the center-of-mass
frame of the fluid plus electric field (the electric field, being invariant under Lorentz boosts, does not pick
out any particular radial frame).
The proper time τ in the fluid frame evolves as
− dτ
dX=
r
ξx, (16.113)
which follows from equation (16.110) and the fact that ∂ητ = 1. The circumferential radius r evolves along
the path of the fluid as
− d ln r
dX=βη
ξx. (16.114)
Although it is straightforward to write down the equations governing how the tetrad frame moves through the
conformal coordinates η and x, there is not much to be gained from this because the conformal coordinates
have no fundamental physical significance.
Next, the defining equations (16.100) for the proper acceleration hη and Hubble parameter hx yield
equations for the evolution of the time and radial components of the conformal Killing vector ξm
− dξη
dX= βx − rhη , (16.115a)
288 General spherically symmetric spacetime
− dξx
dX= − βη + rhx , (16.115b)
in which, in the formula for hη, equation (16.103) has been used to convert the conformal radial derivative
∂x to the conformal time derivative ∂η, and thence to −d/dX by equation (16.110).
Next, the Einstein equations (16.26c) and (16.26c) (with coordinates relabeled per (16.98) in the center-of-
mass frame (16.33) yield evolution equations for the time and radial components of the vierbein coefficients
βm
− dβη
dX= − βx
ξηrhx , (16.116a)
− dβx
dX=βη
ξxrhη , (16.116b)
where again, in the formula for βη, equation (16.103) has been used to convert the conformal radial derivative
∂x to the conformal time derivative ∂η. The 4 evolution equations (16.115) and (16.116) for ξm and βm are
not independent: they are related by ξmβm = v, a constant, equation (16.70). To maintain numerical
precision, it is important to avoid expressing small quantities as differences of large quantities. In practice,
a suitable choice of variables to integrate proves to be ξη + ξx, βη − βx, and βx, each of which can be
tiny in some circumstances. Starting from these variables, the following equations yield ξη − ξx, along with
the interior mass M and the horizon function ∆, equations (16.69) and (16.71), in a fashion that ensures
numerical stability:
ξη − ξx =2v− (ξη + ξx)(βη + βx)
βη − βx, (16.117a)
2M
r= 1 + (βη + βx)(βη − βx) , (16.117b)
∆ = (ξη + ξx)(ξη − ξx) . (16.117c)
The evolution equations (16.115) and (16.116) involve hη and hx. The integrals of motion considered in
§16.4.13 yield explicit expressions for hη and hx not involving any derivatives. For the Hubble parameter
hx, equation (16.105a) gives
rhx = − ξη
ξxrhη +
ξη
v4πε , (16.118)
where ε is the dimensionless enthalpy
ε ≡ r2(1 + w)ρ . (16.119)
For the proper acceleration hη, a somewhat lengthy calculation starting from the integral of motion (16.108),
and simplifying using the integral of motion (16.107) for Q, the expression (16.118) for hx, Maxwell’s equa-
tion (16.54b) [with the relabelling (16.98)], and the conductivity (16.112) in Ohm’s law (16.96), gives
rhη =ξx
8πw⊥(βxξx − wβηξ
η)r2ρ+ [v + (1 + w)4πrσξη ]Q2/r2 − w(4πξηε)2/v
4πε [(ξx)2 − w(ξη)2]. (16.120)
16.4 Self-similar spherically symmetric spacetime 289
Two more equations complete the suite. The first, which represents energy conservation for the fluid, can
be written as an equation governing the entropy S of the fluid
− d lnS
dX=
σQ2
r(1 + w)ρξx, (16.121)
in which the S is (up to an arbitrary constant) the entropy of a comoving volume element V ∝ r3ξx of the
fluid
S ≡ r3ξxρ1/(1+w) . (16.122)
That equation (16.121) really is an entropy equation can be confirmed by rewriting it as
1
V
(
dρV
dτ+ p
dV
dτ
)
= jE =σQ2
r4, (16.123)
in which jE is recognized as the Ohmic dissipation, the rate per unit volume at which the volume element
V is being heated.
The final equation represents electromagnetic energy conservation, equation (16.57a), which can be written
− d lnQ
dX= − 4πrσ
ξx. (16.124)
The (heat) energy going into the fluid is balanced by the (free) energy coming out of the electromagnetic
field.
17
The interiors of spherical black holes
As discussed in Chapter 8, the Reissner-Nordstrom geometry for an ideal charged spherical black hole
contains mathematical wormhole and white hole extensions to other universes. In reality, these extensions
are not expected to occur, thanks to the mass inflation instability discovered by Poisson & Israel (1990).
17.1 The mechanism of mass inflation
17.1.1 Reissner-Nordstrom phase
Figure 17.1 illustrates how the two Einstein equations (16.52) produce the three phases of mass inflation
inside a charged spherical black hole.
During the initial phase, illustrated in the top panel of Figure 17.1, the spacetime geometry is well-
approximated by the vacuum, Reissner-Nordstrom geometry. During this phase the radial energy flux f is
effectively zero, so βr remains constant, according to equation (16.52b). The change in the radial velocity βt,
equation (16.52a), depends on the competition between the Newtonian gravitational force −M/r2, which is
always attractive (tending to make the radial velocity βt more negative), and the gravitational force −4πrp
sourced by the radial pressure p. In the Reissner-Nordstrom geometry, the static electric field produces a
negative radial pressure, or tension, p = −Q2/(8πr4), which produces a gravitational repulsion −4πrp =
Q2/(2r3). At some point (depending on the charge-to-mass ratio) inside the outer horizon, the gravitational
repulsion produced by the tension of the electric field exceeds the attraction produced by the interior mass
M , so that the radial velocity βt slows down. This regime, where the (negative) radial velocity βt is slowing
down (becoming less negative), while βr remains constant, is illustrated in the top panel of Figure 17.1.
If the initial Reissner-Nordstrom phase were to continue, then the radial 4-gradient βm would become
lightlike. In the Reissner-Nordstrom geometry this does in fact happen, and where it happens defines the
inner horizon. The problem with this is that the lightlike 4-vector βm points in one direction for ingoing
frames, and in the opposite direction for outgoing frames. If βm becomes lightlike, then ingoing and outgoing
frames are streaming through each other at the speed of light. This is the infinite blueshift at the inner
horizon first pointed out by Penrose (1968).
17.1 The mechanism of mass inflation 291
βt
βr
outg
oing
ingoing
1. RN
βr
outg
oing
ingoing
2. Inflation
βr
outg
oing
ingoing
3. Collapse
Figure 17.1 Spacetime diagrams of the tetrad-frame 4-vector βm, equation (16.8), illustrating qualitativelythe three successive phases of mass inflation: 1. (top) the Reissner-Nordstrom phase, where inflation ignites;2. (middle) the inflationary phase itself; and 3. (bottom) the collapse phase, where inflation comes to an end.In each diagram, the arrowed lines labeled ingoing and outgoing illustrate two representative examples ofthe 4-vector βt, βr, while the double-arrowed lines illustrate the rate of change of these 4-vectors impliedby Einstein’s equations (16.52). Inside the horizon of a black hole, all locally inertial frames necessarily fallinward, so the radial velocity βt ≡ ∂tr is always negative. A locally inertial frame is ingoing or outgoingdepending on whether the proper radial gradient βr ≡ ∂rr measured in that frame is positive or negative.
292 The interiors of spherical black holes
If there were no matter present, or if there were only one stream of matter, either ingoing or outgoing but
not both, then βm could indeed become lightlike. But if both ingoing and outgoing matter are present, even
in the tiniest amount, then it is physically impossible for the ingoing and outgoing frames to stream through
each other at the speed of light.
If both ingoing and outgoing streams are present, then as they race through each other ever faster, they
generate a radial pressure p, and an energy flux f , which begin to take over as the main source on the right
hand side of the Einstein equations (16.52). This is how mass inflation is ignited.
17.1.2 Inflationary phase
The infalling matter now enters the second, mass inflationary phase, illustrated in the middle panel of
Figure 17.1.
During this phase, the gravitational force on the right hand side of the Einstein equation (16.52a) is
dominated by the pressure p produced by the counter-streaming ingoing and outgoing matter. The mass M
is completely sub-dominant during this phase (in this respect, the designation “mass inflation” is misleading,
since although the mass inflates, it does not drive inflation). The counter-streaming pressure p is positive,
and so accelerates the radial velocity βt (makes it more negative). At the same time, the radial gradient
βr is being driven by the energy flux f , equation (16.52b). For typically low accretion rates, the streams
are cold, in the sense that the streaming energy density greatly exceeds the thermal energy density, even if
the accreted material is at relativistic temperatures. This follows from the fact that for mass inflation to
begin, the gravitational force produced by the counter-streaming pressure p must become comparable to that
produced by the mass M , which for streams of low proper density requires a hyper-relativistic streaming
velocity. For a cold stream of proper density ρ moving at 4-velocity um ≡ ut, ur, 0, 0, the streaming energy
flux would be f ∼ ρutur, while the streaming pressure would be p ∼ ρ(ur)2. Thus their ratio f/p ∼ ut/ur
is slightly greater than one. It follows that, as illustrated in the middle panel of Figure 17.1, the change in
βr slightly exceeds the change in βt, which drives the 4-vector βm, already nearly lightlike, to be even more
nearly lightlike. This is mass inflation.
Inflation feeds on itself. The radial pressure p and energy flux f generated by the counter-streaming ingoing
and outgoing streams increase the gravitational force. But, as illustrated in the middle panel of Figure 17.1,
the gravitational force acts in opposite directions for ingoing and outgoing streams, tending to accelerate
the streams faster through each other. An intuitive way to understand this is that the gravitational force
is always inwards, meaning in the direction of smaller radius, but the inward direction is towards the black
hole for ingoing streams, and away from the black hole for outgoing streams.
The feedback loop in which the streaming pressure and flux increase the gravitational force, which ac-
celerates the streams faster through each other, which increases the streaming pressure and flux, is what
drives mass inflation. Inflation produces an exponential growth in the streaming energy, and along with it
the interior mass, and the Weyl curvature.
17.2 The far future? 293
17.1.3 Collapse phase
It might seem that inflation is locked into an exponential growth from which there is no exit. But the
Einstein equations (16.52) have one more trick up their sleave.
For the counter-streaming velocity to continue to increase requires that the change in βr from equa-
tion (16.52b) continues to exceed the change in βt from equation (16.52a). This remains true as long as the
counter-streaming pressure p and energy flux f continue to dominate the source on the right hand side of
the equations. But the mass term −M/r2 also makes a contribution to the change in βt, equation (16.52a).
As will be seen in the examples of the next two sections, §§?? and ??, at least in the case of pressureless
streams the mass term exponentiates slightly faster than the pressure term. At a certain point, the additional
acceleration produced by the mass means that the combined gravitational force M/r2 + 4πrp exceeds 4πrf .
Once this happens, the 4-vector βm, instead of being driven to becoming more lightlike, starts to become
less lightlike. That is, the counter-streaming velocity starts to slow. At that point inflation ceases, and the
streams quickly collapse to zero radius.
It is ironic that it is the increase of mass that brings mass inflation to an end. Not only does mass not
drive mass inflation, but as soon as mass begins to contribute significantly to the gravitational force, it brings
mass inflation to an end.
17.2 The far future?
The Penrose diagram of a Reissner-Nordstrom or Kerr-Newman black hole indicates that an observer who
passes through the outgoing inner horizon sees the entire future of the outside universe go by. In a sense,
this is “why” the outside universe appears infinitely blueshifted.
This raises the question of whether what happens at the outgoing inner horizon of a real black hole indeed
depends on what happens in the far future. If it did, then the conclusions of previous sections, which are
based in part on the proposition that the accretion rate is approximately constant, would be suspect. A
lot can happen in the far future, such as black hole mergers, the Universe ending in a big crunch, Hawking
evaporation, or something else beyond our current ken.
Ingoing and outgoing observers both see each other highly blueshifted near the inner horizon. An outgoing
observer sees ingoing observers from the future, while an ingoing observer sees outgoing observers from the
past. If the streaming 4-velocity between ingoing and outgoing streams is umba, equation (??), then the proper
time dτb that elapses on stream b observed by stream a equals the blueshift factor utba times the proper time
dτa experienced by stream a,
dτb = utba dτa . (17.1)
17.2.1 Inflationary phase
A physically relevant timescale for the observing stream a is how long it takes for the blueshift to increase
by one e-fold, which is dτa/d lnutba. During the inflationary phase, stream a sees the amount of time elapsed
294 The interiors of spherical black holes
on stream b through one e-fold of blueshift to be
utba
dτad lnut
ba
=2Cbr−λ
. (17.2)
The right hand side of equation (17.2) is derived from utba ≈ 2|uaub|, |dra/dτa| = |βa,t| ≈ βua, ra ≈ r−,
and the approximations (??) and (??) valid during the inflationary phase. The constants Cb and λ on the
right hand side of equation (17.2) are typically of order unity, while r− is the radius of the inner horizon
where mass inflation takes place. Thus the right hand side of equation (17.2) is of the order of one black hole
crossing time. In other words, stream a sees approximately one black hole crossing time elapse on stream b
for each e-fold of blueshift.
For astronomically realistic black holes, exponentiating the Weyl curvature up to the Planck scale will
take typically a few hundred e-folds of blueshift, as illustrated for example in Figure 17.10. Thus what
happens at the inner horizon of a realistic black hole before quantum gravity intervenes depends only on the
immediate past and future of the black hole – a few hundred black hole crossing times – not on the distant
future or past. This conclusion holds even if the accretion rate of one of the ingoing or outgoing streams is
tiny compared to the other, as considered in §??.
From a stream’s own point of view, the entire inflationary episode goes by in a flash. At the onset of
inflation, where β goes through its minimum, at µau2a = µbu
2b = λ according to equation (??), the blueshift
utba is already large
utba =
2λ√µaµb
, (17.3)
thanks to the small accretion rates µa and µb. During the first e-fold of blueshift, each stream experiences a
proper time of order√µaµb times the black hole crossing time, which is tiny. Subsequent e-folds of blueshift
race by in proportionately shorter proper times.
17.2.2 Collapse phase
The time to reach the collapse phase is another matter. According to the estimates in §§?? and ??, reaching
the collapse phase takes of order ∼ 1/µa e-folds of blueshift, where µa is the larger of the accretion rates
of the ingoing and outgoing streams, equation (??). Thus in reaching the collapse phase, each stream has
seen approximately 1/µa black hole crossing times elapse on the other stream. But 1/µa black hole crossing
times is just the accretion time – essentially, how long the black hole has existed. This timescale, the age of
the black hole, is not infinite, but it can hardly be expected that the accretion rate of a black hole would be
constant over its lifetime.
If the accretion rate were in fact constant, and if quantum gravity did not intervene and the streams
remained non-interacting, then indeed streams inside the black hole would reach the collapse phase, where-
upon they would plunge to a spacelike singularity at zero radius. For example, in the self-similar models
illustrated in Figures ?? and 17.10, outgoing baryons hitting the central singularity see ingoing dark matter
accreted a factor of two into the future (specifically, for ingoing baryons and outgoing dark matter hitting the
singularity, the radius of the outer horizon when the dark matter is accreted is twice that when the baryons
17.3 Self-similar models of the interior structure of black holes 295
were accreted; the numbers are 2.11 for M• = 0.03, 1.97 for M• = 0.01, and unknown for M• = 0.003 because
in that case the numerics overflow before the central singularity is reached). The same conclusion applies to
the ultra-hard fluid models illustrated in Figure 17.8: if the accretion rate is constant, then outgoing streams
see only a factor of order unity or a few into the future before hitting the central singularity.
If on the other hard the accretion rate decreases sufficiently rapidly with time, then it is possible that
an outgoing stream never reaches the collapse phase, because the number of e-folds to reach the collapse
phase just keeps increasing as the accretion rate decreases. By contrast, an ingoing stream is always liable to
reach the collapse phase, if quantum gravity does not intervene, because an ingoing stream sees the outgoing
stream from the past, when the accretion rate was liable to have been larger.
However, as already remarked in §??, in astronomically realistic black holes, it is only for large accretion
rates, such as may occur when the black hole first collapses (M• & 0.01 for the models illustrated in
Fig. 17.10), that the collapse phase has a chance of being reached before the Weyl curvature exceeds the
Planck scale.
To summarize, it is only streams accreted during the first few hundred or so black hole crossing times
since a black hole’s formation that have a chance of hitting a central spacelike singularity. Streams accreted
at later times, whether ingoing or outgoing, are likely to meet their fate in the inflationary zone at the inner
horizon, where the Weyl curvature exponentiates to the Planck scale and beyond.
17.3 Self-similar models of the interior structure of black holes
The apparatus is now in hand actually to do some real calculations of the interior structure of black holes.
All the models presented in this section are spherical and self-similar. See Hamilton & Pollack (2005, PRD
71, 084031 & 084031) and Wallace, Hamilton & Polhemus (2008, arXiv:0801.4415) for more.
17.3.1 Boundary conditions at an outer sonic point
Because information can propagate only inward inside the horizon of a black hole, it is natural to set the
boundary conditions outside the horizon. The policy adopted here is to set them at a sonic point, where
the infalling fluid accelerates from subsonic to supersonic. The proper 3-velocity of the fluid through the
self-similar frame is ξx/ξη, equation (16.78) (the velocity ξx/ξη is positive falling inward), and the sound
speed is
sound speed =
√
pb
ρb=√wb , (17.4)
and sonic points occur where the velocity equals the sound speed
ξx
ξη= ±√wb at sonic points . (17.5)
The denominator of the expression (16.120) for the proper acceleration g is zero at sonic points, indicating
that the acceleration will diverge unless the numerator is also zero. What happens at a sonic point depends
296 The interiors of spherical black holes
on whether the fluid transitions from subsonic upstream to supersonic downstream (as here) or vice versa.
If (as here) the fluid transitions from subsonic to supersonic, then sound waves generated by discontinuities
near the sonic point can propagate upstream, plausibly modifying the flow so as to ensure a smooth transition
through the sonic point, effectively forcing the numerator, like the denominator, of the expression (16.120)
to pass through zero at the sonic point. Conversely, if the fluid transitions from supersonic to subsonic, then
sound waves cannot propagate upstream to warn the incoming fluid that a divergent acceleration is coming,
and the result is a shock wave, where the fluid accelerates discontinuously, is heated, and thereby passes
from supersonic to subsonic.
The solutions considered here assume that the acceleration g at the sonic point is not only continuous [so
the numerator of (16.120) is zero] but also differentiable. Such a sonic point is said to be regular, and the
assumption imposes two boundary conditions at the sonic point.
The accretion in real black holes is likely to be much more complicated, but the assumption of a regular
sonic point is the simplest physically reasonable one.
17.3.2 Mass and charge of the black hole
The mass M• and charge Q• of the black hole at any instant can be defined to be those that would be
measured by a distant observer if there were no mass or charge outside the sonic point
M• = M +Q2
2r, Q• = Q at the sonic point . (17.6)
The mass M• in equation (17.6) includes the mass-energy Q2/2r that would be in the electric field outside
the sonic point if there were no charge outside the sonic point, but it does not include mass-energy from any
additional mass or charge that might be outside the sonic point.
In self-similar evolution, the black hole mass increases linearly with time, M• ∝ t, where t is the proper
time at rest far from the black hole. As discussed in §??, this time t equals the proper time τd = rξηd/v
recorded by dark matter clocks that free-fall radially from zero velocity at infinity. Thus the mass accretion
rate M• is
M• ≡dM•dt
=M•τd
=vM•rξη
d
at the sonic point . (17.7)
If there is no mass outside the sonic point (apart from the mass-energy in the electric field), then a freely-
falling dark matter particle will have
βd,x = 1 at the sonic point , (17.8)
which can be taken as the boundary condition on the dark matter at the sonic point, for either massive or
massless dark matter. Equation (17.8) follows from the facts that the geodesic equations in empty space
around a charged black hole (Reissner-Nordstrom metric) imply that βd,x = constant for a radially free-
falling particle (the same conclusion can drawn from the Einstein equation (16.26c)), and that a particle at
rest at infinity satisfies βd,η = ∂d,ηr = 0, and consequently βd,x = 1 from equation (16.69) with r →∞.
As remarked following equation (16.72), the residual gauge freedom in the global rescaling of conformal
17.3 Self-similar models of the interior structure of black holes 297
time η allows the expansion velocity v to be adjusted at will. One choice suggested by equation (17.7) is to
set (but one could equally well set v to the expansion velocity of the horizon, v = r+, for example)
v = M• , (17.9)
which is equivalent to setting
ξηd =
M•r
at the sonic point . (17.10)
Equation (17.10) is not a boundary condition: it is just a choice of units of conformal time η. Equation (17.10)
and the boundary condition (17.8) coupled with the scalar relations (16.69) and (16.70) fully determine the
dark matter 4-vectors βd,m and ξmd at the sonic point.
17.3.3 Equation of state
The density ρb and temperature Tb of an ideal relativistic baryonic fluid in thermodynamic equilibrium are
related by
ρb =π2g
30T 4
b , (17.11)
where
g = gB +7
8gF (17.12)
is the effective number of relativistic particle species, with gB and gF being the number of bosonic and
fermionic species. If the expected increase in g with temperature T is modeled (so as not to spoil self-
similarity) as a weak power law g/gp = T ǫ, with gp the effective number of relativistic species at the Planck
temperature, then the relation between density ρb and temperature Tb is
ρb =π2gp30
T(1+w)/wb , (17.13)
with equation of state parameter wb = 1/(3 + ǫ) slightly less than the standard relativistic value w = 1/3.
In the models considered here, the baryonic equation of state is taken to be
wb = 0.32 . (17.14)
The effective number gp is fixed by setting the number of relativistic particles species to g = 5.5 at T =
10 MeV, corresponding to a plasma of relativistic photons, electrons, and positrons. This corresponds to
choosing the effective number of relativistic species at the Planck temperature to be gp ≈ 2,400, which is
not unreasonable.
The chemical potential of the relativistic baryonic fluid is likely to be close to zero, corresponding to equal
numbers of particles and anti-particles. The entropy Sb of a proper Lagrangian volume element V of the
fluid is then
Sb =(ρb + pb)V
Tb, (17.15)
298 The interiors of spherical black holes
which agrees with the earlier expression (16.122), but now has the correct normalization.
17.3.4 Entropy creation
One fundamentally interesting question about black hole interiors is how much entropy might be created
inside the horizon. Bekenstein first argued that a black hole should have a quantum entropy proportional to
its horizon area A, and Hawking (1974) supplied the constant of proportionality 1/4 in Planck units. The
Bekenstein-Hawking entropy SBH is, in Planck units c = G = ~ = 1,
SBH =A
4. (17.16)
For a spherical black hole of horizon radius r+, the area is A = 4πr2+. Hawking showed that a black hole
has a temperature TH equal to 1/(2π) times the surface gravity g+ at its horizon, again in Planck units,
TH =g+2π
. (17.17)
The surface gravity is defined to be the proper radial acceleration measured by a person in free-fall at the
horizon. For a spherical black hole, the surface gravity is g+ = −Dtβt = M/r2 + 4πrp evaluated at the
horizon, equation (16.52a).
The proper velocity of the baryonic fluid through the sonic point equals ξx/ξη, equation (16.78). Thus
the entropy Sb accreted through the sonic point per unit proper time of the fluid is
dSb
dτ=
(1 + wb)ρb
Tb
4πr2ξx
ξη. (17.18)
The horizon radius r+, which is at fixed conformal radius x, expands in proportion to the conformal factor,
r+ ∝ a, and the conformal factor a increases as d ln a/dτ = ∂η ln a = v/(rξη), so the Bekenstein-Hawking
entropy SBH = πr2+ increases as
dSBH
dτ=
2πr2+v
rξη. (17.19)
Thus the entropy Sb accreted through the sonic point per unit increase of the Bekenstein-Hawking entropy
SBH is
dSb
dSBH=
(1 + wb)ρb4πr3ξx
2πr2+vTb
∣
∣
∣
∣
r=rs
. (17.20)
Inside the sonic point, dissipation increases the entropy according to equation (16.121). Since the entropy
can diverge at a central singularity where the density diverges, and quantum gravity presumably intervenes
at some point, it makes sense to truncate the production of entropy at a “splat” point where the density ρb
hits a prescribed splat density ρ#
ρb = ρ# . (17.21)
Integrating equation (16.121) from the sonic point to the splat point yields the ratio of the entropies at the
17.3 Self-similar models of the interior structure of black holes 299
sonic and splat points. Multiplying the accreted entropy, equation (17.20), by this ratio yields the rate of
increase of the entropy of the black hole, truncated at the splat point, per unit increase of its Bekenstein-
Hawking entropy
dSb
dSBH=
(1 + wb)ρb4πr3ξx
2πr2+vTb
∣
∣
∣
∣
ρb=ρ#
. (17.22)
17.3.5 Holography
The idea that the entropy of a black hole cannot exceed its Bekenstein-Hawking entropy has motivated
holographic conjectures that the degrees of freedom of a volume are somehow encoded on its boundary,
and consequently that the entropy of a volume is bounded by those degrees of freedom. Various counter-
examples dispose of most simple-minded versions of holographic entropy bounds. The most successful entropy
bound, with no known counter-examples, is Bousso’s covariant entropy bound (Bousso 2002, Rev. Mod.
Phys. 74, 825). The covariant entropy bound concerns not just any old 3-dimensional volume, but rather the
3-dimensional volume formed by a null hypersurface, a lightsheet. For example, the horizon of a black hole is
a null hypersurface, a lightsheet. The covariant entropy bound asserts that the entropy that passes (inward
or outward) through a lightsheet that is everywhere converging cannot exceed 1/4 of the 2-dimensional area
of the boundary of the lightsheet.
In the self-similar black holes under consideration, the horizon is expanding, and outgoing lightrays that
sit on the horizon do not constitute a converging lightsheet. However, a spherical shell of ingoing lightrays
that starts on the horizon falls inwards and therefore does form a converging lightsheet, and a spherical shell
of outgoing lightrays that starts just slightly inside the horizon also falls inward and forms a converging
lightsheet. The rate at which entropy Sb passes through such ingoing or outgoing spherical lightsheets per
unit decrease in the area Scov ≡ πr2 of the lightsheet is∣
∣
∣
∣
dSb
dScov
∣
∣
∣
∣
=dSb
dSBH
r2+r2
v
ξx|βη ∓ βx|, (17.23)
in which the ∓ sign is − for ingoing, + for outgoing lightsheets. A sufficient condition for Bousso’s covariant
entropy bound to be satisfied is
|dSb/dScov| ≤ 1 . (17.24)
17.3.6 Black hole accreting a neutral relativistic plasma
The simplest case to consider is that of a black hole accreting a neutral relativistic plasma. In the self-similar
solutions, the charge of the black hole is produced self-consistently by the accreted charge of the baryonic
fluid, so a neutral fluid produces an uncharged black hole.
Figure 17.2 shows the baryonic density ρb and Weyl curvature C inside the uncharged black hole. The
mass and accretion rate have been taken to be
M• = 4× 106 M⊙ , M• = 10−16 , (17.25)
300 The interiors of spherical black holes
hori
zon
Plan
cksc
ale
ρb
−C
dSb/dSBH
1010 1020 1030 1040 105010−12010−11010−10010−9010−8010−7010−6010−5010−4010−3010−2010−10
10010101020
Radius r (Planck units)
(Pla
nck
unit
s)
Figure 17.2 An uncharged baryonic plasma falls into an uncharged spherical black hole. The plot shows inPlanck units, as a function of radius, the plasma density ρb, the Weyl curvature scalar C (which is negative),and the rate dSb/dSBH of increase of the plasma entropy per unit increase in the Bekenstein-Hawking entropy
of the black hole. The mass is M• = 4×106 M⊙, the accretion rate is M• = 10−16, and the equation of stateis wb = 0.32.
which are motivated by the fact that the mass of the supermassive black hole at the center of the Milky Way
is 4× 106 M⊙, and its accretion rate is
Mass of MW black hole
age of Universe≈ 4× 106 M⊙
1010 yr≈ 6× 1060 Planck units
4× 1044 Planck units≈ 10−16 . (17.26)
Figure 17.2 shows that the baryonic plasma plunges uneventfully to a central singularity, just as in the
Schwarzschild solution. The Weyl curvature scalar hits the Planck scale, |C| = 1, while the baryonic proper
density ρb is still well below the Planck density, so this singularity is curvature-dominated.
17.3.7 Black hole accreting a non-conducting charged relativistic plasma
The next simplest case is that of a black hole accreting a charged but non-conducting relativistic plasma.
Figure 17.3 shows a black hole with charge-to-mass Q•/M• = 10−5, but otherwise the same parameters
as in the uncharged black hole of §17.3.6: M• = 4 × 106 M⊙, M• = 10−16, and wb = 0.32. Inside the outer
horizon, the baryonic plasma, repelled by the electric charge of the black hole self-consistently generated by
the accretion of the charged baryons, becomes outgoing. Like the Reissner-Nordstrom geometry, the black
hole has an (outgoing) inner horizon. The baryons drop through the inner horizon, shortly after which the
self-similar solution terminates at an irregular sonic point, where the proper acceleration diverges. Normally
17.3 Self-similar models of the interior structure of black holes 301
hori
zon
inne
rho
rizo
n
ρb
ρe
|C|
dSb/dSBH
1010 1020 1030 1040 105010−12010−11010−10010−9010−8010−7010−6010−5010−4010−3010−2010−10
10010101020
Radius r (Planck units)
(Pla
nck
unit
s)
Figure 17.3 A plasma that is charged but non-conducting. The black hole has an inner horizon like theReissner-Nordstrom geometry. The self-similar solution terminates at an irregular sonic point just beneaththe inner horizon. The mass is M• = 4 × 106 M⊙, accretion rate M• = 10−16, equation of state wb = 0.32,and black hole charge-to-mass Q•/M• = 10−5.
this is a signal that a shock must form, but even if a shock is introduced, the plasma still terminates at an
irregular sonic point shortly downstream of the shock. The failure of the self-similar to continue does not
invalidate the solution, because the failure is hidden beneath the inner horizon, and cannot be communicated
to infalling matter above it.
The solution is nevertheless not realistic, because it assumes that there is no ingoing matter, such as
would inevitably be produced for example by infalling neutral dark matter. Such ingoing matter would
appear infinitely blueshifted to the outgoing baryons falling through the inner horizon, which would produce
mass inflation, as in §17.3.10.
17.3.8 Black hole accreting a conducting relativistic plasma
What happens if the baryonic plasma is not only electrically charged but also electrically conducting? If the
conductivity is small, then the solutions resemble the non-conducting solutions of the previous subsection,
§17.3.7. But if the conductivity is large enough effectively to neutralize the plasma as it approaches the
center, then the plasma can plunge all the way to the central singularity, as in the uncharged case in §17.3.6.
Figure 17.4 shows a case in which the conductivity has been tuned to equal, within numerical accuracy, the
critical conductivity κb = 0.35 above which the plasma collapses to a central singularity. The parameters are
302 The interiors of spherical black holes
hori
zon
Plan
cksc
ale
ρb
ρe
|C|
|dSb/dScov|
dSb/dSBH
1010 1020 1030 1040 105010−12010−11010−10010−9010−8010−7010−6010−5010−4010−3010−2010−10
10010101020
Radius r (Planck units)
(Pla
nck
unit
s)
Figure 17.4 Here the baryonic plasma is charged, and electrically conducting. The conductivity is at (withinnumerical accuracy) the threshold above which the plasma plunges to a central singularity. The mass is
M• = 4 × 106 M⊙, the accretion rate M• = 10−16, the equation of state wb = 0.32, the charge-to-massQ•/M• = 10−5, and the conductivity parameter κb = 0.35. Arrows show how quantities vary a factor of 10into the past and future.
otherwise the same as in previous subsections: a mass of M• = 4× 106 M⊙, an accretion rate M• = 10−16,
an equation of state wb = 0.32, and a black hole charge-to-mass of Q•/M• = 10−5.
The solution at the critical conductivity exhibits the periodic self-similar behavior first discovered in
numerical simulations by Choptuik (1993, PRL 70, 9), and known as “critical collapse” because it happens
at the borderline between solutions that do and do not collapse to a black hole. The ringing of curves in
Figure 17.4 is a manifestation of the self-similar periodicity, not a numerical error.
These solutions are not subject to the mass inflation instability, and they could therefore be prototypical
of the behavior inside realistic rotating black holes. For this to work, the outward transport of angular
momentum inside a rotating black hole must be large enough effectively to produce zero angular momentum
at the center. My instinct is that angular momentum transport is probably not strong enough, but I do not
know this for sure. If angular momentum transport is not strong enough, then mass inflation will take place.
Figure 17.4 shows that the entropy produced by Ohmic dissipation inside the black hole can potentially
exceed the Bekenstein-Hawking entropy of the black hole by a large factor. The Figure shows the rate
dSb/dSBH of increase of entropy per unit increase in its Bekenstein-Hawking entropy, as a function of the
hypothetical splat point above which entropy production is truncated. The rate is almost independent of
the black hole mass M• at fixed splat density ρ#, so it is legitimate to interpret dSb/dSBH as the cumulative
17.3 Self-similar models of the interior structure of black holes 303
hori
zon
Plan
cksc
ale
ρb
ρe|C|
|dSb/dScov|
dSb/dSBH
1010 1020 1030 1040 105010−12010−11010−10010−9010−8010−7010−6010−5010−4010−3010−2010−10
10010101020
Radius r (Planck units)
(Pla
nck
unit
s)
Figure 17.5 This black hole creates a lot of entropy by having a large charge-to-mass Q•/M• = 0.8 and a
low accretion rate M• = 10−28. The conductivity parameter κb = 0.35 is again at the threshold above whichthe plasma plunges to a central singularity. The equation of state is wb = 0.32.
entropy created inside the black hole relative to the Bekenstein-Hawking entropy. Truncated at the Planck
scale, |C| = 1, the entropy relative to Bekenstein-Hawking is dSb/dSBH ≈ 1010.
Generally, the smaller the accretion rate M•, the more entropy is produced. If moreover the charge-to-
mass Q•/M• is large, then the entropy can be produced closer to the outer horizon. Figure 17.5 shows a
model with a relatively large charge-to-mass Q•/M• = 0.8, and a low accretion rate M• = 10−28. The large
charge-to-mass ratio in spite of the relatively high conductivity requires force-feeding the black hole: the
sonic point must be pushed to just above the horizon. The large charge and high conductivity leads to a
burst of entropy production just beneath the horizon.
If the entropy created inside a black hole exceeds the Bekenstein-Hawking entropy, and the black hole later
evaporates radiating only the Bekenstein-Hawking entropy, then entropy is destroyed, violating the second
law of thermodynamics.
This startling conclusion is premised on the assumption that entropy created inside a black hole accumu-
lates additively, which in turn derives from the assumption that the Hilbert space of states is multiplicative
over spacelike-separated regions. This assumption, called locality, derives from the fundamental proposi-
tion of quantum field theory in flat space that field operators at spacelike-separated points commute. This
reasoning is essentially the same as originally led Hawking (1976) to conclude that black holes must destroy
information.
The same ideas that motivate holography also rescue the second law. If the future lightcones of spacelike-
separated points do not intersect, then the points are permanently out of communication, and can behave
304 The interiors of spherical black holes
matterInfalling
Sb >> SBH
Sb = SBH
Sb <
SBH
Figure 17.6 Partial Penrose diagram of the black hole. The entropy passing through the spacelike slice beforethe black hole evaporates exceeds that passing through the spacelike slice after the black hole evaporates,apparently violating the second law of thermodynamics. However, the entropy passing through any null slicerespects the second law.
like alternate quantum realities, like Schrodinger’s dead-and-alive quantum cat. Just as it is not legitimate
to the add the entropies of the dead cat and the live cat, so also it is not legitimate to add the entropies of
regions inside a black hole whose future lightcones do not intersect. The states of such separated regions,
instead of being distinct, are quantum entangled with each other.
Figures 17.4 and 17.5 show that the rate |dSb/dScov| at which entropy passes through ingoing or outgoing
spherical lightsheets is less than one at all scales below the Planck scale. This shows not only that the black
holes obey Bousso’s covariant entropy bound, but also that no individual observer inside the black hole sees
more than the Bekenstein-Hawking entropy on their lightcone. No observer actually witnesses a violation of
the second law.
17.3.9 Black hole accreting a charged massless scalar field
The charged, non-conducting plasma considered in §17.3.7 fell through an (outgoing) inner horizon without
undergoing mass inflation. This can be attributed to the fact that relativistic counter-streaming could not
occur: there was only a single (outgoing) fluid, and the speed of sound in the fluid was less than the speed
of light.
In reality, unless dissipation destroys the inner horizon as in §17.3.8, then relativistic counter-streaming
between ingoing and outgoing fluids will undoubtedly take place, through gravitational waves if nothing else.
One way to allow relativistic counter-streaming is to let the speed of sound be the speed of light. This is
true in a massless scalar (= spin-0) field φ, which has an equation of state wφ = 1. Figure 17.7 shows a black
hole that accretes a charged, non-conducting fluid with this equation of state. The parameters are otherwise
the same as as in previous subsections: a mass of M• = 4 × 106 M⊙, an accretion rate of M• = 10−16, and
a black hole charge-to-mass of Q•/M• = 10−5. As the Figure shows, mass inflation takes place just above
17.3 Self-similar models of the interior structure of black holes 305
hori
zonm
ass
infl
atio
n
ρφ
ρe
|C|
1010 1020 1030 1040 105010−12010−11010−10010−9010−8010−7010−6010−5010−4010−3010−2010−10
10010101020
Radius r (Planck units)
(Pla
nck
unit
s)
Figure 17.7 Instead of a relativistic plasma, this shows a charged scalar field φ whose equation of statewφ = 1 means that the speed of sound equals the speed of light. The scalar field therefore supports relativisticcounter-streaming, as a result of which mass inflation occurs just above the erstwhile inner horizon. Themass is M• = 4 × 106 M⊙, the accretion rate M• = 10−16, the charge-to-mass Q•/M• = 10−5, and theconductivity is zero.
the place where the inner horizon would be. During mass inflation, the density ρφ and the Weyl scalar C
rapidly exponentiate up to the Planck scale and beyond.
One of the remarkable features of the mass inflation instability is that the smaller the accretion rate, the
more violent the instability. Figure 17.8 shows mass inflation in a black hole of charge-to-mass Q•/M• = 0.8
accreting a massless scalar field at rates M• = 0.01, 0.003, and 0.001. The charge-to-mass has been chosen
to be largish so that the inner horizon is not too far below the outer horizon, and the accretion rates have
been chosen to be large because otherwise the inflationary growth rate is too rapid to be discerned easily on
the graph. The density ρφ and Weyl scalar C exponentiate along with, and in proportion to, the interior
mass M , which increases as the radius r decreases as
M ∝ exp(− ln r/M•) . (17.27)
Physically, the scale of length of inflation is set by how close to the inner horizon infalling material approaches
before mass inflation begins. The smaller the accretion rate, the closer the approach, and consequently the
shorter the length scale of inflation.
306 The interiors of spherical black holes
hori
zon
innerhorizon
0.01 hori
zon
innerhorizon
0.003
hori
zon
innerhorizon
ρφ
0.001
ρe
|C|
.1 .2 .5 1 210−12010−11010−10010−9010−8010−7010−6010−5010−4010−3010−2010−10
10010101020
Radius r (geometric units)
(Pla
nck
unit
s)
Figure 17.8 The density ρφ and Weyl curvature scalar |C| inside a black hole accreting a massless scalar field.
The graph shows three cases, with mass accretion rates M• = 0.01, 0.003, and 0.001. The graph illustratesthat the smaller the accretion rate, the faster the density and curvature inflate. Mass inflation destroys theinner horizon: the dashed vertical line labeled “inner horizon” shows the position that the inner horizonwould have if mass inflation did not occur. The black hole mass is M• = 4 × 106 M⊙, the charge-to-mass isQ•/M• = 0.8, and the conductivity is zero.
17.3.10 Black hole accreting charged baryons and dark matter
No scalar field (massless or otherwise) has yet been observed in nature, although it is supposed that the
Higgs field is a scalar field, and it is likely that cosmological inflation was driven by a scalar field. Another
way to allow mass inflation in simple models is to admit not one but two fluids that can counter-stream
relativistically through each other. A natural possibility is to feed the black hole not only with a charged
relativistic fluid of baryons but also with neutral pressureless dark matter that streams freely through the
baryons. The charged baryons, being repelled by the electric charge of the black hole, become outgoing,
while the neutral dark matter remains ingoing.
Figure 17.9 shows that relativistic counter-streaming between the baryons and the dark matter causes
the center-of-mass density ρ and the Weyl curvature scalar C to inflate quickly up to the Planck scale and
beyond. The ratio of dark matter to baryonic density at the sonic point is ρd/ρb = 0.1, but otherwise the
parameters are the generic parameters of previous subsections: M• = 4 × 106 M⊙, M• = 10−16, wb = 0.32,
Q•/M• = 10−5, and zero conductivity. Almost all the center-of-mass energy ρ is in the counter-streaming
energy between the outgoing baryonic and ingoing dark matter. The individual densities ρb of baryons and
ρd of dark matter (and ρe of electromagnetic energy) increase only modestly.
As in the case of the massless scalar field considered in the previous subsection, §17.3.9, the smaller the
17.3 Self-similar models of the interior structure of black holes 307
hori
zonm
ass
infl
atio
n
ρd
ρbρe
|C|
ρ
dSb/dSBH
1010 1020 1030 1040 105010−12010−11010−10010−9010−8010−7010−6010−5010−4010−3010−2010−10
10010101020
Radius r (Planck units)
(Pla
nck
unit
s)
Figure 17.9 Back to the relativistic charged baryonic plasma, but now in addition the black hole accretesneutral pressureless uncharged dark matter, which streams freely through the baryonic plasma. The rel-ativistic counter-streaming produces mass inflation just above the erstwhile inner horizon. The mass isM• = 4 × 106 M⊙, the accretion rate M• = 10−16, the baryonic equation of state wb = 0.32, the charge-to-mass Q•/M• = 10−5, the conductivity is zero, and the ratio of dark matter to baryonic density at the outersonic point is ρd/ρb = 0.1.
accretion rate, the shorter the length scale of inflation. Not only that, but the smaller one of the ingoing or
outgoing streams is relative to the other, the shorter the length scale of inflation. Figure 17.10 shows a black
hole with three different ratios of the dark-matter-to-baryon density ratio at the sonic point, ρd/ρb = 0.3,
0.1, and 0.03, all with the same total accretion rate M• = 10−2. The smaller the dark matter stream, the
faster is inflation. The accretion rate M• and the dark-matter-to-baryon ratio ρd/ρb have been chosen to be
relatively large so that the inflationary growth rate is discernable easily on the graph.
Figure 17.10 shows that, as in Figure 17.9, almost all the center-of-mass energy is in the streaming energy
between the baryons and the dark matter. For one case, ρd/ρb = 0.3, Figure 17.10 shows the individual
densities ρb of baryons, ρd of dark matter, and ρe of electromagnetic energy, all of which remain tiny compared
to the streaming energy.
17.3.11 The black hole particle accelerator
The previous subsection, §17.3.10, showed that almost all the center-of-mass energy during mass inflation
is in the energy of counter-streaming. Thus the black hole acts like an extravagantly powerful particle
accelerator.
Mass inflation is an exponential instability. The nature of the black hole particle accelerator is that an
308 The interiors of spherical black holes
hori
zon
innerhorizon
ρd ρb
ρe |C|ρ
0.3
hori
zon
innerhorizon
ρe |C|ρ
0.1
hori
zon
innerhorizon
ρe |C|ρ
0.03.1 .2 .5 1 2
10−12010−11010−10010−9010−8010−7010−6010−5010−4010−3010−2010−10
10010101020
Radius r (geometric units)
(Pla
nck
unit
s)
Figure 17.10 The center-of-mass density ρ and Weyl curvature |C| inside a black hole accreting baryons
and dark matter at rate M• = 0.01. The graph shows three cases, with dark-matter-to-baryon ratio at thesonic point of ρd/ρb = 0.3, 0.1, and 0.03. The smaller the ratio of dark matter to baryons, the faster thecenter-of-mass density ρ and curvature C inflate. For the largest ratio, ρd/ρb = 0.3 (to avoid confusion,only this case is plotted), the graph also shows the individual proper densities ρb of baryons, ρd of darkmatter, and ρe of electromagnetic energy. During mass inflation, almost all the center-of-mass energy ρ isin the streaming energy: the proper densities of individual components remain small. The black hole massis M• = 4 × 106 M⊙, the baryonic equation of state is wb = 0.32, the charge-to-mass is Q•/M• = 0.8, andthe conductivity is zero.
individual particle spends approximately an equal interval of proper time being accelerated through each
decade of collision energy.
Each baryon in the black hole collider sees a flux ndur of dark matter particles per unit area per unit
time, where nd = ρd/md is the proper number density of dark matter particles in their own frame, and ur is
the radial component of the proper 4-velocity, the γv, of the dark matter through the baryons. The γ factor
in ur is the relavistic beaming factor: all frequencies, including the collision frequency, are speeded up by
the relativistic beaming factor γ. As the baryons accelerate through the collider, they spend a proper time
interval dτ/d lnur in each e-fold of Lorentz factor ur. The number of collisions per baryon per e-fold of ur is
the dark matter flux (ρd/md)ur, multiplied by the time dτ/d ln ur, multiplied by the collision cross-section
σ. The total cumulative number of collisions that have happened in the black hole particle collider equals
this multiplied by the total number of baryons that have fallen into the black hole, which is approximately
equal to the black hole mass M• divided by the mass mb per baryon. Thus the total cumulative number of
17.3 Self-similar models of the interior structure of black holes 309
100 1050 10100
10−2
10−1
1
10
Velocity u
Col
lisi
onra
teρ d
udτ
/dln
u(M
•/M
•)
0.03
0.01
0.003
10−16
Figure 17.11 Collision rate of the black hole particle accelerator per e-fold of velocity u (meaning γv),
expressed in units of the inverse black hole accretion time M•/M•. The models illustrated are the same as
those in Figure 17.10. The curves are labeled with their mass accretion rates: M• = 0.03, 0.01, 0.003, and10−16. Stars mark where the center-of-mass energy of colliding baryons and dark matter particles exceedsthe Planck energy, while disks show where the Weyl curvature scalar C exceeds the Planck scale.
collisions in the black hole collider is
number of collisions
e-fold of ur =M•mb
ρd
mdσur dτ
d lnur. (17.28)
Figure 17.11 shows, for several different accretion rates M•, the collision rate M•ρdurdτ/d lnur of the black
hole collider, expressed in units of the black hole accretion rate M•. This collision rate, multiplied by
M•σ/(mdmb), gives the number of collisions (17.28) in the black hole. In the units c = G = 1 being used
here, the mass of a baryon (proton) is 1 GeV ≈ 10−54 m. If the cross-section σ is expressed in canonical
accelerator units of femtobarns (1 fb = 10−43 m2) then the number of collisions (17.28) is
number of collisions
e-fold of ur = 1045( σ
1 fb
)
(
300 GeV2
mbmd
)
(
M•10−16
)
(
M•ρdurdτ/d lnur
0.03 M•
)
. (17.29)
Particle accelerators measure their cumulative luminosities in inverse femtobarns. Equation (17.29) shows
that the black hole accelerator delivers about 1045 femtobarns, and it does so in each e-fold of collision energy
up to the Planck energy and beyond.
310 The interiors of spherical black holes
17.4 Instability at outer horizon?
A number of papers have suggested that a magical phase transition at, or just outside, the outer horizon
prevents any horizon from forming. Is it true? For example, is there a mass inflation instability at the outer
horizon?
If the were a White Hole on the other side of the outer horizon, then indeed an object entering the outer
horizon would encounter an inflationary instability. But otherwise, no.
PART SEVEN
GENERAL RELATIVISTIC PERTURBATION THEORY
Concept Questions
1. Why do general relativistic perturbation theory using the tetrad formalism as opposed to the coordinate
approach?
2. Why is the tetrad metric γmn assumed fixed in the presence of perturbations?
3. Are the tetrad axes γγm fixed under a perturbation?
4. Is it true that the tetrad components ϕmn of a perturbation are (anti-)symmetric in m↔ n if and only if
its coordinate components ϕµν are (anti-)symmetric in µ↔ ν?
5. Does an unperturbed quantity, such as the unperturbed metric0
gµν , change under an infinitesimal coordi-
nate gauge transformation?
6. How can the vierbein perturbation ϕmn be considered a tetrad tensor field if it changes under an infinites-
imal coordinate gauge transformations?
7. What properties of the unperturbed spacetime allow decomposition of perturbations into independently
evolving Fourier modes?
8. What properties of the unperturbed spacetime allow decomposition of perturbations into independently
evolving scalar, vector, and tensor modes?
9. In what sense do scalar, vector, and tensor modes have spin 0, 1, and 2 respectively?
10. Tensor modes represent gravitational waves that, in vacuo, propagate at the speed of light. Do scalar
and vector modes also propagate at the speed of light in vacuo? If so, do scalar and vector modes also
constitute gravitational waves?
11. If scalar, vector, and tensor modes evolve independently, does that mean that scalar modes can exist and
evolve in the complete absence of tensor modes? If so, does it mean that scalar modes can propagate
causally, in vacuo at the speed of light, without any tensor modes being present?
12. Equation (20.74) defines the mass M of a body as what a distant observer would measure from its
gravitational potential. Similarly equation (20.82) defines the angular momentum L of a body as what a
distant observer would measure from the dragging of inertial frames. In what sense are these definitions
legitimate?
13. Can an observer far from a body detect the difference between the scalar potentials Ψ and Φ produced by
the body?
314 Concept Questions
14. If a gravitational wave is a wave of spacetime itself, distorting the very rulers and clocks that measure
spacetime, how is it possible to measure gravitational waves at all?
15. Have gravitational waves been detected?
16. If gravitational waves carry energy-momentum, then can gravitational waves be present in a region of
spacetime with vanishing energy-momentum tensor, Tmn = 0?
What’s important?
1. Getting your brain around coordinate and tetrad gauge transformations.
2. A central aim of general relativistic perturbation theory is to identify the coordinate and tetrad gauge-
invariant perturbations, since only these have physical meaning.
3. A second central aim is to classify perturbations into independently evolving modes, to the extent that
this is possible.
4. In background spacetimes with spatial translation and rotation symmetry, which includes Minkowski space
and the Friedmann-Roberston-Walker metric of cosmology, modes decompose into independently evolving
scalar (spin-0), vector (spin-1), and tensor (spin-2) modes. In background spacetimes without spatial
translation and rotation symmetry, such as black holes, scalar, vector, and tensor modes scatter off the
curvature of space, and therefore mix with each other.
5. In background spacetimes with spatial translation and rotation symmetry, there are 6 algebraic combina-
tions of metric coefficients that are coordinate and tetrad gauge-invariant, and therefore represent physical
perturbations. There are 2 scalar modes, 2 vector modes, and 2 tensor modes. A spin-m mode varies as
eimχ where χ is the rotational angle about the spatial wavevector k of the mode.
6. In background spacetimes without spatial translation and rotation symmetry, the coordinate and tetrad
gauge-invariant perturbations are not algebraic combinations of the metric coefficients, but rather com-
binations that involve first and second derivatives of the metric coefficients. Gravitational waves are
described by the Weyl tensor, which can be decomposed into 5 complex components, with spin 0, ±1,
and ±2. The spin-±2 components describe propagating gravitational waves, while the spin-0 and spin-±1
components describe the non-propagating gravitational field near a source.
7. The preeminent application of general relativistic perturbation theory is to cosmology. Coupled with
physics that is either well understood (such as photon-electron scattering) or straightforward to model
even without a deep understanding (such as the dynamical behavior of non-baryonic dark matter and
dark energy), the theory has yielded predictions that are in spectacular agreement with observations of
fluctuations in the CMB and in the large scale distribution of galaxies and other tracers of the distribution
of matter in the Universe.
18
Perturbations and gauge transformations
This chapter sets up the basics equations that define perturbations to an arbitrary spacetime in the tetrad
formalism of general relativity, and it examines the effect of tetrad and coordinate gauge transformations
on those perturbations. The perturbations are supposed to be small, in the sense that quantities quadratic
in the perturbations can be neglected. The formalism set up in this chapter provides a foundation used in
subsequent chapters.
18.1 Notation for perturbations
A 0 (zero) overscript signifies an unperturbed quantity, while a 1 (one) overscript signifies a perturbation.
No overscript means the full quantity, including both unperturbed and perturbed parts. An overscript is
attached only where necessary. Thus if the unperturbed part of a quantity is zero, then no overscript is
needed, and none is attached.
18.2 Vierbein perturbation
Let the vierbein perturbation ϕmn be defined so that the perturbed vierbein is
emµ = (δn
m + ϕmn)
0
enµ , (18.1)
with corresponding inverse
emµ = (δm
n − ϕnm)
0
enµ . (18.2)
Since the perturbation ϕmn is already of linear order, to linear order its indices can be raised and lowered
with the unperturbed metric, and transformed between tetrad and coordinate frames with the unperturbed
vierbein. In practice it proves convenient to work with the covariant tetrad-frame components ϕmn of the
vierbein pertubation
ϕmn = γnlφml . (18.3)
18.3 Gauge transformations 317
The perturbation ϕmn can be regarded as a tetrad tensor field defined on the unperturbed background.
18.3 Gauge transformations
The vierbein perturbation ϕmn has 16 degrees of freedom, but only 6 of these degrees of freedom correspond
to real physical perturbations, since 6 degrees of freedom are associated with arbitrary infinitesimal changes
in the choice of tetrad, which is to say arbitrary infinitesimal Lorentz transformations, and a further 4 degrees
of freedom are associated with arbitrary infinitesimal changes in the coordinates.
In the context of perturbation theory, these infinitesimal tetrad and coordinate transformations are called
gauge transformations. Real physical perturbations are perturbations that are gauge-invariant under
both tetrad and coordinate gauge transformations.
18.4 Tetrad metric assumed constant
In the tetrad formalism, tetrad axes γγm are introduced as locally inertial (or other physically motivated)
axes attached to an observer. The axes enable quantities to be projected into the frame of the observer.
In a spacetime buffeted by perturbations, it is natural for an observer to cling to the rock provided by the
locally inertial (or other) axes, as opposed to allowing the axes to bend with the wind. For example, when
a gravitational wave goes by, the tidal compression and rarefaction causes the proper distance between two
freely falling test masses to oscillate. It is natural to choose the tetrad so that it continues to measure proper
times and distances in the perturbed spacetime.
In these notes on general relativistic perturbation theory, the tetrad metric will be taken to be constant
everywhere, and unchanged by a perturbation
γmn =0
γmn = constant . (18.4)
For example, if the tetrad is orthonormal, then the tetrad metric is constant, the Minkowski metric ηmn.
However, the tetrad could also be some other tetrad for which the tetrad metric is constant, such as a spinor
tetrad (§12.1.1), or a Newman-Penrose tetrad (§12.2.1).
18.5 Perturbed coordinate metric
The perturbed coordinate metric is
gµν = γmn em
µen
ν
= γkl(δkm − ϕm
k)0
emµ(δl
n − ϕnl)
0
enν
=0
gµν − (ϕµν + ϕνµ) . (18.5)
318 Perturbations and gauge transformations
Thus the perturbation of the coordinate metric depends only on the symmetric part of the vierbein pertur-
bation ϕmn, not the antisymmetric part
1
gµν = − (ϕµν + ϕνµ) . (18.6)
18.6 Tetrad gauge transformations
Under an infinitesimal tetrad transformation, the covariant vierbein perturbations ϕmn transform as
ϕmn → ϕmn + ǫmn , (18.7)
where ǫmn is the generator of a Lorentz transformation, which is to say an arbitrary antisymmetric tensor
(Exercise 11.2). Thus the antisymmetric part ϕmn − ϕnm of the covariant perturbation ϕmn is arbitrarily
adjustable through an infinitesimal tetrad transformation, while the symmetric part ϕmn + ϕnm is tetrad
gauge-invariant.
It is easy to see when a quantity is tetrad gauge-invariant: it is tetrad gauge-invariant if and only if it
depends only on the symmetric part of the vierbein perturbation, not on the antisymmetric part. Evidently
the perturbation (18.6) to the coordinate metric gµν is tetrad gauge-invariant. This is as it should be, since
the coordinate metric gµν is a coordinate-frame quantity, independent of the choice of tetrad frame.
If only tetrad gauge-invariant perturbations are physical, why not just discard tetrad perturbations (the
antisymmetric part of ϕmn) altogether, and work only with the tetrad gauge-invariant part (the symmetric
part of ϕmn)? The answer is that tetrad-frame quantities such as the tetrad-frame Einstein tensor do change
under tetrad gauge transformations (infinitesimal Lorentz transformations of the tetrad). It is true that
the only physical perturbations of the Einstein tensor are those combinations of it that are tetrad gauge-
invariant. But in order to identify these tetrad gauge-invariant combinations, it is necessary to carry through
the dependence on the non-tetrad-gauge-invariant part, the antisymmetric part of ϕmn.
Much of the professional literature on general relativistic perturbation theory works with the traditional
coordinate formalism, as opposed to the tetrad formalism. The term “gauge-invariant” then means coor-
dinate gauge-invariant, as opposed to both coordinate and tetrad gauge-invariant. This is fine as far as it
goes: the coordinate approach is perfectly able to identify physical perturbations versus gauge perturbations.
However, there still remains the problem of projecting the perturbations into the frame of an observer, so
ultimately the issue of perturbations of the observer’s frame, tetrad perturbations, must be faced.
Concept question 18.1 In perturbation theory, can tetrad gauge transformations be non-infinitesimal?
18.7 Coordinate gauge transformations 319
18.7 Coordinate gauge transformations
A coordinate gauge transformation is a transformation of the coordinates xµ by an infinitesimal shift ǫµ
xµ → x′µ = xµ + ǫµ . (18.8)
You should not think of this as shifting the underlying spacetime around; rather, it is just a change of the
coordinate system, which leaves the underlying spacetime unchanged. Because the shift ǫµ is, like the vierbein
perturbations ϕmn, already of linear order, its indices can be raised and lowered with the unperturbed metric,
and transformed between coordinate and tetrad frames with the unperturbed vierbein. Thus the shift ǫµ
can be regarded as a vector field defined on the unperturbed background. The tetrad components ǫm of the
shift ǫµ are
ǫm =0
emµ ǫ
µ . (18.9)
Physically, the tetrad-frame shift ǫm is the shift measured in locally inertial coordinates
ξm → ξ′m = ξm + ǫm . (18.10)
18.8 Coordinate gauge transformation of a coordinate scalar
Under a coordinate transformation (18.8), a coordinate-frame scalar Φ(x) remains unchanged
Φ(x)→ Φ′(x′) = Φ(x) . (18.11)
Here the scalar Φ′(x′) is evaluated at position x′, which is the same as the original physical position x since
all that has changed is the coordinates, not the physical position. However, in perturbation theory, quantities
are evaluated at coordinate position x, not x′. The value of Φ at x is related to that at x′ by
Φ′(x) = Φ′(x′ − ǫ) = Φ′(x′)− ǫκ ∂Φ′
∂xκ. (18.12)
Since ǫκ is a small quantity, and Φ′ differs from Φ by a small quantity, the last term ǫκ∂Φ′/∂xκ in equa-
tion (18.16) can be replaced by ǫκ∂Φ/∂xκ to linear order. Putting equations (18.11) and (18.12) together
shows that the scalar Φ changes under a coordinate gauge transformation (18.8) as
Φ(x)→ Φ′(x) = Φ(x)− ǫκ ∂Φ
∂xκ. (18.13)
The transformation (18.13) can also be written
Φ(x)→ Φ′(x) = Φ(x) + LǫΦ , (18.14)
where Lǫ is the Lie derivative, §18.13.
320 Perturbations and gauge transformations
18.9 Coordinate gauge transformation of a coordinate vector or tensor
A similar argument applies to coordinate vectors and tensors. Under a coordinate transformation (18.8), a
coordinate-frame 4-vector Aµ(x) transforms in the usual way as
Aµ(x)→ A′µ(x′) = Aκ(x)∂x′µ
∂xκ= Aµ(x) +Aκ(x)
∂ǫµ
∂xκ. (18.15)
As in the scalar case, the vector A′µ(x′) is evaluated at position x′, which is the same as the original physical
position x since all that has changed is the coordinates, not the physical position. Again, in perturbation
theory, quantities are evaluated at coordinate position x, not x′. The value of A′µ at x is related to that at
x′ by
A′µ(x) = A′µ(x′ − ǫ) = A′µ(x′)− ǫκ ∂A′µ
∂xκ. (18.16)
The last term ǫκ∂A′µ/∂xκ in equation (18.16) can be replaced by ǫκ∂Aµ/∂xκ to linear order. Putting
equations (18.15) and (18.16) together shows that the 4-vector Aµ changes under a coordinate gauge trans-
formation (18.8) as
Aµ(x)→ A′µ(x) = Aµ(x) +Aκ ∂ǫµ
∂xκ− ǫκ ∂A
µ
∂xκ. (18.17)
The transformation (18.17) can also be written
Aµ(x)→ A′µ(x) = Aµ(x) + LǫAµ , (18.18)
where Lǫ is the Lie derivative, §18.13.
More generally, under a coordinate gauge transformation (18.8), a coordinate tensor Aκλ...µν... transforms as,
equation (18.32),
Aκλ...µν...(x)→ A′κλ...
µν...(x) = Aκλ...µν...(x) + LǫA
κλ...µν... . (18.19)
18.10 Coordinate gauge transformation of a tetrad vector
A tetrad-frame 4-vector Am is a coordinate-invariant quantity, and therefore acts like a coordinate scalar,
equation (18.13), under a coordinate gauge transformation (18.8)
Am(x)→ A′m(x) = Am(x)− ǫκ ∂Am
∂xκ= Am(x) − ǫk∂kA
m . (18.20)
The change −ǫk∂kAm is a coordinate tensor (specifically, a coordinate scalar), but it is not a tetrad tensor.
More generally, a tetrad-frame tensor Akl...mn... transforms under a coordinate gauge transformation (18.8)
as
Akl...mn...(x)→ A′kl...
mn...(x) = Akl...mn...(x)− ǫa∂aA
kl...mn... . (18.21)
Again, the change −ǫa∂aAkl...mn... is a coordinate tensor (a coordinate scalar), but not a tetrad tensor.
18.11 Coordinate gauge transformation of the vierbein 321
18.11 Coordinate gauge transformation of the vierbein
The inverse vierbein emµ equals the scalar product of the tetrad and coordinate axes, em
µ = γγm · gµ.
Therefore the transformation of the vierbein under a coordinate gauge transformation (18.8) follows from
the transformations of γγm and gµ. The tetrad axes γγm transform in accordance with (18.20) as
γγm → γγ′m = γγm − ǫk∂kγγm
= γγm + ǫkΓmnkγγn . (18.22)
The coordinate axes gµ transform in accordance with (18.32) and (18.34) as (including torsion Sνµκ)
gµ → g′µ = gµ + Lǫgµ
= gµ − gνDµǫν − ǫνDνgµ − gνS
νµκǫ
κ
= gµ −Dµ (gνǫν)− gνS
νµκǫ
κ
= gµ − γγn
(
Dmǫn + Sn
mkǫk) 0
emµ , (18.23)
where the third line follows from the second because the axes gν are by definition covariantly constant,
Dµgν = 0. It follows from (18.22) and (18.23) that the inverse vierbein emµ transforms under an infinitesimal
coordinate gauge transformation (18.8) as
emµ → e′mµ = γγ′m · g′
µ
= emµ +
(
−Dnǫm − Sm
nkǫk + Γm
nkǫk) 0
enµ . (18.24)
From equation (18.24) and the definition (18.2) of the vierbein perturbations ϕmn, it follows that the vierbein
perturbations transform under a coordinate gauge transformation (18.8) as
ϕmn → ϕ′mn = ϕmn + ∂mǫn − (Γknm + Γnmk − Snmk) ǫk , (18.25)
in which ǫk are the tetrad components of the coordinate shift, and Γkmn are tetrad connection coefficients.
If torsion vanishes, Snmk = 0, as general relativity assumes, then the transfomation (18.25) of the vierbein
perturbations under a coordinate gauge transformation reduces to
ϕmn → ϕ′mn = ϕmn + ∂mǫn − (Γknm + Γnmk) ǫk . (18.26)
18.12 Coordinate gauge transformation of the metric
The tetrad metric γmn transforms under an infinitesimal coordinate gauge transformation (18.8) as
γmn → γ′mn = γmn − (Γmnk + Γnmk)ǫk = γmn , (18.27)
where the last expression is true because the tetrad metric γmn is being assumed constant, equation (18.4),
in which case Γmnk + Γnmk = ∂kγnm = 0.
322 Perturbations and gauge transformations
The coordinate metric gµν transforms under an infinitesimal coordinate gauge transformation (18.8) as
gµν → g′µν = gµν + Lǫgµν = gµν − (Dµǫν +Dνǫµ)− (Sµνκ + Sνµκ)ǫκ . (18.28)
18.13 Lie derivative
The change in the coordinate 4-vector Aµ on the right hand side of equation (18.17) is called the Lie
derivative of Aµ along the direction ǫκ, and it is designated by the operator Lǫ
LǫAµ ≡ Aκ ∂ǫ
µ
∂xκ− ǫκ∂A
µ
∂xκ. (18.29)
The Lie derivative has the important property of being a tensor, which is one reason that it merits a special
name. As its name suggests, the Lie derivative acts like a derivative: it is linear, and it satisfies the Leibniz
rule. Translating from ordinary partial derivatives to covariant derivatives yields the following expression
for the Lie derivative in covariant form
LǫAµ = AκDκǫ
µ − ǫκDκAµ +AκǫλSµ
κλ is a coordinate vector , (18.30)
where Sµκλ is the torsion. If torsion vanishes, as GR assumes, then the Lie derivative of a 4-vector is
LǫAµ = AκDκǫ
µ − ǫκDκAµ is a coordinate vector . (18.31)
More generally, under a coordinate gauge transformation (18.8), a coordinate tensor Aκλ...µν... transforms as
Aκλ...µν...(x)→ A′κλ...
µν...(x) = Aκλ...µν...(x) + LǫA
κλ...µν... (18.32)
where the Lie derivative is defined by
LǫAκλ...µν... ≡ Aαλ...
µν...
∂ǫκ
∂xα+Aκα...
µν...
∂ǫλ
∂xα... −Aκλ...
αν...
∂ǫα
∂xµ−Aκλ...
µα...
∂ǫα
∂xν− ǫα ∂A
κλ...µν...
∂xα. (18.33)
In covariant form, equation (18.33) is
LǫAκλ...µν... = Aαλ...
µν...Dαǫκ +Aκα...
µν...Dαǫλ ... −Aκλ...
αν...Dµǫα −Aκλ...
µα...Dνǫα ... − ǫαDαA
κλ...µν... (18.34)
+(
Aαλ...µν...S
καβ +Aκα...
µν...Sλαβ ... −Aκλ...
αν...Sαµβ −Aκλ...
µα...Sανβ
)
ǫβ is a coordinate tensor ,
which in the case of vanishing torsion, as GR assumes, reduces to
LǫAκλ...µν... = Aαλ...
µν...Dαǫκ +Aκα...
µν...Dαǫλ ... −Aκλ...
αν...Dµǫα −Aκλ...
µα...Dνǫα ... − ǫαDαA
κλ...µν... is a coordinate tensor .
(18.35)
19
Scalar, vector, tensor decomposition
In the particular case that the unperturbed spacetime is spatially homogeneous and isotropic, which includes
not only Minkoswki space but also the important case of the cosmological Friedmann-Robertson-Walker
metric, perturbations decompose into independently evolving scalar (spin-0), vector (spin-1), and tensor
(spin-2) modes.
Similarly to Fourier decomposition, decomposition into scalar, vector, and tensor modes is non-local, in
principle requiring knowledge of perturbation amplitudes simultaneously throughout all of space. In practical
problems however, an adequate decomposition is possible as long as the scales probed are sufficiently larger
than the wavelengths of the modes probed. Ultimately, the fact that an adequate decomposition is possible
is a consequence of the fact that gravitational fluctuations in the real Universe appear to converge at the
cosmological horizon, so that what happens locally is largely independent of what is happening far away.
19.1 Decomposition of a vector in flat 3D space
Theorem: In flat 3-dimensional space, a 3-vector field w(x) can be decomposed uniquely (subject to the
boundary condition that w vanishes sufficiently rapidly at infinity) into a sum of scalar and vector parts
w = ∇w‖scalar
+ w⊥vector
. (19.1)
In this context, the term vector signifies a 3-vector w⊥ that is transverse, that is to say, it has vanishing
divergence,
∇ ·w⊥ = 0 . (19.2)
Here ∇ ≡ ∂/∂x ≡ ∇i ≡ ∂/∂xi is the gradient in flat 3D space. The scalar and vector parts are also known
as spin-0 and spin-1, or gradient and curl, or longitudinal and transverse. The scalar part ∇w‖ contains 1
degree of freedom, while the vector part w⊥ contains 2 degrees of freedom. Together they account for the 3
degrees of freedom of the vector w.
324 Scalar, vector, tensor decomposition
Proof: Take the divergence of equation (19.1)
∇ ·w = ∇2w‖ . (19.3)
The operator∇2 on the right hand side of equation (19.3) is the 3D Laplacian. The solution of equation (19.3)
is
w‖(x) = −∫
∇′ ·w(x′)
|x′ − x|d3x′
4π. (19.4)
The solution (19.4) is valid subject to boundary conditions that the vector w vanish sufficiently rapidly at
infinity. In cosmology, the required boundary conditions, which are set at the Big Bang, are apparently
satisfied because fluctuations at the Big Bang were small. Equation (19.1) then immediately implies that
the vector part is w⊥ = w −∇w‖.
19.2 Fourier version of the decomposition of a vector in flat 3D space
When the background has some symmetry, it is natural to expand perturbations in eigenmodes of the
symmetry. If the background space is flat, then it is translation symmetric. Eigenmodes of the translation
operator ∇ are Fourier modes.
A function a(x) in flat 3D space and its Fourier transform a(k) are related by (the disposition of factors
of 2π in the following definition follows the convention most commonly adopted by cosmologists)
a(k) =
∫
a(x)eik·x d3x , a(x) =
∫
a(k)e−ik·x d3k
(2π)3. (19.5)
You may not be familiar with the practice of using the same symbol a in both real and Fourier space; but
a is the same vector in Hilbert space, with components ax = a(x) in real space, and ak = a(k) in Fourier
space.
Taking the gradient ∇ in real space is equivalent to multiplying by −ik in Fourier space
∇→ −ik . (19.6)
Thus the decomposition (19.1) of the 3D vector w translates into Fourier space as
w = −ikw‖scalar
+ w⊥vector
, (19.7)
where the vector part w⊥ satisfies
k ·w⊥ = 0 . (19.8)
In other words, in Fourier space the scalar part ∇w‖ of the vector w is the part parallel (longitudinal) to
the wavevector k, while the vector part w⊥ is the part perpendicular (transverse) to the wavevector k.
19.3 Decomposition of a tensor in flat 3D space 325
19.3 Decomposition of a tensor in flat 3D space
Similarly, the 9 components of a 3× 3 spatial matrix hij can be decomposed into 3 scalars, 2 vectors, and 1
tensor:
hij = δij φscalar
+∇i∇jhscalar
+ εijk∇khscalar
+∇ihjvector
+∇j hivector
+ hTij
tensor
. (19.9)
In this context, the term tensor signifies a 3× 3 matrix hTij that is traceless, symmetric, and transverse:
hT ii = 0 , hT
ij = hTji , ∇ih
Tij = 0 . (19.10)
The transverse-traceless-symmetric matrix hTij has two degrees of freedom. The vector components hi and
hi are by definition transverse,
∇ihi = ∇ihi = 0 . (19.11)
The tildes on h and hi simply distinguish those symbols; the tildes have no other significance. The trace of
the 3× 3 matrix hij is
hii = 3φ+∇2h . (19.12)
Spinor decomposition.
20
Flat space background
General relativistic perturbation theory is simplest in the case that the unperturbed background space
is Minkowski space. In Cartesian coordinates xµ ≡ t, x, y, z, the unperturbed coordinate metric is the
Minkowski metric0
gµν = ηµν . (20.1)
In this chapter the tetrad is taken to be orthonormal, and aligned with the unperturbed coordinate axes, so
that the unperturbed vierbein is the unit matrix
0
emµ = δµ
m . (20.2)
Let overdot denote partial differentiation with respect to time t,
overdot ≡ ∂
∂t, (20.3)
and let ∇ denote the spatial gradient
∇ ≡ ∂
∂x≡ ∇i ≡
∂
∂xi. (20.4)
Sometimes it will also be convenient to use ∇m to denote the 4-dimensional spacetime derivative
∇m ≡ ∂
∂t,∇
. (20.5)
20.1 Classification of vierbein perturbations
The aims of this section are two-fold. First, decompose perturbations into scalar, vector, and tensor parts.
Second, identify the coordinate and tetrad gauge-invariant perturbations. It will be found, equations (20.13),
that there are 6 coordinate and tetrad gauge-invariant perturbations, comprising 2 scalars Ψ and Φ, 1 vector
Wi containing 2 degrees of freedom, and 1 tensor hij containing 2 degrees of freedom.
20.1 Classification of vierbein perturbations 327
The vierbein perturbations ϕmn decompose into 6 scalars, 4 vectors, and 1 tensor
ϕtt = ψscalar
, (20.6a)
ϕti = ∇iwscalar
+ wivector
, (20.6b)
ϕit = ∇iwscalar
+ wivector
, (20.6c)
ϕij = δij Φscalar
+∇i∇jhscalar
+ εijk∇khscalar
+∇ihjvector
+∇j hivector
+ hijtensor
. (20.6d)
The tildes on w and h simply distinguish those symbols; the tildes have no other significance. The vector
components are by definition transverse (have vanishing divergence), while the tensor component hij is by
definition traceless, symmetric, and transverse. For a single Fourier mode whose wavevector k is taken
without loss of generality to lie in the z-direction, equations (20.6) are
ϕmn =
ψ wx wy ∇zw
wx Φ + hxx hxy +∇zh ∇zhx
wy hxy −∇zh Φ− hxx ∇zhy
∇zw ∇zhx ∇zhy Φ +∇2zh
. (20.7)
To identify coordinate gauge-invariant quantities, it is necessary to consider infinitesimal coordinate gauge
transformations (18.8). The tetrad-frame components ǫm of the coordinate shift of the coordinate gauge
transformation decompose into 2 scalars and 1 vector
ǫm = ǫtscalar
, ∇iǫscalar
+ ǫivector
. (20.8)
In the flat space background space being considered the coordinate gauge transformation (18.26) of the
vierbein perturbation simplifies to
ϕmn → ϕ′mn = ϕmn +∇mǫn . (20.9)
In terms of the scalar, vector, and tensor potentials introduced in equations (20.6), the gauge transforma-
tions (20.9) are
ϕtt → ψ + ǫtscalar
, (20.10a)
ϕti → ∇i(w + ǫ)scalar
+ (wi + ǫi)vector
, (20.10b)
ϕit → ∇i(w + ǫt)scalar
+ wivector
, (20.10c)
ϕij → δij Φscalar
+∇i∇j(h+ ǫ)scalar
+ εijk∇khscalar
+∇i(hj + ǫj)vector
+∇j hivector
+ hijtensor
. (20.10d)
328 Flat space background
Equations (20.10a) imply that under an infinitesimal coordinate gauge transformation
ψ → ψ + ǫt , (20.11a)
w → w + ǫ , wi → wi + ǫi , (20.11b)
w → w + ǫt , wi → wi , (20.11c)
Φ→ Φ , h→ h+ ǫ , h→ h , hi → hi + ǫi , hi → hi , hij → hij . (20.11d)
Eliminating the coordinate shift ǫm from the transformations (20.11) yields 12 coordinate gauge-invariant
combinations of the potentials
ψ − ˙w , w − h , wi − hi , wi , Φ , h , hi , hij . (20.12)
Physical perturbations are not only coordinate but also tetrad gauge-invariant. A quantity is tetrad gauge-
invariant if and only if it depends only on the symmetric part of the vierbein pertubations, not on the
antisymmetric part, §18.6. There are 6 combinations of the coordinate gauge-invariant perturbations (20.12)
that are symmetric, and therefore not only coordinate but also tetrad gauge-invariant. These 6 coordinate
and tetrad gauge-invariant perturbations comprise 2 scalars, 1 vector, and 1 tensor
Ψscalar
≡ ψ − w − ˙w + h , (20.13a)
Φscalar
, (20.13b)
Wivector
≡ wi + wi − hi − ˙hi , (20.13c)
hijtensor
. (20.13d)
Since only the 6 tetrad and coordinate gauge-invariant potentials Ψ, Φ, Wi, and hij have physical signifi-
cance, it is legitimate to choose a particular gauge, a set of conditions on the non-gauge-invariant potentials,
arranged to simplify the equations, or to bring out some physical aspect. Three gauges considered later are
harmonic gauge (§20.7), Newtonian gauge (§20.9), and synchronous gauge (§20.10). However, for the next
several sections, no gauge will be chosen: the exposition will continue to be completely general.
20.2 Metric, tetrad connections, and Einstein and Weyl tensors
This section gives expressions in a completely general gauge for perturbed quantities in flat background
Minkowski space.
20.2 Metric, tetrad connections, and Einstein and Weyl tensors 329
The perturbed coordinate metric gµν , equation (18.5), is
gtt = −(1 + 2ψ) , (20.14a)
gti = −∇i(w + w)− (wi + wi) , (20.14b)
gij = δij (1− 2 Φ)− 2∇i∇jh−∇i(hj + hj)−∇j(hi + hi)− 2 hij . (20.14c)
The coordinate metric is tetrad gauge-invariant, but not coordinate gauge-invariant.
The perturbed tetrad connections Γkmn are
Γtit = −∇i(ψ − ˙w) + ˙wi , (20.15a)
Γtij = δij Φ−∇i∇j(w − h)− 12 (∇iWj +∇jWi) +∇jwi + hij , (20.15b)
Γijt = 12 (∇iWj −∇jWi)−
∂
∂t(εijl∇lh−∇ihj +∇j hi) , (20.15c)
Γijk = (δjk∇i − δik∇j)Φ−∇k(εijl∇lh−∇ihj +∇j hi) +∇ihjk −∇jhik . (20.15d)
Being purely tetrad-frame quantities, the tetrad connections are automatically coordinate gauge-invariant.
However, they are not tetrad gauge-invariant, as is evident from the fact that they (all) depend on antisym-
metric parts of the vierbein perturbations ϕmn.
One of the advantages of working with tetrads is that tetrad-frame quantities such as the tetrad connec-
tions Γkmn and the tetrad-frame Riemann tensor Rklmn are by construction independent of the choice of
coordinates, and are therefore automatically coordinate gauge-invariant. In the tetrad formalism, you do
not have to work too hard (is that really ever true?) to construct coordinate gauge-invariant combinations
of the vierbein perturbations ϕmn: the tetrad-frame connections and Riemann tensor will automatically
give you the coordinate gauge-invariant combinations. You can check that in the present case the tetrad
connections (20.15) depend only on, and on all 12 of, the coordinate gauge-invariant combinations (20.12).
The perturbed tetrad-frame Einstein tensor Gmn is
Gtt = 2∇2Φscalar
, (20.16a)
Gti = 2∇iΦscalar
+ 12 ∇2Wivector
, (20.16b)
Gij = 2 δij Φscalar
− (∇i∇j − δij∇2)(Ψ− Φ)scalar
+ 12 (∇iWj +∇jWi)
vector
− hijtensor
, (20.16c)
where is the d’Alembertian, the 4-dimensional wave operator
≡ −∇m∇m =∂2
∂t2−∇2 . (20.17)
Being a tetrad-frame quantity, the tetrad-frame Einstein tensor is automatically coordinate gauge-invariant.
Equations (20.16) show that the tetrad-frame Einstein tensor Gmn is also tetrad gauge-invariant, since it
depends only on the tetrad-gauge invariant combinations (20.13) of the vierbein perturbations. The property
330 Flat space background
that the Einstein tensor is tetrad as well as coordinate gauge-invariant is a feature of empty background
space, and does not persist to more general spacetimes, such as the Friedmann-Robertson-Walker spacetime.
In a frame with the wavector k taken along the z-axis, the perturbed Einstein tensor is
Gmn =
2∇2zΦ
12 ∇2
zWx12 ∇2
zWy 2∇zΦ
12 ∇2
zWx 2 Φ +∇2z(Ψ− Φ)−h+ −h×
12 ∇zWx
12 ∇2
zWy −h× 2 Φ +∇2z(Ψ− Φ) + h+
12 ∇zWy
2∇zΦ12 ∇zWx
12 ∇zWy 2 Φ
(20.18)
where h+ and h× are the two polarizations of gravitational waves, discussed further in §20.14,
h+ ≡ hxx = − hyy , h× ≡ hxy = hyx . (20.19)
The tetrad-frame complexified Weyl tensor is
Ctitj = 14 (∇i∇j − 1
3 δij∇2)(Ψ + Φ)scalar
+ 18
[
− (∇iWj +∇jWi) + i(εikl∇j + εjkl∇i)∇kWl
]
vector
+ 14
[
hij − εiklεjmn∇k∇mhln − i(εikl∇khjl + εjkl∇khil)]
tensor
. (20.20)
Like the tetrad-frame Einstein tensor, the tetrad-frame Weyl tensor is both coordinate and tetrad gauge-
invariant, depending only on the coordinate and tetrad gauge-invariant potentials Ψ, Φ, Wi, and hij .
20.3 Spinor components of the Einstein tensor
Scalar, vector, and tensor perturbations correspond respectively to perturbations of spin 0, 1, and 2. An
object has spin m if it is unchanged by a rotation of 2π/m about a prescribed direction. In perturbed
Minkowski space, the prescribed direction is the direction of the wavevector k in the Fourier decomposition
of the modes. The spin components may be projected out by working in a spinor tetrad, §12.1.1.
In a frame where the wavevector k is taken along the z-axis, the spinor components of the perturbed
Einstein tensor Gmn are (compare equations (20.16))
Gtt = 2∇2zΦ
spin-0
, Gtz = 2∇zΦspin-0
, Gzz = 2 Φspin-0
, (20.21a)
G+− −Gzz = ∇2z(Ψ− Φ)spin-0
, (20.21b)
Gt± = 12 ∇2
zW±
spin-±1
, Gz± = 12 ∇zW±
spin-±1
, (20.21c)
G±± = −h±±
spin-±2, (20.21d)
20.4 Too many Einstein equations? 331
where W± are the spin-±1 components of the vector perturbation Wi
W± = 1√2
(Wx ± iWy) , (20.22)
and h±± are the spin-±2 components of the tensor perturbation hij
h±± = hxx ± i hxy = h+ ± i h× . (20.23)
The spin +2 and −2 components h++ and h−− of the tensor perturbation are called the right- and left-handed
circular polarizations. The spin +2 and −2 circular polarizations h++ and h−− have respective shapes ei2χ
and e−i2χ, under a right-handed rotation by angle χ about the z-axis, which may be compared to the cos 2χ
and sin 2χ shapes of the linear polarizations h+ and h×.
20.4 Too many Einstein equations?
The Einstein equations are as usual (units c = G = 1)
Gmn = 8πTmn . (20.24)
There are 10 Einstein equations, but the Einstein tensor (20.16) depends on only 6 independent potentials:
the two scalars Ψ and Φ, the vector Wi, and the tensor hij . The system of Einstein equations is thus over-
complete. Why? The answer is that 4 of the Einstein equations enforce conservation of energy-momentum,
and can therefore be considered as governing the evolution of the energy-momentum as opposed to being
equations for the gravitational potentials. For example, the form of equations (20.16a) and (20.16b) for Gtt
and Gti enforces conservation of energy
DmGmt = 0 , (20.25)
while the form of equations (20.16b) and (20.16c) for Gti and Gij enforces conservation of momentum
DmGmi = 0 . (20.26)
Normally, the equations governing the evolution of the energy-momentum Tmna of each species a of mass-
energy would be set up so as to ensure overall conservation of energy-momentum. If this is done, then
the conservation equations (20.25) and (20.26) can be regarded as redundant. Since equations (20.25) and
(20.26) are equations for the time evolution of Gtt and Gti, one might think that the Einstein equations
for Gtt and Gti would become redundant, but this is not quite true. In fact the Einstein equations for
Gtt and Gti impose constraints that must be satisfied on the initial spatial hypersurface. Conservation
of energy-momentum guarantees that those constraints will continue to be satisfied on subsequent spatial
hypersurfaces, but still the initial conditions must be arranged to satisfy the constraints. Because the Einstein
equations for Gtt and Gti must be satisfied as constraints on the initial conditions, but thereafter can be
ignored, the equations are called constraint equations. The Einstein equation for Gtt is called the energy
constraint, or Hamiltonian constraint. The Einstein equations for Gti are called the momentum constraints.
332 Flat space background
Exercise 20.1 Energy and momentum constraints. Confirm the argument of this section. Suppose
that the spatial Einstein equations are true, Gij = 8πT ij. Show that if the time-time and time-space
Einstein equations Gtm = 8πT tm are initially true, then conservation of energy-momentum implies that
these equations must necessarily remain true at all times. [Hint: Conservation of energy-momentum requires
that DmTmn = 0, and the Bianchi identities require that the Einstein tensor satisfies DmG
mn = 0, so
Dm(Gmn − 8πTmn) = 0 . (20.27)
By expanding out these equations in full, or otherwise, show that the solution satisfying Gij − 8πT ij = 0 at
all times, and Gtm − 8πT tm = 0 initially, is Gtm − 8πT tm = 0 at all times.]
Concept question 20.2 Which Einstein equations are redundant? It has been argued in this
section that, if the energy-momentum tensor Tmn is arranged to satisfy energy-conservation DmTmn = 0 as
it should, then the time-time and time-space Einstein equations must be satisfied by the initial conditions,
but thereafter become redundant. Question: Can any 4 of the 10 Einstein equations be dropped, or just the
time-time and time-space Einstein equations?
20.5 Action at a distance?
The tensor component of the Einstein equations shows that, in a vacuum Tmn = 0, the tensor perturbations
hij propagate at the speed of light, satisfying the wave equation
hij = 0 . (20.28)
The tensor perturbations represent propagating gravitational waves.
It is to be expected that scalar and vector perturbations would also propagate at the speed of light, yet
this is not obvious from the form of the Einstein tensor (20.16). Specifically, there are 4 components of the
Einstein tensor (20.16) that apparently depend only on spatial derivatives, not on time derivatives. The 4
corresponding Einstein equations are
∇2Φ = 4πTttscalar
, (20.29a)
∇2Wi = 16πTtivector
, (20.29b)
∇2(Ψ− Φ) = − 8πQijTijscalar
, (20.29c)
where Qij in equation (20.29c) is the quadrupole operator defined below, equation (20.87). These condi-
tions must be satisfied everywhere at every instant of time, giving the impression that signals are traveling
instantaneously from place to place.
20.6 Comparison to electromagnetism 333
20.6 Comparison to electromagnetism
The previous two sections §§20.4, 20.5 brought up two issues:
1. There are 10 Einstein equations, but only 6 independent gauge-invariant potentials Ψ, Φ, Wi, and hij .
The additional 4 Einstein equations serve to enforce conservation of energy-momentum.
2. Only 2 of the gauge-invariant potentials, the tensor potentials hij , satisfy causal wave equations. The
remaining 4 gauge-invariant potentials Ψ, Φ, and Wi, satisfy equations (20.29) that depend on the instan-
taneous distribution of energy-momentum throughout space, on the face of it violating causality.
These facts may seem surprising, but in fact the equations of electromagnetism have a similar structure, as
will now be shown. In this section, the spacetime is assumed to be flat Minkowski space. The discussion
in this section is based in part on the exposition by Bertschinger (1995 “Cosmological Dynamics,” 1993 Les
Houches Lectures, arXiv:astro-ph/9503125).
In accordance with the usual procedure, the electromagnetic field may be defined in terms of an elec-
tromagnetic 4-potential Am, whose time and spatial parts constitute the scalar potential φ and the vector
potential A:
Am ≡ φ,A . (20.30)
The electric and magnetic fields E and B may be defined in terms of the potentials φ and A by
E ≡ −∇φ− ∂A
∂t, (20.31a)
B ≡∇×A . (20.31b)
Given their definition (20.31), the electric and magnetic fields automatically satisfy the two source-free
Maxwell’s equations
∇ ·B = 0 , (20.32a)
∇×E +∂B
∂t= 0 . (20.32b)
The remaining two Maxwell’s equations, the sourced ones, are
∇ ·E = 4πq , (20.33a)
∇×B − ∂E
∂t= 4πj , (20.33b)
where q and j are the electric charge and current density, the time and space components of the electric
4-current density jm
jm ≡ q, j . (20.34)
The electromagnetic potentials φ and A are not unique, but rather are defined only up to a gauge transfor-
mation by some arbitrary gauge field χ
φ→ φ+∂χ
∂t, A→ A−∇χ . (20.35)
334 Flat space background
The gauge transformation (20.35) evidently leaves the electric and magnetic fields E and B, equations (20.31),
invariant.
Following the path of previous sections, §20.1 and thereafter, decompose the vector potential A into its
scalar and vector parts
A = ∇A‖scalar
+ A⊥vector
, (20.36)
in which the vector part by definition satisfies the transversality condition ∇ · A⊥ = 0. Under a gauge
transformation (20.35), the potentials transform as
φ→ φ+∂χ
∂t, (20.37a)
A‖ → A‖ − χ , (20.37b)
A⊥ → A⊥ . (20.37c)
Eliminating the gauge field χ yields 3 gauge-invariant potentials, comprising 1 scalar Φ, and 1 vector A⊥containing 2 degrees of freedom:
Φscalar
≡ φ+∂A‖∂t
, (20.38a)
A⊥vector
. (20.38b)
This shows that the electromagnetic field contains 3 independent degrees of freedom, consisting of 1 scalar
and 1 vector.
Concept question 20.3 Are gauge-invariant potentials Lorentz invariant? The potentials Φ and
A⊥, equations (20.38), are by construction gauge-invariant, but is this construction Lorentz invariant? Do
Φ and A⊥ constitute the components of a 4-vector?
In terms of the gauge-invariant potentials Φ and A⊥, equations (20.38), the electric and magnetic fields
are
E = −∇Φ− ∂A⊥∂t
, (20.39a)
B = ∇×A⊥ . (20.39b)
The sourced Maxwell’s equations (20.33) thus become, in terms of Φ and A⊥,
−∇2Φscalar
= 4πqscalar
, (20.40a)
∇Φscalar
+ A⊥vector
= 4π∇j‖scalar
+ 4πj⊥vector
, (20.40b)
20.6 Comparison to electromagnetism 335
where ∇j‖ and j⊥ are the scalar and vector parts of the current density j. Equations (20.40) bear a striking
similarity to the Einstein equations (20.16). Only the vector part A⊥ satisfies a wave equation,
A⊥ = 4πj⊥ , (20.41)
while the scalar part Φ satisfies an instantaneous equation (20.40a) that seemingly violates causality. And just
as Einstein’s equations (20.16) enforce conservation of energy-momentum, so also Maxwell’s equations (20.40)
enforce conservation of electric charge
∂q
∂t+ ∇ · j = 0 , (20.42)
or in 4-dimensional form
∇mjm = 0 . (20.43)
The fact that only the vector part A⊥ satisfies a wave equation (20.41) reflects physically the fact that
electromagnetic waves are transverse, and they contain only two propagating degrees of freedom, the vector,
or spin ±1, components.
Why do Maxwell’s equations (20.40) have this structure? Although equation (20.41) appears to be a local
wave equation for the vector part A⊥ of the potential sourced by the vector part j⊥ of the current, in fact
the wave equation is non-local because the decomposition of the potential and current into scalar and vector
parts is non-local (it involves the solution of a Laplacian equation, eq. (19.3)). It is only the sum j = ∇j‖+j⊥of the scalar and vector parts of the current density that is local. Therefore, the Maxwell’s equation (20.40b)
must have a scalar part to go along with the vector part, such that the source on the right hand side, the
current density j, is local. Given this Maxwell equation (20.40b), the Maxwell equation (20.40a) then serves
precisely to enforce conservation of electric charge, equation (20.42).
Just as it is possible to regard the Einstein equations (20.16a) and (20.16b) as constraint equations
whose continued satisfaction is guaranteed by conservation of energy-momentum, so also the Maxwell equa-
tion (20.40a) for Φ can be regarded as a constraint equation whose continued satisfaction is guaranteed
by conservation of electric charge. For charge conservation (20.42) coupled with the spatial Maxwell equa-
tion (20.40b) ensures that
∂
∂t
(
4πq +∇2Φ)
= 0 , (20.44)
the solution of which, subject to the condition that 4πq +∇2Φ = 0 initially, is 4πq +∇2Φ = 0 at all times,
which is precisely the Maxwell equation (20.40a).
Exercise 20.4 Is it possible to discard the scalar part of the spatial Maxwell equation (20.40b),
rather than equation (20.40a) for Φ? Project out the scalar part of equation (20.40b) by taking its
divergence,
∇2(
4πj‖ − Φ)
= 0 . (20.45)
336 Flat space background
Argue that the Maxwell equation (20.40a), coupled with charge conservation (20.42), ensures that equa-
tion (20.45) is true, subject to boundary condition that the current j vanish sufficiently rapidly at spatial
infinity, in accordance with the decomposition theorem of §19.1.
Since only gauge-invariant quantities have physical significance, it is legitimate to impose any condition
on the gauge field χ. A gauge in which the potentials φ and A individually satisfy wave equations is Lorenz
(not Lorentz!) gauge, which consists of the Lorentz-invariant condition
∇mAm = 0 . (20.46)
Under a gauge transformation (20.35), the left hand side of equation (20.46) transforms as
∇mAm → ∇mA
m + χ , (20.47)
and the Lorenz gauge condition (20.46) can be accomplished as a particular solution of the wave equation
for the gauge field χ. In terms of the potentials φ and A‖, the Lorenz gauge condition (20.46) is
∂φ
∂t+∇2A‖ = 0 . (20.48)
In Lorenz gauge, Maxwell’s equations (20.40) become
φ = 4πq , (20.49a)
A = 4πj , (20.49b)
which are manifestly wave equations for the potentials φ and A.
Does the fact that the potentials φ and A in one particular gauge, Lorenz gauge, satisfy wave equations
necessarily guarantee that the electric and magnetic fields E and B satisfy wave equations? Yes, because it
follows from the definitions (20.31) of E and B that if the potentials φ and A satisfy wave equations, then
so also must the fields E and B themselves; but the fields E and B are gauge-invariant, so if they satisfy
wave equations in one gauge, then they must satisfy the same wave equations in any gauge.
In electromagnetism, the most physical choice of gauge is one in which the potentials φ and A coincide
with the gauge-invariant potentials Φ and A⊥, equations (20.38). This gauge, known as Coulomb gauge,
is accomplished by setting
A‖ = 0 , (20.50)
or equivalently
∇ ·A = 0 . (20.51)
The gravitational analogue of this gauge is the Newtonian gauge discussed in the next section but one, §20.9.
Does the fact that in Lorenz gauge the potentials φ and A propagate at the speed of light (in the absence
of sources, jm = 0) imply that the gauge-invariant potentials Φ and A⊥ propagate at the speed of light?
No. The gauge-invariant potentials Φ and A⊥, equations (20.38), are related to the Lorenz gauge potentials
φ and A by a non-local decomposition.
20.7 Harmonic gauge 337
20.7 Harmonic gauge
The fact that all locally measurable gravitational perturbations do propagate causally, at the speed of light
in the absence of sources, can be demonstrated by choosing a particular gauge, harmonic gauge, equa-
tion (20.52), which can be considered an analogue of the Lorenz gauge of electromagnetism, equation (20.46).
In harmonic gauge, all 10 of the tetrad gauge-variant (i.e. symmetric) combinations ϕmn+ϕnm of the vierbein
perturbations satisfy wave equations (20.56), and therefore propagate causally. This does not imply that
the scalar, vector, and tensor components of the vierbein perturbations individually propagate causally, be-
cause the decomposition into scalar, vector, and tensor modes is non-local. In particular, the coordinate and
tetrad-gauge invariant potentials Ψ, Φ, Wi, and hij defined by equations (20.13) do not propagate causally.
The situation is entirely analogous to that of electromagnetism, §20.6, where in Lorenz gauge the potentials
φ and A propagate causally, equations (20.49), yet the gauge-invariant potentials Φ and A⊥ defined by
equations (20.38) do not.
Harmonic gauge is the set of 4 coordinate conditions
∇m(ϕmn + ϕnm)−∇nϕmm = 0 . (20.52)
The conditions (20.52) are arranged in a form that is tetrad gauge-invariant (the conditions depend only on
the symmetric part of ϕmn). The quantities on the left hand side of equations (20.52) transform under a
coordinate gauge transformation, in accordance with (20.9), as
∇m(ϕmn + ϕnm)−∇nϕmm → ∇m(ϕmn + ϕnm)−∇nϕm
m + ǫn . (20.53)
The change ǫn resulting from the coordinate gauge transformation is the 4-dimensional wave operator
acting on the coordinate shift ǫn. Indeed, the harmonic gauge conditions (20.52) follow uniquely from the
requirements (a) that the change produced by a coordinate gauge transformation be ǫn, as suggested by
the analogous electromagnetic transformation (20.47), and (b) that the conditions be tetrad gauge-invariant.
The harmonic gauge conditions (20.52) can be accomplished as a particular solution of the wave equation for
the coordinate shift ǫn. In terms of the potentials defined by equations (20.6) and (20.13), the 4 harmonic
gauge conditions (20.52) are
Ψ + 3Φ + (w + w − h) = 0 , (20.54a)
Wi + (hi + hi) = 0 , (20.54b)
−Ψ + Φ + h = 0 , (20.54c)
or equivalently
− 4 Φ = (w + w) , (20.55a)
− Wi = (hi + hi) , (20.55b)
Ψ− Φ = h . (20.55c)
Substituting equations (20.55) into the Einstein tensor Gmn leads, after some calculation, to the result that
338 Flat space background
in harmonic gauge
− 12 (ϕmn + ϕnm) = Rmn , (20.56)
where Rmn = Gmn− 12 ηmnG is the Ricci tensor. Equation (20.56) shows that in harmonic gauge, all tetrad
gauge-invariant (i.e. symmetric) combinations ϕmn+ϕnm of the vierbein potentials propagate causally, at
the speed of light in vacuo, Rmn = 0. Although the result (20.56) is true only in a particular gauge,
harmonic gauge, it follows that all quantities that are (coordinate and tetrad) gauge-invariant, and that can
be constructed from the vierbein potentials ϕmn and their derivatives (and are therefore local), must also
propagate at the speed of light.
20.8 What is the gravitational field?
In electromagnetism, the electromagnetic fields are the electric field E and the magnetic field B. These fields
have the property that they are gauge-invariant, and measurable locally. The electromagnetic potentials Φ
and A⊥, equations (20.38), are gauge-invariant, but they are not measurable locally.
What are the analogous gauge-invariant and locally measurable quantities for the gravitational field in
perturbed Minkowski space? The answer is, the Weyl tensor Cklmn, the trace-free or tidal part of the
Riemann tensor, the expression (20.20) for which depends only on the coordinate and tetrad gauge-invariant
potentials.
20.9 Newtonian (Copernican) gauge
If the unperturbed background is Minkowski space, then the most physical gauge is one in which the 6
perturbations retained coincide with the 6 coordinate and tetrad gauge-invariant perturbations (20.13).
This gauge is called Newtonian gauge. Because in Newtonian gauge the perturbations are precisely the
physical perturbations, if the perturbations are physically weak (small), then the perturbations in Newtonian
gauge will necessarily be small.
I think Newtonian gauge should be called Copernican gauge. Even though the solar system is a highly non-
linear system, from the perspective of general relativity it is a weakly perturbed gravitating system. Applied
to the solar system, Newtonian gauge effectively keeps the coordinates aligned with the classical Sun-centred
Copernican coordinate frame. By contrast, the coordinates of synchronous gauge (§20.10), which are chosen
to follow freely-falling bodies, would quickly collapse or get wound up by orbital motions if applied to the
solar system, and would cease to provide a useful description.
Newtonian gauge sets
w = w = wi = h = h = hi = hi = 0 , (20.57)
20.10 Synchronous gauge 339
so that the retained perturbations are the 6 coordinate and tetrad gauge-invariant perturbations (20.13)
Ψscalar
= ψ , (20.58a)
Φscalar
, (20.58b)
Wivector
= wi , (20.58c)
hijtensor
. (20.58d)
The Newtonian line-element is, in a form that keeps the Newtonian tetrad manifest,
ds2 = −[
(1 + Ψ) dt]2
+ δij[
(1− Φ)dxi − hikdx
k −W idt][
(1− Φ)dxj − hjl dx
l −W jdt]
, (20.59)
which reduces to the Newtonian metric
ds2 = − (1 + 2 Ψ) dt2 − 2Wi dt dxi +[
δij(1 − 2 Φ)− 2 hij
]
dxi dxj . (20.60)
Since scalar, vector, and tensor perturbations evolve independently, it is legitimate to consider each in
isolation. For example, if one is interested only in scalar perturbations, then it is fine to keep only the
scalar potentials Ψ and Φ non-zero. Furthermore, as discussed in §20.13, since the difference Ψ − Φ in
scalar potentials is sourced by anisotropic relativistic pressure, which is typically small, it is often a good
approximation to set Ψ = Φ.
The tetrad-frame 4-velocity of a person at rest in the tetrad frame is by definition um = 1, 0, 0, 0, and
the corresponding coordinate 4-velocity uµ is, in Newtonian gauge,
uµ = etµ = 1−Ψ,Wi . (20.61)
This shows that Wi can be interpreted as a 3-velocity at which the tetrad frame is moving through the coor-
dinates. This is the “dragging of inertial frames” discussed in §20.12. The proper acceleration experienced
by a person at rest in the tetrad frame, with tetrad 4-velocity um = 1, 0, 0, 0, is
Dui
Dτ= utDtu
i = ut(
∂tui + Γi
ttut)
= Γitt = ∇iΨ . (20.62)
This shows that the “gravity,” or minus the proper acceleration, experienced by a person at rest in the tetrad
frame is minus the gradient of the potential Ψ.
Concept question 20.5 If the decomposition into scalar, vector, and tensor modes is non-local, how can
it be legimate to consider the evolution of the modes in isolation from each other?
20.10 Synchronous gauge
One of the earliest gauges used in general relativistic perturbation theory, and still (in its conformal version)
widely used in cosmology, is synchronous gauge. As will be seen below, equations (20.69) and (20.70),
340 Flat space background
synchronous gauge effectively chooses a coordinate system and tetrad that is attached to the locally inertial
frames of freely falling observers. This is fine as long as the observers move only slightly from their initial
positions, but the coordinate system will fail when the system evolves too far, even if, as in the solar system,
the gravitational perturbations remain weak and therefore treatable in principle with perturbation theory.
Synchronous gauge sets the time components ϕmn with m = t or n = t of the vierbein perturbations to
zero
ψ = w = w = wi = wi = 0 , (20.63)
and makes the additional tetrad gauge choices
h = hi = 0 , (20.64)
with the result that the retained perturbations are the spatial perturbations
Φscalar
, hscalar
, hivector
, hijtensor
. (20.65)
In terms of these spatial perturbations, the gauge-invariant perturbations (20.13) are
Ψscalar
= h , (20.66a)
Φscalar
, (20.66b)
Wivector
= − hi , (20.66c)
hijtensor
. (20.66d)
The synchronous line-element is, in a form that keeps the synchronous tetrad manifest,
ds2 = − dt2 + δij[
(1−Φ)dxi − (∇k∇ih+∇khi + hi
k)dxk][
(1−Φ)dxj − (∇l∇jh+∇lhj + hj
l )dxl]
, (20.67)
which reduces to the synchronous metric
ds2 = − dt2 + [(1 − 2 Φ)δij − 2∇i∇jh−∇ihj −∇jhi − 2 hij ] dxi dxj . (20.68)
In synchronous gauge, a person at rest in the tetrad frame has coordinate 4-velocity
uµ = etµ = 1, 0, 0, 0 , (20.69)
so that the tetrad rest frame coincides with the coordinate rest frame. Moreover a person at rest in the
tetrad frame is freely falling, which follows from the fact that the acceleration experienced by a person at
rest in the tetrad frame is zero
Duk
Dτ= ut
(
∂tuk + Γk
ttut)
= Γktt = 0 , (20.70)
in which ∂tuk = 0 because the 4-velocity at rest in the tetrad frame is constant, uk = 1, 0, 0, 0, and Γk
tt = 0
from equations (20.15a) with the synchronous gauge choices (20.63) and (20.64).
20.11 Newtonian potential 341
20.11 Newtonian potential
The next few sections examine the physical meaning of each of the gauge-invariant potentials Ψ, Φ, Wi, and
hij by looking at the potentials at large distances produced by a finite body containing energy-momentum,
such as the Sun.
Einstein’s equations Gmn = 8πTmn applied to the time-time component Gtt of the Einstein tensor, equa-
tion (20.16a), imply Poisson’s equation
∇2Φ = 4πρ , (20.71)
where ρ is the mass-energy density
ρ ≡ Ttt . (20.72)
The solution of Poisson’s equation (20.71) is
Φ(x) = −∫
ρ(x′) d3x′
|x′ − x| . (20.73)
Consider a finite body, for example the Sun, whose energy-momentum is confined within a certain region.
Define the mass M of the body to be the integral of the mass-energy density ρ,
M ≡∫
ρ(x′) d3x′ . (20.74)
Equation (20.74) agrees with what the definition of the mass M would be in the non-relativistic limit, and
as seen below, equation (20.77), it is what a distant observer would infer the mass of the body to be based
on its gravitational potential Φ far away. Thus equation (20.74) can be taken as the definition of the mass
of the body even when the energy-momentum is relativistic. Choose the origin of the coordinates to be at
the centre of mass, meaning that∫
x′ ρ(x′) d3x′ = 0 . (20.75)
Consider the potential Φ at a point x far outside the body. Expand the denominator of the integral on the
right hand side of equation (20.73) as a Taylor series in 1/x where x ≡ |x|
1
|x′ − x| =1
x
∞∑
ℓ=0
(
x′
x
)ℓ
Pℓ(x · x′) =1
x+
x · x′
x2+ ... (20.76)
where Pℓ(µ) are Legendre polynomials. Then
Φ(x) = − 1
x
∫
ρ(x′) d3x′ − 1
x2x ·∫
x′ ρ(x′) d3x′ −O(x−3)
= − M
x−O(x−3) . (20.77)
Equation (20.77) shows that the potential far from a body goes as Φ = −M/x, reproducing the usual
Newtonian formula.
342 Flat space background
20.12 Dragging of inertial frames
In Newtonian gauge, the vector potential W ≡Wi is the velocity at which the locally inertial tetrad frame
moves through the coordinates, equation (20.61). This is called the dragging of inertial frames. As shown
below, a body of angular momentum L drags frames around it with an angular velocity that goes to 2L/x3
at large distances x.
Einstein’s equations applied to the vector part of the time-space component Gti of the Einstein tensor,
equation (20.16b), imply
∇2W = − 16πf , (20.78)
where W ≡Wi is the gauge-invariant vector potenial, and f is the vector part of the energy flux T ti
f ≡ fi = f i ≡ T ti
vector= −Tti
vector. (20.79)
The solution of equation (20.78) is
W (x) = 4
∫
f(x′) d3x′
|x′ − x| . (20.80)
As in the previous section, §20.11, consider a finite body, such as the Sun, whose energy-momentum is
confined within a certain region. Work in the rest frame of the body, defined to be the frame where the
energy flux f integrated over the body is zero,
∫
f(x′) d3x′ = 0 . (20.81)
Define the angular momentum L of the body to be
L ≡∫
x′ × f(x′) d3x′ . (20.82)
Equation (20.82) agrees with what the definition of angular momentum would be in the non-relativistic limit,
where the mass-energy flux of a mass density ρ moving at velocity v is f = ρv. As will be seen below, the
angular momentum (20.82) is what a distant observer would infer the angular momentum of the body to be
based on the potential W far away, and equation (20.82) can be taken to be the definition of the angular
momentum of the body even when the energy-momentum is relativistic. As will be proven momentarily,
equation (20.83), the integral∫
x′ifj(x′) d3x′ is antisymmetric in ij. To show this, write fj = εjkl∇kφl for
some potential φl, which is valid because fj is the vector (curl) part of the energy flux. Then
∫
x′ifj(x′) d3x′ =
∫
x′iεjkl∇′kφl(x
′) d3x′ = −∫
εjklφl(x′)∇′
kx′i d
3x′ =
∫
εijlφl(x′) d3x′ , (20.83)
where the third expression follows from the second by integration by parts, the surface term vanishing
because of the assumption that the energy-momentum of the body is confined within a certain region.
20.13 Quadrupole pressure 343
Taylor expanding equation (20.80) using equation (20.76) gives
W (x) =4
x
∫
f(x′) d3x′ +4
x2
∫
(x · x′)f(x′) d3x+O(x−3)
=2
x2
∫
[(x · x′)f(x′)− (x · f(x′))x′] d3x+O(x−3)
=2
x2L× x +O(x−3) , (20.84)
where the first integral on the right hand side of the first line of equation (20.84) vanishes because the frame
is the rest frame of the body, equation (20.81), and the second integral on the right hand side of the first line
equals the first integral on the second line thanks to the antisymmetry of∫
x′f(x′) d3x, equation (20.83).
The vector potential W ≡ Wi points in the direction of rotation, right-handedly about the axis of angular
momentum L. Equation (20.84) says that a body of angular momemtum L drags frames around it at angular
velocity Ω at large distances x
W = Ω× x , Ω =2L
x3. (20.85)
20.13 Quadrupole pressure
Einstein’s equations applied to the part of the Einstein tensor (20.16c) involving Ψ− Φ imply
∇2(Ψ− Φ) = − 8πQijTij , (20.86)
where Qij is the quadrupole operator (an integro-differential operator) defined by
Qij ≡ 32 ∇i∇j ∇−2 − 1
2 δij , (20.87)
with ∇−2 the inverse spatial Laplacian operator. In Fourier space, the quadrupole operator is
Qij = 32 kikj − 1
2 δij . (20.88)
The quadrupole operator Qij yields zero when acting on δij , and the Laplacian operator ∇2 when acting on
∇i∇j
Qijδij = 0 , Qij∇i∇j = ∇2 . (20.89)
The solution of equation (20.86) is
Ψ− Φ = −∫ [
3
2
(xi − x′i)(xj − x′j)|x− x′|2 − 1
2δij
]
Tij(x′) d3x′
|x− x′| . (20.90)
At large distance in the z-direction from a finite body
Ψ− Φ = − 1
x
∫
[
Tzz − 12 (Txx + Tyy)
]
d3x′ +O(x−2) . (20.91)
Equation (20.86) shows that the source of the difference Ψ − Φ between the two scalar potentials is the
344 Flat space background
quadrupole pressure. Since the quadrupole pressure is small if either there are no relativistic sources, or any
relativistic sources are isotropic, it is often a good approximation to set Ψ = Φ. An exception is where there
is a significant anisotropic relativistic component. For example, the energy-momentum tensor of a static
electric field is relativistic and anisotropic. However, this is still not enough to ensure that Ψ differs from Φ:
as found in Exercise 20.6, if the energy-momentum of a body is spherically symmetric, then Ψ−Φ vanishes
outside (but not inside) the body.
One situation where the difference between Ψ and Φ is appreciable is the case of freely-streaming neutrinos
at around the time of recombination in cosmology. The 2008 analysis of the CMB by the WMAP team claims
to detect a non-zero value of Ψ− Φ from a slight shift in the third acoustic peak.
Exercise 20.6 Argue that the traceless part of the energy-momentum tensor of a spherically symmetric
distribution must take the form
Tij(r) =(
rirj − 13 δij
)(
p(r) − p⊥(r))
, (20.92)
where p(r) and p⊥(r) are the radial and transverse pressures at radius r. From equation (20.90), show that
Ψ− Φ at radial distance x from the centre of a spherically symmetric distribution is
Ψ(x)− Φ(x) = −∫ ∞
x
(r2 − x2)(
p(r)− p⊥(r)) 4πdr
r. (20.93)
Notice that the integral is over r > x, that is, only energy-momentum outside radius x produces non-
vanishing Ψ− Φ. In particular, if the body has finite extent, then Ψ− Φ vanishes outside the body.
20.14 Gravitational waves
The tensor perturbations hij describe propagating gravitational waves. The two independent components of
the tensor perturbations describe two polarizations. The two components are commonly designated h+ and
h×, equations (20.19). Gravitational waves induce a quadrupole tidal oscillation transverse to the direction
of propagation, and the subscripts + and × represent the shape of the quadrupole oscillation, as illustrated
by Figure 20.1. The h+ polarization has a cos 2χ shape, while the h× polarization has a sin 2χ shape, where
χ is the azimuthal angle with respect to the y-axis about the direction x of propagation.
Einstein’s equations applied to the tensor component of the spatial Einstein tensor (20.16c) imply that
gravitational waves are sourced by the tensor component of the energy-momentum
hij = − 8π Tijtensor
. (20.94)
The solution of the wave equation (20.94) can be obtained from the Green’s function of the d’Alembertian
wave operator . The Green’s function is by definition the solution of the wave equation with a delta-
function source. There are retarded solutions, which propagate into the future along the future light cone,
20.14 Gravitational waves 345
Figure 20.1 The two polarizations of gravitational waves. The (top) polarization h+ has a cos 2χ shapeabout the direction of propagation (into the paper), while the (bottom) polarization h× has a sin 2χ shape.A gravitational wave causes a system of freely falling test masses to oscillate relative to a grid of points afixed proper distance apart.
and advanced solutions, which propagate into the past along the past light cone. In the present case, the
solutions of interest are the retarded solutions, since these represent gravitational waves emitted by a source.
Because of the time and space translation symmetry of the d’Alembertian, the delta-function source of the
Green’s function can without loss of generality be taken at the origin t = x = 0. Thus the Green’s function
F is the solution of
F = δ4(x) , (20.95)
where δ4(x) ≡ δ(t)δ3(x) is the 4-dimensional Dirac delta-function. The solution of equation (20.95) subject
to retarded boundary conditions is (a standard exercise in mathematics) the retarded Green’s function
F =δ(x− t)Θ(t)
4πx, (20.96)
where x ≡ |x| and Θ(t) is the Heaviside function, Θ(t) = 0 for t < 0 and Θ(t) = 1 for t ≥ 0. The solution of
the sourced gravitational wave equation (20.94) is thus
hij(t,x) = − 2
∫ Tij(t′,x′)
tensor
d3x′
|x′ − x| , (20.97)
346 Flat space background
where t′ is the retarded time
t′ ≡ t− |x′ − x| , (20.98)
which lies on the past light cone of the observer, and is the time at which the source emitted the signal. The
solution (20.97) resembles the solution of Poisson’s equation, except that the source is evaluated along the
past light cone of the observer.
As in §§20.11 and 20.12, consider a finite body, whose energy-momentum is confined within a certain region,
and which is a source of gravitational waves. The Hulse-Taylor binary pulsar is a fine example. Far from the
body, the leading order contribution to the tensor potential hij is, from the multipole expansion (20.76),
hij(t,x) = − 2
x
∫
Tij(t′,x′)
tensor
d3x′ . (20.99)
The integral (20.99) is hard to solve in general, but there is a simple solution for gravitational waves whose
wavelengths are large compared to the size of the body. To obtain this solution, first consider that conser-
vation of energy-momentum implies that
∂2T tt
∂t2−∇i∇jT
ji =∂
∂t
(
∂T tt
∂t+∇iT
ti
)
−∇i
(
∂T ti
∂t+∇jT
ji
)
= 0 . (20.100)
Multiply by xixj and integrate∫
xixj ∂2T tt
∂t2d3x =
∫
xixj ∇k∇lTkl d3x =
∫
T kl∇k∇l(xixj) d3x = 2
∫
T ij d3x , (20.101)
where the third expression follows from the second by a double integration by parts. For wavelengths that
are long compared to the size of the body, the first expression of equations (20.101) is∫
xixj∂2T tt
∂t2d3x ≈ ∂2
∂t2
∫
xixj Ttt d3x =
∂2Iij∂t2
(20.102)
where Iij is the second moment of the mass
Iij ≡∫
xixj Ttt d3x . (20.103)
The tensor (spin-2) part of the energy-momentum is trace-free. The trace-free part –Iij of the second moment
Iij is the quadrupole moment of the mass distribution (this definition is conventional, but differs by a factor
of 2/3 from what is called the quadrupole moment in spherical harmonics)
–Iij ≡ Iij − 13 δij I
kk =
∫
(xixj − 13 δij x
2)T tt d3x . (20.104)
Substituting the last expression of equations (20.101) into equation (20.99) gives the quadrupole formula for
gravitational radiation at wavelengths long compared to the size of the emitting body
hij(t,x) = − 1
x–Iij
tensor
. (20.105)
20.15 Energy-momentum carried by gravitational waves 347
If the gravitational wave is moving in the z-direction, then the tensor components of the quadrupole moment–Iij are
–I+ = 12 (Ixx − Iyy) , –I× = 1
2 (Ixy + Iyx) . (20.106)
20.15 Energy-momentum carried by gravitational waves
The gravitational wave equation (20.28) in empty space appears to describe gravitational waves propagating
in a region where the energy-momentum tensor Tmn is zero. However, gravitational waves do carry energy-
momentum, just as do other kinds of waves, such as electromagnetic waves. The energy-momentum is
quadratic in the tensor perturbation hij , and so vanishes to linear order.
To determine the energy-momentum in gravitational waves, calculate the Einstein tensor Gmn to second
order, imposing the vacuum conditions that the unperturbed and linear parts of the Einstein tensor vanish
0
Gmn =1
Gmn = 0 . (20.107)
The parts of the second-order perturbation that depend on the tensor perturbation hij are, in a frame where
the wavevector k is along the z-axis,
2
Gtt = − (hij)(hij) +
1
4
( ∂2
∂t2+∇2
z
)
h2 , (20.108a)
2
Gtz = − (hij)(∇zhij) +
1
2
∂
∂t∇zh
2 , (20.108b)
2
Gzz = − (∇zhij)(∇zhij) +
1
4
( ∂2
∂t2+∇2
z
)
h2 , (20.108c)
where
h2 ≡ hijhij = 2(h2
++ h2
×) = 2h++h−−. (20.109)
Being tetrad-frame quantities, the expressions (20.108) are automatically coordinate gauge-invariant, and
they are also tetrad gauge-invariant since they depend only on the (coordinate and) tetrad gauge-invariant
perturbation hij . The rightmost set of terms on the right hand side of each of equations (20.108) are total
derivatives (with respect to time t or space z). These terms yield surface terms when integrated over a
region, and tend to average to zero when integrated over a region much larger than a wavelength. On the
other hand, the leftmost set of terms on the right hand side of each of equations (20.108) do not average
to zero; for example, the terms for Gtt and Gzz are negative everywhere, being minus a sum of squares. A
negative energy density? The interpretation is that these terms are to be taken over to the right hand side
348 Flat space background
of the Einstein equations, and re-interpreted as the energy-momentum T gwmn in gravitational waves
T gwtt ≡
1
8π
[
(hij)(hij)− 1
4
( ∂2
∂t2+∇2
z
)
h2
]
, (20.110a)
T gwtz ≡
1
8π
[
(hij)(∇zhij)− 1
2
∂
∂t∇zh
2
]
, (20.110b)
T gwzz ≡
1
8π
[
(∇zhij)(∇zhij)− 1
4
( ∂2
∂t2+∇2
z
)
h2
]
. (20.110c)
The terms involving total derivatives, although they vanish when averaged over a region larger than many
wavelengths, ensure that the energy-momentum T gwmn in gravitational waves satisfies conservation of energy-
momentum in the flat background space
∇mT gwmn = 0 . (20.111)
Averaged over a region larger than many wavelengths, the energy-momentum in gravitational waves is
〈T gwmn〉 =
1
8π(∇mhij)(∇nh
ij) . (20.112)
Equation (20.112) may also be written explicitly as a sum over the two linear or circular polarizations
〈T gwmn〉 =
1
4π[(∇mh+)(∇nh+) + (∇mh×)(∇nh×)]
=1
8π[(∇mh++)(∇nh−−) + (∇nh++)(∇mh−−)] . (20.113)
PART EIGHT
COSMOLOGICAL PERTURBATIONS
Concept Questions
1. Why do the wavelengths of perturbations in cosmology expand with the Universe, whereas perturbations
in Minkowski space do not expand?
2. What does power spectrum mean?
3. Why is the power spectrum a good way to characterize the amplitude of fluctuations?
4. Why is the power spectrum of fluctuations of the Cosmic Microwave Background (CMB) plotted as a
function of harmonic number?
5. What causes the acoustic peaks in the power spectrum of fluctuations of the CMB?
6. Are there acoustic peaks in the power spectrum of matter (galaxies) today?
7. What sets the scale of the first peak in the power spectrum of the CMB? [What sets the physical scale?
Then what sets the angular scale?]
8. The odd peaks (including the first peak) in the CMB power spectrum are compression peaks, while the
even peaks are rarefaction peaks. Why does a rarefaction produce a peak, not a trough?
9. Why is the first peak the most prominent? Why do higher peaks generally get progressively weaker?
10. The third peak is about as strong as the second peak? Why?
11. The matter power spectrum reaches a maximum at a scale that is slightly larger than the scale of the first
baryonic acoustic peak. Why?
12. The physical density of species x at the time of recombination is proportional to Ωxh2 where Ωx is the
ratio of the actual to critical density of species x at the present time, and h ≡ H0/100 kms−1 Mpc−1 is
the present-day Hubble constant. Explain.
13. How does changing the baryon density Ωbh2 affect the CMB power spectrum?
14. How does changing the non-baryonic cold dark matter density Ωch2, without changing the baryon density
Ωbh2, affect the CMB power spectrum?
15. What effects do neutrinos have on perturbations?
16. How does changing the curvature Ωk affect the CMB power spectrum?
17. How does changing the dark energy ΩΛ affect the CMB power spectrum?
21
An overview of cosmological perturbations
Undoubtedly the preeminent application of general relativistic perturbation theory is to cosmology. Flucta-
tions in the temperature and polarization of the Cosmic Microwave Background (CMB) provide an observa-
tional window on the Universe at 400,000 years old that, coupled with other astronomical observations, has
yielded impressively precise measurements of cosmological parameters.
The theory of cosmological perturbations is based principally on general relativistic perturbation theory
coupled to the physics of 5 species of energy-momentum: photons, baryons, non-baryonic cold dark matter,
neutrinos, and dark energy.
Dark energy was not important at the time of recombination, where the CMB that we see comes from,
but it is important today. If dark energy has a vacuum equation of state, p = −ρ, then dark energy does
not cluster (vacuum energy density is a constant), but it affects the evolution of the cosmic scale factor,
and thereby does affect the clustering of baryons and dark matter today. Moreover the evolution of the
gravitational potential along the line-of-sight to the CMB does affect the observed power spectrum of the
CMB, the so-called integrated Sachs-Wolfe effect.
Unfortunately, it is beyond the scope of these notes to treat cosmological perturbations in full. For that,
consult Scott Dodelson’s incomparable text “Modern Cosmology”.
1. Inflationary initial conditions. The theory of inflation has been remarkably successful in accounting for
many aspects of observational cosmology, even though a fundamental understanding of the inflaton scalar
field that supposedly drove inflation is missing. The current paradigm holds that primordial fluctuations
were generated by vacuum quantum fluctuations in the inflaton field at the time of inflation. The theory
makes the generic predictions that the gravitational potentials generated by vacuum fluctuations were (a)
Gaussian, (b) adiabatic (meaning that all species of mass-energy fluctuated together, as opposed to in
opposition to each other), and (c) scale-free, or rather almost scale-free (the fact that inflation came to
an end modifies slightly the scale-free character). The three predictions fit the observed power spectrum
of the CMB astonishingly well.
2. Comoving Fourier modes. The spatial homogeneity of the Friedmann-Robertson-Walker background
spacetime means that its perturbations are characterized by Fourier modes of constant comoving wavevec-
An overview of cosmological perturbations 353
tor. Each Fourier mode generated by inflation evolved independently, and its wavelength expanded with
the Universe.
3. Scalar, vector, tensor modes. Spatial isotropy on top of spatial homogenity means that the pertur-
bations comprised independently evolving scalar, vector, and tensor modes. Scalar modes dominate the
fluctuations of the CMB, and caused the clustering of matter today. Vector modes are usually assumed
to vanish, because there is no mechanism to generate the vorticity that sources vector modes, and the
expansion of the Universe tends to redshift away any vector modes that might have been present. Inflation
generates gravitational waves, which then propagate essentially freely to the present time. Gravitational
waves leave an observational imprint in the “B” (curl) mode of polarization of the CMB, whereas scalar
modes produce only an “E” (gradient) mode of polarization.
4. Power spectrum. The primary quantity measurable from observations is the power spectrum, which
is the variance of fluctuations of the CMB or of matter (as traced by galaxies, galaxy clusters, the Lyman
alpha forest, peculiar velocities, weak lensing, or 21 centimeter observations at high redshift). The statis-
tics of a Gaussian field are completely characterized by its mean and variance. The mean characterizes
the unperturbed background, while the variance characterizes the fluctuations. For a 3-dimensional sta-
tistically homogeneous and isotropic field, the variance of Fourier modes δk defines the power spectrum
P (k)
〈δkδk′〉 = 1kk′P (k) , (21.1)
where 1kk′ is the unit matrix in the Hilbert space of Fourier modes
1kk′ ≡ (2π)3δ3D(k + k′) . (21.2)
The “momentum-conserving” Dirac delta-function in equation (21.2) is a consequence of spatial translation
symmetry. Isotropy implies that the power spectrum P (k) is a function only of the absolute value k ≡ |k|of the wavevector. For a statistically rotation invariant field projected on the sky, such as the CMB, the
variance of spherical harmonic modes Θℓm ≡ δTℓm/T defines the power spectrum Cℓ
〈ΘℓmΘℓ′m′〉 = 1ℓm,ℓ′m′Cℓ (21.3)
where 1ℓm,ℓ′m′ is the unit matrix in the Hilbert space of spherical harmonics (distinguish the three usages
of δ in this paragraph: δ meaning fluctuation, δD meaning Dirac delta-function, and δ meaning Kronecker
delta, as in the following equation)
1ℓm,ℓ′m′ ≡ δℓℓ′δm,−m′ . (21.4)
Again, the “angular momentum-preserving” condition (21.4) that ℓ = ℓ′ and m+m′ = 0 is a consequence
of rotational symmetry. The same rotational symmetry implies that the power spectrum Cℓ is a function
only of the harmonic number ℓ, not of the directional harmonic number m.
5. Reheating. Early Universe inflation evidently came to an end. It is presumed that the vacuum energy
released by the decay of the inflaton field, an event called reheating, somehow efficiently produced the
matter and radiation fields that we see today. After reheating, the Universe was dominated by relativistic
354 An overview of cosmological perturbations
fields, collectively called “radiation”. Reheating changed the evolution of the cosmic scale factor from
acceleration to deceleration, but is presumed not to have generated additional fluctuations.
6. Photon-baryon fluid and the sound horizon. Photon-electron (Thomson) scattering kept photons
and baryons tightly coupled to each other, so that they behaved like a relativistic fluid. As long as the
radiation density exceeded the baryon density, which remained true up to near the time of recombination,
the speed of sound in the photon-baryon fluid was√
p/ρ ≈√
13 of the speed of light. Fluctuations with
wavelengths outside the sound horizon grew by gravity. As time went by, the sound horizon expanded
in comoving radius, and fluctuations thereby came inside the sound horizon. Once inside the sound
horizon, sound waves could propagate, which tended to decrease the gravitational potential. However,
each individual sound wave itself continued to oscillate, its oscillation amplitude δT/T relative to the
background temperature T remaining approximately constant. The relativistic suppression of the potential
at small scales is responsible for the fact that the power spectrum of matter declines at small scales.
7. Acoustic peaks in the power spectrum. The oscillations of the photon-baryon fluid produced the
characteristic pattern of peaks and troughs in the CMB power spectrum observed today. The same
peaks and troughs occur in the matter power spectrum, but are much less prominent, at a level of about
10% as opposed to the order unity oscillations observed in the CMB power spectrum. For adiabatic
fluctuations, the amplitude of the temperature fluctuations follows a pattern ∼ − cos(kηs) where ηs is the
comoving sound horizon. The n’th peak occurs at a wavenumber k where kηs ≈ nπ. In the observed
CMB power spectrum, the relevant value of the sound horizon ηs is its value ηs,∗ at recombination.
Thus the wavenumber k of the first peak of the observed CMB power spectrum occurs where kηs,∗ ≈ π.
Two competing forces cause a mode to evolve: a gravitational force that amplifies compression, and
a restoring pressure force that counteracts compression. When a mode enters the sound horizon for
the first time, the compressing gravitational force beats the restoring pressure force, so the first thing
that happens is that the mode compresses further. Consequently the first peak is a compression peak.
This sets the subsequent pattern: odd peaks are compression peaks, while even peaks are rarefaction
peaks. The observed temperature fluctuations of the CMB are produced by a combination of intrinsic
temperature fluctuations, Doppler shifts, and gravitational redshifting out of potential wells. The Doppler
shift produced by the velocity of a perturbation is 90 out of phase with the temperature fluctuation, and
so tends to fill in the troughs in the power spectrum of the temperature fluctuation. This is the main
reason that the observed CMB power spectrum remains above zero at all scales.
8. Logarithmic growth of matter fluctuations. Non-baryonic cold dark matter interacts weakly except
by gravity, and is needed to explain the observed clustering of matter in the Universe today in spite of the
small amplitude of temperature fluctuations in the CMB. The adjective “cold” refers to the requirement
that the dark matter became non-relativistic (p = 0) at some early time. If the dark matter is both
non-baryonic and cold, then it did not participate in the oscillations of the photon-baryon fluid. During
the radiation-dominated phase prior to matter-radiation equality, dark matter matter fluctuations inside
the sound horizon grow logarithmically. The logarithmic growth translates into a logarithmic increase
in the amplitude of matter fluctuations at small scales, and is a characteristic signature of non-baryonic
An overview of cosmological perturbations 355
cold dark matter. Unfortunately this signature is not readily discernible in the power spectrum of matter
today, because of nonlinear clustering.
9. Epoch of matter-radiation equality. The density of non-relativistic matter decreases more slowly than
the density of relativistic radiation. There came a point where the matter density equaled the radiation
density, an epoch called matter-radiation equality, after which the matter density exceeded the radiation
density. The observed ratio of the density of matter and radiation (CMB) today require that matter-
radiation equality occurred at a redshift of zeq ≈ 3200, a factor of 3 higher in redshift than recombination
at z∗ ≈ 1100. After matter-radiation equality, dark matter perturbations grew more rapidly, linearly
instead of just logarithmically with cosmic scale factor. A larger dark matter density causes matter-
radiation equality to occur earlier. The sound horizon at matter-radiation equality corresponds to a scale
roughly around the 2.5’th peak in the CMB power spectrum. For adiabatic fluctuations, the way that the
temperature and gravitational perturbations interact when a mode first enters the sound horizon means
that the temperature oscillation is 5 times larger for modes that enter the horizon well into the radiation-
dominated epoch versus well into the matter-dominated epoch. The effect enhances the amplitude of
observed CMB peaks higher than 2.5 relative to those lower than 2.5. The observed relative strengths
of the 3rd versus the 2nd peak of the CMB power spectrum provides a measurement of the redshift of
matter-radiation equality, and direct evidence for the presence of non-baryonic cold dark matter.
10. Sound speed. The density of baryons decreased more slowly than the density of radiation, so that at
around recombination the baryon density was becoming comparable to the radiation density. The sound
speed√
p/ρ depends on the ratio of pressure p, which was essentially entirely that of the photons, to the
density ρ, which was produced by both photons and baryons. The sound speed consequently decreased
below√
13 . Increasing the baryon-to-photon ratio at recombination has several observational effects on
the acoustic peaks of the CMB power spectrum, making it a prime measurable parameter from the CMB.
First, an increased baryon fraction increases the gravitational forcing (baryon loading), which enhances
the compression (odd) peaks while reducing the rarefaction (even) peaks. Second, increasing the baryon
fraction reduces the sound speed, which: (a) decreases the amplitude of the radiation dipole relative to the
radiation monopole, so increasing the prominence of the peaks; and (b) reduces the oscillation frequency
of the photon-baryon fluid, which shifts the peaks to larger scales. The reduced sound speed also causes
an adiabatic reduction of the amplitudes of all modes by the square root of the sound speed, but this effect
is degenerate with an overall reduction in the initial amplitudes of modes produced by inflation.
11. Recombination. As the temperature cooled below about 3,000 K, electrons combined with hydrogen
and helium nuclei into neutral atoms. This drastically reduced the amount of photon-electron scattering,
releasing the CMB to propagate almost freely. At the same time, the baryons were released from the
photons. Without radiation pressure to support them, fluctuations in the baryons began to grow like the
dark matter fluctuations.
12. Neutrinos. Probably all three species of neutrino have mass less than 0.3 eV and were therefore relativistic
up to and at the time of recombination. Each of the 3 species of neutrino had an abundance comparable
to that of photons, and therefore made an important contribution to the relativistic background and its
356 An overview of cosmological perturbations
fluctuations. Unlike photons, neutrinos streamed freely, without scattering. The relativistic free-streaming
of neutrinos provided the main source of the quadrupole pressure that produces a non-vanishing difference
Ψ − Φ between the scalar potentials. However, the neutrino quadrupole pressure was still only ∼ 10% of
the neutrino monopole pressure. To the extent that the neutrino quadrupole pressure can be approximated
as negligible, the neutrinos and their fluctuations can be treated the same as photons.
13. CMB fluctuations. The CMB fluctuations seen on the sky today represent a projection of fluctuations
on a thin but finite shell at a redshift of about 1100, corresponding to an age of the Universe of about
400,000 yr. The temperature, and the degrees of polarization in two different directions, provide 3 inde-
pendent observables at each point on the sky. The isotropy of the unperturbed radiation means that it is
most natural to measure the fluctuations in spherical harmonics, which are the eigenmodes of the rotation
operator. Similarly, it is natural to measure the CMB polarization in spin harmonics.
14. Matter fluctuations. After recombination, perturbations in the non-baryonic and baryonic matter grew
by gravity, essentially unaffected any longer by photon pressure. If one or more of the neutrino types
had a mass small enough to be relativistic but large enough to contribute appreciable density, then its
relativistic streaming could have suppressed power in matter fluctuations at small scales, but observations
show no evidence of such suppression, which places an upper limit of about an eV on the mass of the
most massive neutrino. The matter power spectrum measured from the clustering of galaxies contains
acoustic oscillations like the CMB power spectrum, but because the non-baryonic dark matter dominates
the baryons, the oscillations are much smaller.
15. Integrated Sachs-Wolfe effect. Variations in the gravitational potential along the line-of-sight to the
CMB affect the CMB power spectrum at large scales. This is called the integrated Sachs-Wolfe (ISW)
effect. If matter dominates the background, then the gravitational potential Φ has the property that it
remains constant in time for (subhorizon) linear fluctuations, and there is no ISW effect. In practice,
ISW effects are produced by at least three distinct causes. First, an early-time ISW effect is produced
by the fact that the Universe at recombination still has an appreciable component of radiation, and is
not yet wholly matter-dominated. Second, a late-time ISW effect is produced either by curvature or by a
cosmological constant. Third, a non-linear ISW effect is produced by non-linear evolution of the potential.
22
∗Cosmological perturbations in a flatFriedmann-Robertson-Walker background
For simplicity, this book considers only a flat (not closed or open) Friedmann-Robertson-Walker background.
The comoving cosmological horizon size at recombination was much smaller than today, and consequently the
cosmological density Ω was much closer to 1 at recombination than it is today. Since observations indicate
that the Universe today is within 1% of being spatially flat, it is an excellent approximation to treat the
Universe at the time of recombination as being spatially flat.
With some modifications arising from cosmological expansion, perturbation theory on a flat FRW back-
ground is quite similar to perturbation theory in flat (Minkowski) space, Chapter 20.
The strategy is to start in a completely general gauge, and to discover how the conformal Newtonian
gauge, which is used in subsequent Chapters, emerges naturally as that gauge in which the perturbations
are precisely the physical perturbations.
22.1 Unperturbed line-element
It is convenient to choose the coordinate system xµ = η, xi to consist of conformal time η together with
Cartesian comoving coordinates x ≡ xi ≡ x, y, z. The coordinate metric of the unperturbed background
FRW geometry is then
ds2 = a(η)2(
− dη2 + dx2 + dy2 + dz2)
, (22.1)
where a(η) is the cosmic scale factor. The unperturbed coordinate metric is thus the conformal Minkowski
metric0
gµν = a(η)2ηµν . (22.2)
The tetrad is taken to be orthonormal, with the unperturbed tetrad axes γγm ≡ γγ0.γγ1,γγ2,γγ3 being aligned
with the unperturbed coordinate axes gµ ≡ gη, gx, gy, gz so that the unperturbed vierbein and inverse
vierbein are respectively 1/a and a times the unit matrix
0
emµ =
1
aδµm ,
0
em
µ = a δmµ . (22.3)
358 ∗Cosmological perturbations in a flat Friedmann-Robertson-Walker background
Let overdot denote partial differentiation with respect to conformal time η,
overdot ≡ ∂
∂η, (22.4)
so that for example a ≡ da/dη. The coordinate time derivative ∂/∂η is to be distinguished from the directed
time derivative ∂0 ≡ e0µ ∂/∂xµ. Let ∇i denote the comoving gradient
∇i ≡∂
∂xi, (22.5)
which should be distinguished from the directed derivative ∂i ≡ eiµ ∂/∂xµ.
22.2 Comoving Fourier modes
Since the unperturbed Friedmann-Robertson-Walker spacetime is spatially homogeneous and isotropic, it
is natural to work in comoving Fourier modes. Comoving Fourier modes have the key property that they
evolve independently of each other, as long as perturbations remain linear. Equations in Fourier space are
obtained by replacing the comoving spatial gradient ∇i by −i times the comoving wavevector ki (the choice
of sign is the standard convention in cosmology)
∇i → −iki . (22.6)
By this means, the spatial derivatives become algebraic, so that the partial differential equations governing
the evolution of perturbations become ordinary differential equations.
In what follows, the comoving gradient ∇i will be used interchangeably with −iki, whichever is most
convenient.
22.3 Classification of vierbein perturbations
The tetrad-frame components ϕmn of the vierbein perturbation of the FRW geometry decompose in much
the same way as in flat Minkowski case into 6 scalars, 4 vectors (8 degrees of freedom), and 1 tensor (2
degrees of freedom) (the following equations are essentially the same as those (20.6) for the flat Minkowski
background),
ϕ00 = ψscalar
, (22.7a)
ϕ0i = ∇iwscalar
+ wivector
, (22.7b)
ϕi0 = ∇iwscalar
+ wivector
, (22.7c)
ϕij = δij φscalar
+∇i∇jhscalar
+ εijk∇khscalar
+∇ihjvector
+∇j hivector
+ hijtensor
. (22.7d)
22.3 Classification of vierbein perturbations 359
The tetrad-frame components ǫm of the coordinate shift of the coordinate gauge transformation (18.8)
similarly decompose into 2 scalars and 1 vector (the following equation is essentially the same as that (20.8)
for the flat Minkowski background)
ǫm = ǫ0scalar
, ∇iǫscalar
+ ǫivector
. (22.8)
The vierbein perturbations ϕmn transform under a coordinate gauge transformation (18.8) as, equation (18.25),
ϕ00 → ψ +1
a
∂ǫ0∂η
scalar
, (22.9a)
ϕ0i → ∇i
(
w +1
a
( ∂
∂η− a
a
)
ǫ
)
scalar
+
(
wi +1
a
( ∂
∂η− a
a
)
ǫi
)
vector
, (22.9b)
ϕi0 → ∇i
(
w +1
aǫ0
)
scalar
+ wivector
, (22.9c)
ϕij → δij
(
φ− a
a2ǫ0
)
scalar
+∇i∇j
(
h+1
aǫ
)
scalar
+ εijk∇khscalar
+∇i
(
hj +1
aǫj
)
vector
+∇j hivector
+ hijtensor
, (22.9d)
or equivalently
ψ → ψ +1
a
∂ǫ0∂η
, (22.10a)
w → w +1
a
( ∂
∂η− a
a
)
ǫ , wi → wi +1
a
( ∂
∂η− a
a
)
ǫi , (22.10b)
w → w +1
aǫ0 , wi → wi , (22.10c)
φ→ φ− a
a2ǫ0 , h→ h+
1
aǫ , h→ h , hi → hi +
1
aǫi , hi → hi , hij → hij . (22.10d)
Eliminating the coordinate shift ǫm from the transformations (22.10) yields 12 coordinate gauge-invariant
combinations of the perturbations
ψ −( ∂
∂η+a
a
)
w , w − h , wi − hi , wi , φ+a
aw , h , hi , hij . (22.11)
Six combinations of these coordinate gauge-invariant perturbations depend only on the symmetric part
ϕmn + ϕnm of the vierbein perturbations, and are therefore tetrad gauge-invariant as well as coordinate
gauge-invariant. These 6 coordinate and tetrad gauge-invariant perturbations comprise 2 scalars, 1 vector,
360 ∗Cosmological perturbations in a flat Friedmann-Robertson-Walker background
and 1 tensor
Ψscalar
≡ ψ −( ∂
∂η+a
a
)
(w + w − h) , (22.12a)
Φscalar
≡ φ+a
a(w + w − h) , (22.12b)
Wivector
≡ wi + wi − hi − ˙hi , (22.12c)
hijtensor
. (22.12d)
22.4 Metric, tetrad connections, and Einstein tensor
This section gives expressions in a completely general gauge for perturbed quantities in the flat Friedmann-
Robertson-Walker background geometry.
The perturbed coordinate metric gµν is
gηη = −a2(1 + 2ψ) , (22.13a)
gηi = −a2[
∇i(w + w) + (wi + wi)]
, (22.13b)
gij = a2[
(1− 2φ)δij − 2∇i∇jh−∇i(hj + hj)−∇j(hi + hi)− 2 hij
]
. (22.13c)
The coordinate metric is tetrad gauge-invariant, but not coordinate gauge-invariant.
The perturbed tetrad connections Γkmn are
Γ0i0 =1
a
[
−∇i
(
ψ −( ∂
∂η+a
a
)
w
)
+( ∂
∂η+a
a
)
wi
]
, (22.14a)
Γ0ij =1
a
[
(
− aa
+ F)
δij −∇i∇j(w − h)−1
2(∇iWj +∇jWi) +∇jwi + hij
]
, (22.14b)
Γij0 =1
a
[
1
2(∇iWj −∇jWi)−
∂
∂η(εijl∇lh−∇ihj +∇j hi)
]
, (22.14c)
Γijk =1
a
[
(δjk∇i − δik∇j)(
φ+a
aw)
− a
a(δikδjl − δjkδil)wl
−∇k(εijl∇lh−∇ihj +∇j hi) +∇ihjk −∇jhik
]
, (22.14d)
where F is defined by
F ≡ a
aψ + φ . (22.15)
Being purely tetrad-frame quantities, the tetrad connections are automatically coordinate gauge-invariant,
22.4 Metric, tetrad connections, and Einstein tensor 361
but they are not tetrad gauge-invariant. The quantity F defined by equation (22.15) is not coordinate gauge-
invariant, but the combination − a/a2 + F/a that appears in the expression (22.14b) for Γ0ij is coordinate
and tetrad gauge-invariant.
Exercise 22.1 Coordinate gauge invariance of − a/a2 + F/a. Argue that under a coordinate gauge
transformation of the conformal time, η → η′ = η + ǫη, the cosmic scale factor a(η) and its derivative
a ≡ da/dη transforms as (see §18.8)
a→ a+ Lǫa = a− aǫη = a+a
aǫ0 , a→ a+ Lǫa = a− aǫη = a+
a
aǫ0 . (22.16)
Check that this behaviour is consistent with the gauge transformation (18.28) of gηη, equation (22.13a).
Hence show that − a/a2 + F/a is coordinate gauge invariant. THIS IS NOT WORKING.
The perturbed tetrad-frame Einstein tensor Gmn is
G00 =1
a2
3a2
a2− 6
a
aF + 2∇2Φ
scalar
, (22.17a)
G0i =1
a2
2∇i
(
F +( a
a− 2
a2
a2
)
w
)
scalar
+1
2∇2Wi + 2
( a
a− 2
a2
a2
)
wi
vector
, (22.17b)
Gij =1
a2
(
− 2a
a+a2
a2+ 2
( ∂
∂η+ 2
a
a
)
F + 2( a
a− 2
a2
a2
)
ψ
)
δij
scalar
− (∇i∇j − δij∇2)(Ψ − Φ)scalar
+1
2
( ∂
∂η+ 2
a
a
)
(∇iWj +∇jWi)
vector
−( ∂2
∂η2+ 2
a
a
∂
∂η−∇2
)
hij
tensor
. (22.17c)
Being tetrad-frame quantities, all components of the tetrad-frame Einstein tensor are automatically coordi-
nate gauge-invariant. The time-time G00 and space-space Gij components are not only coordinate but also
tetrad gauge-invariant, as follows from the fact that these components depend only on symmetric combi-
nations of the vierbein potentials. Specifically, the quantities 3(a2/a4) − 6(a/a3)F on the right hand side
of equation (22.17a) for G00, and the coefficient of δij (including the overall 1/a2 factor) on the right hand
side of equation (22.17c) for Gij , are coordinate and tetrad gauge-invariant. However, the time-space com-
ponents G0i are not tetrad gauge-invariant, as is evident from the fact that equation (22.17b) involves the
non-tetrad-gauge-invariant perturbations w and wi. Physically, under a tetrad boost by a velocity v of linear
order, the time-space components G0i change by first order v, but G00 and Gij change only to second order
v2. Thus to linear order, only G0i changes under a tetrad boost. Note that G0i changes under a tetrad boost
(w and wi), but not under a tetrad rotation (h and hi).
362 ∗Cosmological perturbations in a flat Friedmann-Robertson-Walker background
22.5 ADM gauge choices
The ADM (3+1) formalism, Chapter 13, chooses the tetrad time axis γγ0 to be orthogonal to hypersurfaces
of constant time, η = constant, equivalent to requiring that the tetrad time axis be orthogonal to each of
the spatial coordinate axes, γγ0 · gi = 0, equation (13.2). The ADM choice is equivalent to setting
w = wi = 0 . (22.18)
The ADM choice simplifies the tetrad-frame connections (22.14) and the time-space component G0i of the
tetrad-frame frame Einstein tensor, equation (22.17b).
Another gauge choice that significantly simplifies the tetrad connections (22.14), though does not affect
the Einstein tensor (22.17), is
h = hi = 0 . (22.19)
If the wavevector k is taken along the coordinate z-direction, then the gauge choice hi = 0 is equivalent to
choosing the tetrad 3-axis (z-axis) γγ3 to be orthogonal to the coordinate x and y-axes, γγ3 · gx = γγ3 · gy = 0.
The gauge choice h = 0 is equivalent to rotating the tetrad axes about the 3-axis (z-axis) so that γγ1 · gy =
γγ2 · gx.
22.6 Conformal Newtonian gauge
Conformal Newtonian gauge sets
w = w = wi = h = h = hi = hi = 0 , (22.20)
so that the retained perturbations are the 6 coordinate and tetrad gauge-invariant perturbations (22.12)
Ψscalar
= ψ , (22.21a)
Φscalar
= φ , (22.21b)
Wivector
= wi , (22.21c)
hijtensor
. (22.21d)
In conformal Newtonian gauge, the scalar perturbations of the Einstein equations are the energy density,
energy flux, monopole pressure, and quadrupole pressure equations
− 3a
aF − k2Φ = 4πGa2
1
T 00 , (22.22a)
ikF = 4πGa2 ki
1
T 0i , (22.22b)
F + 2a
aF +
( a
a− 2
a2
a2
)
Ψ− k2
3(Ψ − Φ) =
4
3Gπa2 δij
1
T ij , (22.22c)
k2(Ψ − Φ) = 8πGa2(
32 kikj − 1
2 δij
)
1
T ij , (22.22d)
22.7 Synchronous gauge 363
where F is the coordinate and tetrad gauge-invariant quantity
F ≡ a
aΨ + Φ . (22.23)
All 4 of the scalar Einstein equations (22.22) are expressed in terms of gauge-invariant variables, and are
therefore fully gauge-invariant.
The energy and momentum equations (22.22a) and (22.22b) can be combined to eliminate F , yielding an
equation for Φ alone
−k2Φ = 4πa2
(
1
T 00 + 3a
a
ki
ik
1
T 0i
)
. (22.24)
The quantity in parentheses on the right hand side of equation (22.24) is the source for the scalar potential
Φ, and can be interpreted as a measure of the “true” energy fluctuation. This is however just a matter of
interpretation: the individual perturbations1
T 00 and1
T 01 are both individually gauge-invariant, and therefore
have physical meaning.
If the energy-momentum tensors of the various matter components are arranged so as to conserve overall
energy-momentum, as they should, then 2 of the 4 equations (22.22a)–(22.22d) are redundant, since they
serve simply to enforce conservation of energy and scalar momentum. Thus it suffices to take, as the equations
governing the potentials Φ and Ψ, any 2 of the equations (22.22a)–(22.22d). One is free to retain whichever 2
of the equations is convenient. Usually the 1st equation, the energy equation (22.22a), and the 4th equation,
the quadrupole pressure equation (22.22d), are most convenient to retain. But sometimes the 2nd equation,
the scalar momentum equation (22.22b), is more convenient in place of the energy equation (22.22a).
22.7 Synchronous gauge
One gauge that remains in common use in cosmology, but is not used here, is synchronous gauge, discussed
in the case of Minkowski background space in §20.10. The cosmological synchronous gauge choices are the
same as for the Minkowski background, equations (20.63) and (20.64):
ψ = w = w = wi = wi = h = hi = 0 . (22.25)
The gauge-invariant perturbations (22.12) in synchronous gauge are
Ψscalar
= h+a
ah , (22.26a)
Φscalar
= φ− a
ah , (22.26b)
Wivector
= − hi , (22.26c)
hijtensor
. (22.26d)
23
Cosmological perturbations: a simplest set ofassumptions
1. Consider only scalar modes.
2. Consider explicitly only two species: non-baryonic cold dark matter, and radiation consisting of photons
and neutrinos lumped together. Neglect the contribution of baryons to the mass density.
3. Treat the radiation as almost isotropic, so it is dominated by its first two moments, the monopole and
dipole. In practice, electron-photon scattering keeps photons almost isotropic. Unlike photons, neutrinos
stream freely, but they inherit an approximately isotropic distribution from an early time when they were
in thermodynamic equilibrium.
4. Include damping from photon-electron (Thomson) scattering by allowing the radiation a small quadrupole
moment, the diffusion approximation.
23.1 Perturbed FRW line-element
Perturbed FRW line-element in conformal Newtonian gauge
ds2 = a2[
−(1 + 2Ψ)dη2 + δij(1− 2Φ)dxidxj]
, (23.1)
where a(η) is the cosmic scale factor, a function only of conformal time η.
23.2 Energy-momenta of ideal fluids
In the simplest approximation, matter, radiation, and dark energy can each be treated as ideal fluids. The
energy-momentum tensor of an ideal fluid with proper density ρ and isotropic pressure p in its own rest
frame, and moving with bulk 4-velocity um relative to the conformal Newtonian tetrad frame, is
Tmn = (ρ+ p)umun + p ηmn . (23.2)
In the situation under consideration, the fluids satisfy equations of state
p = wρ (23.3)
23.2 Energy-momenta of ideal fluids 365
with w constant. Specifically, w = 0 for non-relativistic matter, w = 1/3 for relativistic radiation, and
w = −1 for dark energy with constant density. Furthermore, all the fluids are moving with non-relativistic
bulk velocities, including the radiation, which is almost isotropic, and therefore has a small bulk velocity
even though the individual particles of radiation move at the speed of light. The bulk tetrad-frame 4-velocity
um is thus, to linear order
um = 1, vi , (23.4)
where vi is the non-relativistic spatial bulk 3-velocity (the spatial tetrad metric is Euclidean, so vi = vi).
The density ρ can be written in terms of the unperturbed density ρ and a fluctuation δ defined by
ρ = ρ[1 + (1 + w)δ] . (23.5)
The factor 1+w is included in the definition of the fluctuation δ because it simplifies the resulting perturbation
equations (23.12). As you will discover in Exercise 23.1, the fluctuation δ can be interpreted physically as
the entropy fluctuation,
δ =1
1 + w
δρ
ρ=δs
s. (23.6)
For matter, w = 0, the entropy fluctuation coincides with the density fluctuation, δ = δρ/ρ. For dark
energy, w = −1, the density fluctuation is necessarily zero, δρ/ρ = 0. To linear order in the velocity vi, the
tetrad-frame energy-momentum tensor (23.2) of the ideal fluid is then
T 00 ≡ ρ[1 + (1 + w)δ] , (23.7a)
T 0i ≡ (1 + w)ρ vi , (23.7b)
T ij = wρ[1 + (1 + w)δ] δij . (23.7c)
To linear order in the fluctuation δ, velocity vi, and potentials Ψ and Φ, and in conformal Newtonian gauge,
conservation of energy and momentum requires
DmTm0 =
1−Ψ
a
[
∂ρ
∂η+ (1 + w)∇i(ρvi) + 3(1 + w)ρ
( a
a− Φ
)
]
= 0 , (23.8a)
DmTmi =
1
a
[
(1 + w)∂ρvi
∂η+ 4(1 + w)
a
aρvi + w∇iρ+ (1 + w)ρ∇iΨ
]
= 0 . (23.8b)
The energy conservation equation (23.8a) has an unperturbed part,
Dm
0
Tm0 =1
a
[
∂ρ
∂η+ 3(1 + w)ρ
a
a
]
= 0 , (23.9)
which implies the usual result that the mean density evolves as a power law with cosmic scale factor,
ρ ∝ a−3(1+w). Subtracting appropriate amounts of the unperturbed energy conservation equation (23.9)
from the perturbed energy-momentum conservation equations (23.8) yields equations for the fluctuation δ
366 Cosmological perturbations: a simplest set of assumptions
and velocity vi:
δ +∇ivi = 3Φ , (23.10a)
vi + (1− 3w)a
avi + w∇iδ = −∇iΨ . (23.10b)
Now decompose the 3-velocity vi into its scalar v and vector v⊥,i parts. Up to this point, the scalar part of
a vector has been taken to be the gradient of a potential. But here it is advantageous to absorb a factor of k
into the definition of the scalar part v of the velocity, so that instead of vi = −ikiv + v⊥,i in Fourier space,
the velocity is given in Fourier space by
vi = −ikiv + v⊥,i . (23.11)
The advantage of this choice is that v is dimensionless, as are δ and Ψ and Φ. The scalar parts of the
perturbation equations (23.10) are then
δ − kv = 3Φ , (23.12a)
v + (1− 3w)a
av + wkδ = −kΨ . (23.12b)
Combining the two equations (23.12) for the fluctuation δ and velocity v yields a second-order differential
equation for δ − 3Φ,[
d2
dη2+ (1− 3w)
a
a
d
dη+ k2w
]
(δ − 3Φ) = −k2(Ψ + 3wΦ) . (23.13)
For positive w, equation (23.13) is a wave equation for a damped, forced oscillator with sound speed√w. The
resulting generic behaviour for the particular cases of matter (w = 0) and radiation (w = 13 ) is considered
in §23.6 and §23.7 below.
A more careful treatment, deferred to Chapter 24, accounts for the complete momentum distribution of
radiation by expanding the temperature perturbation Θ ≡ δT/T in multipole moments, equation (24.48).
The radiation fluctuation δr and scalar bulk velocity vr are related to the first two multipole moments of
the temperature perturbation, the monopole Θ0 and the dipole Θ1, by
δr = 3Θ0 , (23.14a)
vr = 3Θ1 . (23.14b)
The factor of 3 arises because the unperturbed radiation distribution is in thermodynamic equilibrium, for
which the entropy density is s ∝ T 3, so δr = 3δT/T .
The energy-momentum perturbation1
Tmn that goes into the Einstein equations (22.22) are, from equa-
tions (23.7) with the unperturbed part subtracted,
1
T 00 ≡ (1 + w)ρ δ , (23.15a)1
T 0i ≡ (1 + w)ρ vi , (23.15b)1
T ij = w(1 + w)ρ δ δij . (23.15c)
23.3 Diffusive damping 367
Exercise 23.1 Entropy fluctuation. The purpose of this exercise is to discover that the fluctuation
δ defined by equation (23.5) can be interpreted as the entropy fluctuation. According to the first law of
thermodynamics, the entropy density s of a fluid of energy density ρ, pressure p, and temperature T in a
volume V satisfies
d(ρV ) + pdV = Td(sV ) . (23.16)
If the fluid is ideal, so that ρ, p, T , and s are independent of volume V , then integrating the first law (23.16)
implies that
ρV + pV = TsV . (23.17)
This implies that the entropy density s is related to the other variables by
s =ρ+ p
T. (23.18)
Show that, for ideal fluid with equation of state p/ρ = w = constant, the first law (23.16) together with the
expression (23.18) for entropy implies that
T ∝ ρw/(1+w) , (23.19)
and hence
s ∝ ρ1/(1+w) . (23.20)
Conclude that small variations of the entropy and density are related by
δs
s=
1
1 + w
δρ
ρ, (23.21)
confirming equation (23.6). [Hint: Do not confuse what is being asked here with adiabatic expansion. The
results (23.19) are properties of the fluid, independent of whether the fluid is changing adiabatically. For
adiabatic expansion, the fluid satisfies the additional condition sV = constant.]
23.3 Diffusive damping
The treatment of matter and radiation as ideal fluids misses a feature that has a major impact on observed
fluctuations in the CMB, namely the diffusive damping of sound waves that results from the finite mean
free path to electron-photon scattering. As recombination approaches, the scattering mean free path length-
ens, until at recombination photons are able to travel freely across the Universe, ready to be observed by
astronomers. The damping is greater at smaller scales, and is responsible for the systematic decrease in the
CMB power spectrum to smaller scales.
As expounded in Chapter 24, in §24.13 and following, the damping can be taken into account to lowest
368 Cosmological perturbations: a simplest set of assumptions
order, the diffusion approximation, by admitting a small quadrupole moment Θ2 to the photon distribu-
tion. A detailed analysis of the collisional Boltzmann equation for photons, §24.12, reveals that the photon
quadrupole Θ2 is given in the diffusion approximation by equation (24.94). There is an additional source of
damping that arises from viscous baryon drag, §24.15, but this effect vanishes in the limit of small baryon
density, and is neglected in the present Chapter.
The diffusive damping resulting from a small photon quadrupole conserves the energy and momentum
of the photon fluid, so that covariant momentum conservation DmTmn = 0 continues to hold true within
the photon fluid. By contrast, viscous baryon drag, §24.15, neglected in this Chapter, transfers momentum
between photons and baryons.
Define the dimensionless quadrupole q by
T ijquadrupole = (1 + w)ρq
(
32 kikj − 1
2 δij
)
, (1 + w)ρq ≡(
kikj − 13 δij
)
T ij . (23.22)
For photons, the dimensionless quadrupole q is related to the photon quadrupole harmonic Θ2 by, equa-
tion (24.55d),
q = − 2Θ2 . (23.23)
In the presence of a quadrupole pressure, the energy conservation equation (23.8a) is unchanged, but the
momentum conservation equation (23.8b) is modified by the change Ψ→ Ψ + q:
DmTmi =
1
a
[
(1 + w)∂ρvi
∂η+ 4(1 + w)
a
aρvi + w∇iρ+ (1 + w)ρ∇i(Ψ + q)
]
= 0 . (23.24)
The consequent equations (23.10b), (23.12b), (23.13) are similarly modified by Ψ → Ψ + q. In particular,
the velocity equation (23.12b) is modified to
v + (1− 3w)a
av + wkδ = −k(Ψ + q) . (23.25)
23.4 Equations for the simplest set of assumptions
Non-baryonic cold dark matter, subscripted c:
δc − k vc = 3 Φ , (23.26a)
vc +a
avc = −kΨ . (23.26b)
Radiation, which includes both photons and neutrinos:
Θ0 − kΘ1 = Φ , (23.27a)
Θ1 +k
3Θ0 = − k
3(Ψ− 2Θ2) , (23.27b)
Θ2 = − 4k
9neσTaΘ1 . (23.27c)
23.4 Equations for the simplest set of assumptions 369
Einstein energy and quadrupole pressure equations:
− k2Φ− 3a
aF = 4πGa2(ρcδc + 4ρrΘ0) , (23.28a)
k2(Ψ− Φ) = − 32πGa2ρrΘ2 , (23.28b)
where
F ≡ a
aΨ + Φ . (23.29)
In place of the Einstein energy equation (23.28a) it is sometimes convenient to use the Einstein momentum
equation
−kF = 4πGa2(ρcvc + 4ρrΘ1) , (23.30)
which, because the matter and radiation equations (23.26) and (23.27) already satisfy covariant energy-
momentum conservation, is not an independent equation.
The radiation quadrupole Θ2, equation (23.27c), derived in Chapter 24, equation (24.94), is proportional
to the comoving mean free path lT to electron-photon (Thomson) scattering,
lT ≡1
neσTa, (23.31)
where σT is the Thomson cross-section. The quadrupole becomes important only near recombination, where
the increasing mean free path to electron-photon scattering leads to dissipation of photon-baryon sound
waves. In the simple treatment of this Chapter, neutrinos are being lumped with photons, and of course
neutrinos do not scatter, but rather stream freely. However, radiation is gravitationally sub-dominant near
recombination, so not much error arises from treating neutrinos as gravitationally the same as photons near
recombination. To make comparison with observed CMB fluctuations, the important thing is to follow
the evolution of the photon multipoles, and for this purpose the radiation quadrupole defined by equa-
tion (23.27c), without any correction for neutrino-to-photon ratio, is the appropriate choice.
In much of the remainder of this Chapter, that is, excepting in §23.7 and §23.14, the radiation quadrupole
will be set to zero,
Θ2 = 0 , (23.32)
which is equivalent to neglecting the effect of damping. If the radiation quadrupole vanishes, then the
Einstein quadrupole pressure equation (23.28b) implies that the scalar potentials Ψ and Φ are equal,
Ψ = Φ . (23.33)
In any case, the radiation quadrupole is always small, so that Ψ ≈ Φ to a good approximation.
370 Cosmological perturbations: a simplest set of assumptions
23.5 Unperturbed background
In the unperturbed background, the unperturbed dark matter density ρc and radiation density ρr evolve
with cosmic scale factor as
ρc ∝ a−3 , ρr ∝ a−4 . (23.34)
The Hubble parameter H is defined in the usual way to be
H ≡ 1
a
da
dt=
a
a2, (23.35)
in which overdot represents differentiation with respect to conformal time, a ≡ da/dη. The Friedmann
equations for the background imply that the Hubble parameter for a universe dominated by dark matter
and radiation is
H2 =8πG
3(ρc + ρr) =
H2eq
2
(
a3eq
a3+a4eq
a4
)
(23.36)
where aeq and Heq are the cosmic scale factor and the Hubble parameter at the time of matter-radiation
equality, ρc = ρr.
The comoving horizon distance η is defined to be the comoving distance that light travels starting from
zero expansion. This is
η =
∫ a
0
da
a2H=
2√
2
aeqHeq
(√
1 +a
aeq− 1
)
=2√
2
aeqHeq
(
a/aeq
1 +√
1 + a/aeq
)
. (23.37)
In the radiation- and matter-dominated epochs respectively, the comoving horizon distance η is
η =
√2
aeqHeq
(
a
aeq
)
∝ a radiation-dominated ,
2√
2
aeqHeq
(
a
aeq
)1/2
∝ a1/2 matter-dominated .
(23.38)
The ratio of the comoving horizon distance η to the comoving cosmological horizon distance 1/(aH) is
ηaH =2√
1 + a/aeq
1 +√
1 + a/aeq
, (23.39)
which is evidently a number of order unity, varying between 1 in the radiation-dominated epoch a ≪ aeq,
and 2 in the matter-dominated epoch a≫ aeq.
Exercise 23.2 Matter-radiation equality.
1. Argue that the redshift zeq of matter-radiation equality is given by
1 + zeq =a0
aeq= ? Ωmh
2 , (23.40)
23.6 Generic behaviour of non-baryonic cold dark matter 371
where Ωm is the matter density today relative to critical. What is the factor, and what is its numerical
value? The factor depends on the energy-weighted effective number of relativistic species gρ, Exercise 10.18.
Should this gρ be that now, or that at matter-radiation equality?
2. Show that the ratio Heq/H0 of the Hubble parameter at matter-radiation equality to that today is
Heq
H0=√
2Ωm(1 + zeq)3/2 . (23.41)
Solution. The redshift zeq of matter-radiation equality is given by
1 + zeq =Ωm
Ωr=
45c5~3ΩmH20
4π3Ggρ(kT0)4= 8.093× 104 Ωmh
2
gρ= 3200
( gρ
3.36
)−1(
Ωmh2
0.133
)
, (23.42)
where T0 = 2.725 K is the present-day CMB temperature, and gρ = 2 + 6 78
(
411
)4/3= 3.36 is the energy-
weighted effective number of relativistic species at matter-radiation equality.
23.6 Generic behaviour of non-baryonic cold dark matter
Combining equations (23.26) for the dark matter overdensity and velocity gives
(
d2
dη2+a
a
d
dη
)
(δc − 3Φ) = −k2Ψ = −k2Φ , (23.43)
where the last expression follows because Ψ = Φ to a good approximation. In the absence of a driving
potential, Φ = 0, the dark matter velocity would redshift as vc ∝ 1/a, and the dark matter density equa-
tion (23.26a) would then imply that δc = kvc ∝ a−1. In the radiation-dominated epoch, where η ∝ a,
this leads to a logarithmic growth in the overdensity δc, even though there is no driving potential, and the
velocity is redshifting to a halt. In the matter-dominated epoch, where η ∝ a1/2, the dark matter overdensity
δc would freeze out at a constant value, in the absence of a driving potential.
Exercise 23.3 Generic behaviour of dark matter. Find the homogeneous solutions of equation (23.43).
Hence find the retarded Green’s function of the equation. Write down the general solution of equation (23.43)
as an integral over the Green’s function.
Solution. The general solution of equation (23.43) is
δc(a)−3Φ(a) = A0+A1 ln
(√1 + a+ 1√1 + a− 1
)
+2k2
∫ a
0
ln
[
(√
1 + a+ 1)
(√
1 + a− 1)
(√
1 + a′ − 1)
(√
1 + a′ − 1)
]
Φ(a′)a′ da′√1 + a′
, (23.44)
where A0 and A1 are constants.
372 Cosmological perturbations: a simplest set of assumptions
23.7 Generic behaviour of radiation
Combining equations (23.27) for the radiation monopole, dipole, and quadrupole gives(
3d2
dη2+ 2√
3κd
dη+ k2
)
(Θ0 − Φ) = − k2 (Ψ + Φ) = − 2 k2Φ , (23.45)
where in the last expression the approximation Ψ = Ψ has again been invoked. The coefficient κ of the linear
derivative term in equation (23.45) is a damping coefficient
κ ≡ 4k2lT
9√
3, (23.46)
where lT is the comoving electron-photon scattering (Thomson) mean free path, equation (23.31), The mean
free path is small, Exercise 23.6, except near recombination. In the absence of a driving potential, Φ = 0,
and in the absence of damping, κ = 0, the radiation oscillates as Θ0 ∝ e±iωη with frequency ω =√
13 k. In
other words, the solutions are sound waves, moving at the sound speed
cs =ω
k=
√
1
3. (23.47)
Define the conformal sound time by
ηs ≡ csη =η√3. (23.48)
In terms of the conformal sound time ηs, the differential equation (23.45) becomes(
d2
dη2s
+ 2κd
dηs+ k2
)
(Θ0 − Φ) = − 2k2Φ . (23.49)
As you will discover in Exercises 23.7 and 23.5, equation (23.49) describes damped oscillations forced by the
potential on the right hand side.
Exercise 23.4 Generic behaviour of radiation. Find the homogeneous solutions of equation (23.49)
in the case of zero damping, κ = 0. Hence find the retarded Green’s function of the equation. Write down
the general solution of equation (23.49) as an integral over the Green’s function. Convince yourself that
Θ0 − Φ oscillates about −2Φ.
Solution. The general solution of equation (23.49) is, with α ≡ kηs,
Θ0(α)− Φ(α) = B0 cosα+B1 sinα− 2
∫ α
0
sin(α− α′)Φ(α′) dα′ , (23.50)
where B0 and B1 are constants.
Exercise 23.5 Behaviour of radiation in the presence of damping. Now suppose that there is a
small damping coefficient, κ≪ k. Try a solution of the form Θ0 −Φ ∝ eR
ω dηs in equation (23.49). Suppose
that the frequency ω changes slowly over a period, ω′ ≪ ω2, so that ω′ can be set to zero. Show that
23.8 Regimes 373
the homogeneous solutions of equation (23.49) are approximately Θ0 − Φ ∝ e−R
κ dηs ± ikηs . Hence find the
retarded Green’s function, and write down the general solution to equation (23.49).
Solution. See §24.17. The general solution of equation (23.49) is, with α ≡ kηs and β ≡∫
κ dηs,
Θ0(α) − Φ(α) = e−β (B0 cosα+B1 sinα) − 2
∫ α
0
e−(β−β′) sin(α− α′)Φ(α′) dα′ , (23.51)
where B0 and B1 are constants.
23.8 Regimes
In the remainder of this Chapter, approximate analytic solutions are developed that describe the evolution
of perturbations in the matter and radiation in various regimes. The regimes are:
1. Superhorizon scales, §23.9.
2. Radiation-dominated:
a. adiabatic initial conditions, §23.10;
b. isocurvature initial conditions, §23.11.
3. Subhorizon scales, §23.12.
4. Matter-dominated, §23.13.
5. Recombination §23.14.
6. Post-recombination §23.15.
7. Matter with dark energy §23.16.
8. Matter with dark energy and curvature §23.17.
23.9 Superhorizon scales
At sufficiently early times, any mode is outside the horizon, kη < 1. In the superhorizon limit kη ≪ 1, the
evolution equations (23.27)–(23.28) reduce to
δc = 3Φ , (23.52a)
Θ0 = Φ , (23.52b)
− 3a
aF = 4πGa2(ρcδc + 4ρrΘ0) . (23.52c)
The first two of these equations evidently imply that the dark matter overdensity δc and radiation monopole
Θ0 are related to the potential Φ by
δc = 3Φ + constant , (23.53a)
Θ0 = Φ + constant . (23.53b)
374 Cosmological perturbations: a simplest set of assumptions
aeq arec
Cosmic scale factor →
Com
ovin
gdi
stan
ce→
superhorizon
radiation
radi
atio
nba
ckgr
ound
mat
ter
fluc
tuat
ion
matter
horizon
Figure 23.1 Various regimes in the evolution of fluctuations. The line increasing diagonally from bottomleft to top right is the comoving horizon distance η. Above this line are superhorizon fluctuations, whosecomoving wavelengths exceed the horizon distance, while below the line are subhorizon fluctuations, whosecomoving wavelengths are less than the horizon distance. The vertical line at cosmic scale factor aeq ≈a0/3200 marks the moment of matter-radiation equality. Before matter-radiation equality (to the left), thebackground mass-energy is dominated by radiation, while after matter-radiation equality (to the right), thebackground mass-energy is dominated by matter. Once a fluctuation enters the horizon, the non-baryonicmatter fluctuation tends to grow, whereas the radiation fluctuation tend to decay, so there is an epochprior to matter-radiation equality where gravitational perturbations are dominated by matter rather thanradiation fluctuations, even though radiation dominates the background energy density. The vertical lineat a∗ ≈ a0/1100 marks recombination, where the temperature has cooled to the point that baryons changefrom being mostly ionized to mostly neutral, and the Universe changes from being opaque to transparent.The observed CMB comes from the time of recombination.
In effect, the dark matter velocity vc and radiation dipole Θ1 are negligibly small at superhorizon scales,
vc = Θ1 = 0 . (23.54)
Plugging the solutions (23.53) into the Einstein energy equation (23.52c), and replacing derivatives with
respect to conformal time η with derivatives with respect to cosmic scale factor a,
∂
∂η= a
∂
∂a= a2H
∂
∂a, (23.55)
23.9 Superhorizon scales 375
10−3 10−2 10−1 1 10 102 103
.0
.5
1.0
Cosmic scale factor a / aeq
Φ/Φ
(lat
e)
isocurvature
adiabatic
Figure 23.2 Evolution of the scalar potential Φ at superhorizon scales, from radiation-dominated to matter-dominated. The scale for the potential is normalized to its value Φ(late) at late times a≫ aeq.
with the Hubble parameter H from equation (23.36) gives the first order differential equation, in units
aeq = Heq = 1,
2a(1 + a)Φ′ + (6 + 5a)Φ + 4C0 + C1a = 0 , (23.56)
where prime ′ denotes differentiation with respect to cosmic scale factor, d/da, and the constants C0 and C1
are
C0 = Θ0(0)− Φ(0) , C1 = δc(0)− 3Φ(0) . (23.57)
The constants C0 and C1 are set by initial conditions. There are adiabatic and isocurvature initial conditions.
Inflation generically produces adiabatic fluctuations, in which matter and radiation fluctuate together
δc(0) = 3 Θ0(0) = − 32 Φ(0) adiabiatic . (23.58)
Notice that a positive energy fluctuation corresponds to a negative potential, consistent with Newtonian
intuition. Isocurvature initial conditions are defined by the vanishing of the initial potential, Φ(0) = 0. This,
together with equations (23.56) and (23.57), implies the isocurvature initial conditions
Φ(0) = Θ0(0) = 0 , δc(0) = − 8 Φ′(0) = − 8 Θ′0(0) isocurvature . (23.59)
376 Cosmological perturbations: a simplest set of assumptions
Subject to the condition that Φ remains finite at a → 0, the adiabatic and isocurvature solutions to equa-
tion (23.56) are, in units aeq = 1,
Φad =Φ(0)
10
(
9 +2
a− 8
a2− 16
a3+
16√
1 + a
a3
)
=Φ(0)
10
[
3(14 + 9a) + (38 + 9a)√
1 + a]
(
1 +√
1 + a)3 , (23.60a)
Φiso =8 Φ′(0)
5
(
1− 2
a+
8
a2− 16
a3+
16√
1 + a
a3
)
=8 Φ′(0)
5
a(
6 + a+ 4√
1 + a)
(
1 +√
1 + a)4 , (23.60b)
in which the last expressions in each case are written in a form that is numerically well-behaved for all a.
Figure 23.2 shows the evolution of the potential Φ from equations (23.60), normalized to the value Φ(late)
at late times a ≫ aeq. For adiabatic fluctuations, the potential changes by a factor of 9/10 from initial to
final value, while for isocurvature fluctuations the potential evolves from zero to a final value of 85Φ′(0):
Φad(late) = 910Φ(0) = − 3
5δc(0) adiabatic , (23.61a)
Φiso(late) = 85Φ′(0) = − 1
5δc(0) isocurvature . (23.61b)
The superhorizon solutions (23.53) for the dark matter overdensity δc and radiation monopole Θ0 are,
δc = 3Θ0 = 3(
Φad − 32Φ(0)
)
adiabatic , (23.62a)
δc = 3Φiso − 8Φ′(0) , Θ0 = Φiso isocurvature . (23.62b)
23.10 Radiation-dominated, adiabatic initial conditions
For adiabatic initial conditions, fluctuations that enter the horizon before matter-radiation equality, kηeq ≫1, are dominated by radiation. In the regime where radiation dominates both the unperturbed energy and
its fluctuations, the relevant equations are, from equations (23.27), (23.28), and (23.30),
Θ0 − kΘ1 = Φ , (23.63a)
− k2Φ− 3a
aF = 16πGa2ρrΘ0 , (23.63b)
−kF = 16πGa2ρrΘ1 , (23.63c)
in which, because it simplifies the mathematics, the Einstein momentum equation is used as a substitute
for the radiation dipole equation. In the radiation-dominated epoch, the horizon is proportional to the
cosmic scale factor, η ∝ a, equation (23.38). Inserting Θ0 and Θ1 from the Einstein energy and momentum
equations (23.63b) and (23.63c) into the radiation monopole equation (23.63a) gives a second order differential
equation for the potential Φ
Φ +4
ηΦ +
k2
3Φ = 0 . (23.64)
23.10 Radiation-dominated, adiabatic initial conditions 377
0 1 2 3 4 5 6 7 8−2.0−1.5−1.0−.5
.0
.51.01.52.02.53.0
kηs /π
Θ0
−Φ
and
−2
Φ
adiabatic initial conditions
large scales (kηs,eq << 1)
small scales (kηs,eq >> 1)
0 1 2 3 4 5 6 7 8−.5
.0
.5
1.0
kηs /π
Θ0
−Φ
and
−2
Φ
isocurvature initial conditions
large scales (kηs,eq << 1)
small scales (kηs,eq >> 1)
Figure 23.3 The difference Θ0 − Φ between the radiation monopole and the Newtonian scalar potentialoscillates about − 2Φ, in accordance with equation (23.45). The difference (Θ0 − Φ) − (−2Φ) = Θ0 + Φ,which is the temperature Θ0 redshifted by the potential Φ, is the monopole contribution to temperaturefluctuation of the CMB. The top panel is for adiabatic initial conditions, equations (23.58), while the bottompanel is for isocurvature initial conditions, equation (23.59). The units of Φ and Θ0 are such that Φ(0) = −1for adiabatic fluctuations, and δc(0) = 1 for isocurvature fluctuations. In each case, the thin lines showthe evolution of small scale fluctuations, which enter the horizon during the radiation-dominated epoch wellbefore matter-radiation equality, while the thick lines show the evolution of large-scale fluctuations, whichenter the horizon during the matter-dominated epoch well after matter-radiation equality.
Equation (23.64) describes damped sound waves moving at sound speed√
13 times the speed of light. The
sound horizon, the comoving distance that sound can travel, is η/√
3, the horizon distance η multiplied by
378 Cosmological perturbations: a simplest set of assumptions
the sound speed. The growing and decaying solutions to equation (23.64) are
Φgrow =3j1(α)
α=
3(sinα− α cosα)
α3, Φdecay = − j−2(α)
α=
cosα+ α sinα
α3, (23.65)
where the dimensionless parameter α is the wavevector k multiplied by the sound horizon η/√
3,
α ≡ kη√3
=
√
2
3
k
aeqHeq
a
aeq, (23.66)
and jl(α) ≡√
π/(2α)Jl+1/2(α) are spherical Bessel functions. The physically relevant solution that satisfies
adiabatic initial conditions, remaining finite as α→ 0, is the growing solution
Φ = Φ(0)Φgrow . (23.67)
The solution (23.65) shows that, after a mode enters the sound horizon the scalar potential Φ oscillates with
an envelope that decays as α−2. Physically, relativistically propagating sound waves tend to suppress the
gravitational potential Φ.
For the growing solution (23.67), the radiation monopole Θ0 is
Θ0 = Φ(0)3
α3
[
(1− α2) sinα− α(
1− 12α
2)
cosα]
. (23.68)
The thin lines in the top panel of Figure 23.45 show the growing mode potential Φ and the radiation monopole
Θ0, equation (23.68). The Figure plots these two quantities in the form −2Φ and Θ0 − Φ, to bring out the
fact that Θ0 −Φ oscillates about −2Φ, in accordance with equation (23.45). After a mode is well inside the
sound horizon, α≫ 1, the radiation monopole oscillates with constant amplitude,
Θ0 =3Φ(0)
2cosα for α≫ 1 . (23.69)
The dark matter fluctuations are driven by the gravitational potential of the radiation. The solution
of the dark matter equation (23.43) driven by the potential (23.67) and satisfying adiabatic initial condi-
tions (23.62a) is
δc = 3Φ− 9Φ(0)
(
γ − 1
2+ lnα− Ciα+
sinα
α
)
, (23.70)
where the potential Φ is the growing mode solution (23.67), γ ≡ 0.5772... is Euler’s constant, and Ci(α) ≡∫ α
∞ cosxdx/x is the cosine integral. Once the mode is well inside the sound horizon, α≫ 1, the dark matter
density δc, equation (23.70), evolves as
δc = − 9Φ(0)
(
γ − 1
2+ lnα
)
for α≫ 1 , (23.71)
which grows logarithmically. This logarithmic growth translates into a logarithmic increase in the amplitude
of matter fluctuations at small scales, and is a characteristic signature of non-baryonic cold dark matter.
23.11 Radiation-dominated, isocurvature initial conditions 379
23.11 Radiation-dominated, isocurvature initial conditions
For isocurvature initial conditions, the matter fluctuation contributes from the outset, |ρcδc| > |ρrΘ0| even
while radiation dominates the background density, ρc ≪ ρr.
To develop an approximation adequate for isocurvature fluctuations entering the horizon well before
matter-radiation equality, kηeq ≫ 1, regard the Einstein energy equation (23.28a) as giving the radiation
monopole Θ0, and the Einstein momentum equation (23.30) as giving the radiation dipole Θ1. Insert these
into the radiation monopole equation (23.27a), and eliminate the δc terms using the dark matter density
equation (23.26a). The result is, in units aeq = Heq = 1,
2a(1 + a)Φ′′ + (8 + 9a)Φ′ + 2(
1 +2k2a
3
)
Φ + δc = 0 , (23.72)
where prime ′ denotes differentiation with respect to cosmic scale factor a. Equation (23.72) is valid in
all regimes, for any combination of matter and radiation. For isocurvature initial conditions, the radiation
monopole and potential vanish initially, Θ0(0) = Φ(0) = 0, whereas the dark matter overdensity is finite,
δc(0) 6= 0. For small scales that enter the horizon well before matter-radiation equality, kηeq ≫ 1, the
potential Φ is small, while δc has some approximately constant non-zero value, up to and through the time
when the mode enters the horizon, kη ≈ ka ≈ 1. In the radiation-dominated epoch, a≪ 1, and with Φ ≈ 0
(but k large and ka ∼ 1, so k2aΦ is not small) equation (23.72) simplifies to
2aΦ′′ + 8Φ′ +4k2a
3Φ + δc = 0 . (23.73)
The solution of equation (23.73) for constant δc = δc(0) is, with α given by equation (23.66),
Φ = − δc(0)√
2/3 k
1 + 12α
2 − cosα− α sinα
α3. (23.74)
The solution (23.51) for the radiation monopole Θ0 driven by the potential (23.74) is
Θ0 = − δc(0)√
2/3 k
(− 1 + α2) cosα+ (− 1 + 12α
2)(− 1 + α sinα)
α3. (23.75)
Equations (23.74) and (23.75) are the solution for small scale modes with isocurvature initial conditions
that enter the horizon well before matter-radiation equality. After a mode is well inside the sound horizon,
α≫ 1, the radiation monopole (23.75) oscillates with constant amplitude,
Θ0 = − δc(0)
2√
2/3 ksinα α≫ 1 . (23.76)
Whereas for adiabatic initial conditions the radiation monopole oscillated as cosα well inside the hori-
zon, equation (23.69), for isocurvature initial conditions it oscillates as sinα well inside the horizon, equa-
tion (23.76).
380 Cosmological perturbations: a simplest set of assumptions
23.12 Subhorizon scales
After a mode enters the horizon, the radiation fluctuation Θ0 oscillates, but the non-baryonic cold dark
matter fluctuation δc grows monotonically. In due course, the dark matter density fluctuation ρcδc dominates
the radiation density fluctuation ρrΘ0, and this necessarily occurs before matter-radiation equality; that is,
|ρcδc| > |ρrΘ0| even though ρc < ρr. This is true for both adiabatic and isocurvature initial conditions; of
course, for isocurvature initial conditions, the dark matter density fluctuation dominates from the outset.
Even before the dark matter density fluctuation dominates, the cumulative contribution of the dark matter
to the potential Φ begins to be more important than that of the radiation, because the potential sourced by
the radiation oscillates, with an effect that tends to cancel when averaged over an oscillation.
Regard the Einstein energy equation (23.28a) as giving the dark matter overdensity δc, and the Einstein
momentum equation (23.30) as giving the dark matter velocity vc. Insert these into the dark matter density
equation (23.26a) and eliminate the Θ0 terms using the radiation monopole equation (23.27a). The result
is, in units aeq = 1,
2a2(1 + a)Φ′′ + a(6 + 7a)Φ′ − 2Φ− 4Θ0 = 0 , (23.77)
where prime ′ denotes differentiation with respect to cosmic scale factor a. Equation (23.77) is valid in all
regimes, for any combination of matter and radiation. Once the mode is well inside the horizon, kη ≫ 1,
the radiation monopole Θ0 oscillates about an average value of −Φ (since Θ0 − Φ oscillates about −2Φ, as
noted in §23.7):
〈Θ0〉 = −Φ . (23.78)
Inserting this cycle-averaged value of Θ0 into equation (23.77) gives the Meszaros differential equation
2(1 + a)a2Φ′′ + (6 + 7a)aΦ′ + 2Φ = 0 . (23.79)
The solutions of Meszaros’ differential equation (23.79) are
Φ = − 3
4
(
aeqHeq
k
)2δc
a/aeq, (23.80)
where the dark matter overdensity δc is a linear combination
δc = Cgrow δc,grow + Cdecay δc,decay (23.81)
of growing and decaying solutions, in units aeq = 1,
δc,grow = 1 +3
2a , δc,decay =
(
1 +3
2a)
ln(
√1 + a+ 1√1 + a− 1
)
− 3√
1 + a . (23.82)
For adiabatic initial conditions, the desired solution is the one that matches smoothly on to the the logarith-
mically growing solution for the dark matter overdensity δc given by equation (23.70). For modes that enter
the horizon well before matter-radiation equality, the matching may be done in the radiation-dominated
epoch a≪ 1, where the growing and decaying modes (23.82) simplify to
δc,grow = 1 , δdecay = − ln(a/4)− 3 , for a≪ 1 . (23.83)
23.13 Matter-dominated 381
Matching to the solution for δc well inside the horizon, equation (23.71), determines the constants
Cgrow = − 9 Φ(0)
[
γ − 7
2+ ln
(
4
√
2
3
k
aeqHeq
)
]
, Cdecay = 9 Φ(0) adiabatic . (23.84)
For isocurvature initial conditions, for modes that enter the horizon well before matter-radiation equality,
only the growing mode is present,
Cgrow = δc(0) , Cdecay = 0 isocurvature . (23.85)
The dark matter overdensity δc then evolves as the linear combination (23.81) of growing and decaying
modes (23.82). For modes that enter the horizon well before matter-radiation equality, the constants are
set by equation (23.84) for adiabatic intial conditions, or equation (23.85) for isocurvature intial conditions.
The solution remains valid from the radiation-dominated through into the matter-dominated epoch. At late
times well into the matter-dominated epoch, a≫ 1, the growing mode of the Meszaros solution dominates,
δc,grow = 32a , δc,decay = 4
15a−3/2 for a≫ 1 , (23.86)
so that the dark matter overdensity δc at late times is
δc = 32 Cgrow a for a≫ 1 . (23.87)
The potential Φ, equation (23.80), at late times is constant,
Φ = − 9
8k2Cgrow for a≫ 1 . (23.88)
For modes that enter the horizon well before matter-radiation equality, the radiation monopole Θ0 at late
times a≫ 1 is, with α ≡ kη/√
3,
Θ0 = −Φ +3
2Φ(0) cosα adiabatic , (23.89a)
Θ0 = −Φ− δc(0)
2√
2/3 ksinα isocurvature . (23.89b)
23.13 Matter-dominated
After matter-radiation equality, but before curvature or dark energy become important, non-relativistic
matter dominates the mass-energy density of the Universe.
In the matter-dominated epoch, the relevant equations are, from equations (23.26), (23.28a), and (23.30),
δc − k vc = 3 Φ , (23.90a)
− k2Φ− 3a
aF = 4πGa2ρcδc , (23.90b)
−kF = 4πGa2ρcvc , (23.90c)
382 Cosmological perturbations: a simplest set of assumptions
in which, because it simplifies the mathematics, the Einstein momentum equation is used as a substitute
for the matter velocity equation. In the matter-dominated epoch, the horizon is proportional to the square
root of the cosmic scale factor, η ∝ a1/2, equation (23.38). Inserting δc and vc from the Einstein energy and
momentum equations (23.90b) and (23.90c) into the matter density equation (23.90a) yields a second order
differential equation for the potential Φ
Φ +6
ηΦ = 0 . (23.91)
The general solution of equation (23.91) is a linear combination
Φ = Cgrow Φgrow + Cdecay Φdecay (23.92)
of growing and decaying solutions
Φgrow = 1 , Φdecay = α−5 , (23.93)
where the dimensionless paramater α is the wavevector k multiplied by the sound horizon η/√
3,
α ≡ kη√3
= 2
√
2
3
ka1/2
a3/2eq Heq
. (23.94)
The constants Cgrow and Cdecay in the solution (23.92) depend on conditions established before the matter-
dominated epoch. The corresponding growing and decaying modes for the dark matter overdensity δc are
δc,grow = −(
2 +α2
2
)
Φgrow , δc,decay =
(
3− α2
2
)
Φdecay . (23.95)
For modes well inside the horizon, α ≫ 1, the behaviour of the growing and decaying modes agrees with
that (23.86) of the Meszaros solution, as it should. Any admixture of the decaying solution tends quickly to
decay away, leaving the growing solution. The solution (23.51) for the radiation monopole Θ0 driven by the
potential (23.92) is a sum of a homogeneous solution and a particular solution,
Θ0 = B0 cosα+B1 sinα+ Cgrow Θ0,grow + Cdecay Θ0,decay , (23.96)
with growing and decaying modes
Θ0,grow = −Φgrow , Θ0,decay =Φdecay
12
12− 2α2 + α4 + α5 [cosα (Siα− π/2)− sinαCiα]
. (23.97)
23.14 Recombination
Exercise 23.6 Electron-scattering mean free path.
23.15 Post-recombination 383
1. Define the neutron faction Xn by
Xn ≡nn
np + nn, (23.98)
where the proton and neutron number densities np and nn are taken to include protons and neutrons in
all nuclei. For a H plus 4He composition, the proton and neutron number densities are
np = nH + 2n4He , nn = 2n4He . (23.99)
Show that the primordial 4He mass fraction defined by Y4He ≡ ρ4He/(ρH + ρ4He) satisfies
Y4He = 2Xn . (23.100)
The observed primordial 4He abundance is Y4He = 0.24, implying
Xn = 0.12 . (23.101)
2. Define the ionization fraction Xe by
Xe ≡ne
np(23.102)
where again the proton number density np is taken to include protons in all nuclei, not just in hydrogen.
Show thatne
ρb=Xe(1−Xn)
mp(23.103)
where mp is the mass of a proton or neutron.
3. Show that the (dimensionless) ratio of the comoving electron-photon scattering mean free path lT to the
comoving cosmological horizon distance c/(aeqHeq) at matter-radiation equality is
aeqHeqlTc
≡ aeqHeq
cneσTa=
16πGmp
3cσTXe(1−Xn)Heq
Ωm
Ωb
(
a
aeq
)2
=28.94 h−1
Xe(1 −Xn)
H0
Heq
Ωm
Ωb
(
a
aeq
)2
=0.002
Xe
(
a
aeq
)2
, (23.104)
the Hubble parameter Heq at matter-radiation equality being related to the present-day Hubble parameter
H0 by equation (23.41).
23.15 Post-recombination
Recombination frees baryons and photons from each other’s grasp.
Exercise 23.7 Growth of baryon fluctuations after recombination.
384 Cosmological perturbations: a simplest set of assumptions
23.16 Matter with dark energy
Some time after recombination, dark energy becomes important. Observational evidence suggests that the
dominant energy-momentum component of the Universe today is dark energy, with an equation of state
consistent with that of a cosmological constant, pΛ = −ρΛ. In what follows, dark energy is taken to have
constant density, and therefore to be synonymous with a cosmological constant. Since dark energy has a
constant energy density whereas matter density declines as a−3, dark energy becomes important only well
after recombination.
Dark energy does not cluster gravitationally, so the Einstein equations for the perturbed energy-momentum
depend only on the matter fluctuation. However, dark energy does affect the evolution of the cosmic scale
factor a. In fact, if matter is taken to be the only source of perturbation, then covariant energy-momentum
conservation, as enforced by the Einstein equations, implies that the only addition that can be made to the
unperturbed background is dark energy, with constant energy density. To see this, consider the equations
governing the matter overdensity δm and scalar velocity vm (now subscripted m, since post-recombination
matter includes baryons as well as non-baryonic cold dark matter), together with the Einstein energy and
momentum equations:
δm − k vm = 3 Φ , (23.105a)
vm +a
avm = −kΦ , (23.105b)
− k2Φ− 3a
aF = 4πGa2ρmδm , (23.105c)
−kF = 4πGa2ρmvm . (23.105d)
The factor 4πGa2ρm on the right hand side of the two Einstein equations can be written
4πGa2ρm =3a3
0H20Ωm
2a, (23.106)
where a0 and H0 are the present-day cosmic scale factor and Hubble parameter, and Ωm is the present-
day matter density (a constant). Allow the Hubble parameter H(a) ≡ a/a2 to be an arbitrary function
of cosmic scale factor a. Inserting δm and velocity vm from the Einstein energy and momentum equa-
tions (23.105c) and (23.105d) into the matter equations (23.105a) and (23.105b), and taking the overdensity
equation (23.105a) minus 3a/a times the velocity equation (23.105b), yields the condition
a4 dH2
da+ 3a3
0H20Ωm = 0 , (23.107)
whose solution is
H2
H20
=Ωm
(a/a0)3+ ΩΛ (23.108)
for some constant ΩΛ. This shows that, as claimed, if only matter perturbations are present, then the unper-
turbed background can contain, besides matter, only dark energy with constant density ρΛ = H20ΩΛ/(
83πG).
23.17 Matter with dark energy and curvature 385
The result is a consequence of the fact that the Einstein equations enforce covariant conservation of energy-
momentum.
With the Hubble parameter given by equation (23.108), the matter and Einstein equations (23.105) yield
a second order differential equation for the potential Φ, in units a0 = 1:
2a(Ωm + a3ΩΛ)Φ′′ + (7Ωm + 10a3ΩΛ)Φ′ + 6a2ΩΛΦ = 0 . (23.109)
The growing and decaying solutions to equation (23.109) are, in units a0 = 1,
Φgrow =5ΩmH
20
2
H(a)
a
∫ a
0
da′
a′3H(a′)3, Φdecay =
H
a. (23.110)
The factor 52ΩmH
20 in the growing solution is chosen so that Φgrow → 1 as a → 0. The growing solution
Φgrow can be expressed as an elliptic integral. The corresponding growing and decaying solutions for the
matter overdensity δm are, again in units a0 = 1,
δm,grow =
(
3− 2k2a
3ΩmH20
)
Φgrow − 5 , δm,decay =
(
3− 2k2a
3ΩmH20
)
Φdecay . (23.111)
For modes well inside the horizon, kη ∼ ka1/2/H0 ≫ 1, the relation (23.111) agrees with that (23.117) below.
23.17 Matter with dark energy and curvature
Curvature may also play a role after recombination. Observational evidence as of 2010 is consistent with the
Universe having zero curvature, but it is possible that there may be some small curvature.
If the curvature is non-zero, then strictly the unperturbed metric should be an FRW metric with curvature.
However, a flat background FRW metric remains a good approximation for modes whose scales are small
compared to the curvature, that is, for modes that are well inside the horizon, kη ≫ 1. For modes well inside
the horizon, the time derivative of the potential can be neglected, Φ = 0. With matter, curvature, and dark
energy present, and for modes well inside the horizon, the equations go over to the Newtonian limit:
δm − k vm = 0 , (23.112a)
vm +a
avm = −kΦ , (23.112b)
− k2Φ = 4πGa2ρmδm . (23.112c)
The factor 4πGa2ρm in the Einstein equation can be written as equation (23.106). The matter and Einstein
equations (23.112) yield a second order equation for the matter overdensity δm, in units a0 = 1:
δm +a
aδm −
3ΩmH20
2
δma
= 0 . (23.113)
Equation (23.113) can be recast as a differential equation with respect to cosmic scale factor a:
δ′′m +
(
H ′
H+
3
a
)
δ′m −3ΩmH
20
2
δma5H2
= 0 , (23.114)
386 Cosmological perturbations: a simplest set of assumptions
where H ≡ a/a2 is the Hubble parameter, and prime ′ denotes differentiation with respect to a. In the case
of matter plus curvature plus dark energy, the Hubble parameter H satisfies, again in units a0 = 1,
H2
H20
= Ωma−3 + Ωka
−2 + ΩΛ , (23.115)
where Ωm, Ωk, and ΩΛ are the (constant) present-day values of the matter, curvature, and dark energy
densities. The growing and decaying solutions to equation (23.114) are
δm,grow ≡ a g(a) =5ΩmH
20
2H(a)
∫ a
0
da′
a′3H(a′)3, δm,decay =
H
H0. (23.116)
The potential Φ is related to the matter overdensity δm by, again in units a0 = 1, equation (23.112c),
Φ = −3ΩmH20
2k2
δma
. (23.117)
The observationally relevant solution is the growing mode. The growing mode is conventionally given a
special notation, the growth factor g(a), because of its importance to relating the amplitude of clustering at
various times, from recombination up to the present. For the growing mode,
δ ∝ a g(a) , Φ ∝ g(a) . (23.118)
The normalization factor 52ΩmH
20 in equation (23.116) is chosen so that in the matter-dominated phase
shortly after recombination (small a), the growth factor g(a) is
g(a) = 1 . (23.119)
Thus as long as the Universe remains matter-dominated, the potential Φ remains constant. Curvature or
dark energy causes the potential Φ to decrease.
It should be emphasized that the growing and decaying solutions (23.116) are valid only for the case of
matter plus curvature plus constant density dark energy, where the Hubble parameter takes the form (23.115).
If another kind of mass-energy is considered, such as dark energy with non-constant density, then equations
governing perturbations of the other kind must be adjoined, and the Einstein equations modified accordingly.
The growth factor g(a) may expressed analytically as an elliptic function. A good analytic approximation
is (Carroll, Press & Turner 1992, Ann. Rev. Astron. Astrophys. 30, 449)
g ≈ 5Ωm
2[
Ω4/7m − ΩΛ +
(
1 + 12Ωm
) (
1 + 170ΩΛ
)
] , (23.120)
where Ωa are densities at the epoch being considered (such as the present, a = a0).
24
∗Cosmological perturbations: a more carefultreatment of photons and baryons
The “simple” treatment of cosmological perturbations in the previous Chapter is sufficient to reveal that the
photon-baryon fluid at the time of recombination shows a characteristic pattern of oscillations. Translating
this pattern into something that can be compared to observations of the CMB requires a more careful
treatment that follows the evolution of photons using a collisional Boltzmann equation.
Cosmologists conventionally refer to atomic matter — anything that acts by either strong or electromag-
netic forces — as baryons (from the greek baryos, meaning heavy), even though they mean by that not only
baryons (protons, neutrons, and other nuclei), but also electrons, which are leptons (from the greek leptos,
meaning light). The designation baryons does not include relativistic species, such as photons and neutrinos,
nor does it include non-baryonic dark matter or dark energy. The designation “baryons” is nonsensical, but
has stuck.
Although baryons are gravitationally sub-dominant, having a mass density about 1/5 that of the non-
baryonic dark matter, they play an important role in CMB fluctuations. Most importantly, before recombi-
nation atomic matter is ionized, and the free electrons scatter photons, preventing photons from travelling
far. Recombination occurs when the temperature drops to the point that electrons combine into neutral
atoms, releasing the photons to travel freely across the Universe, into astronomers’ telescopes.
Electron-photon scattering keeps photons and baryons tightly coupled, so that up to the time of recom-
bination they oscillate together as a photon-baryon fluid. As recombination approaches, the baryon density
becomes increasingly important relative to the photon density. The baryons, which provide mass but no
pressure, decrease the sound speed of the photon-baryon fluid, and their gravity enhances sound wave com-
pressions while weakening rarefactions. As recombination approaches, the mean free path of photons to
scattering increases, which tends to damp sound waves at short scales. All these effects — a decreased sound
speed, an enhancement of compression over rarefaction, and damping at small scales — produce observable
signatures in the power spectrum of temperature fluctuations in the CMB.
388 ∗Cosmological perturbations: a more careful treatment of photons and baryons
24.1 Lorentz-invariant spatial and momentum volume elements
DO THIS BETTER. To define an occupation number in a Lorentz-invariant fashion, it is first necessary
to define Lorentz-invariant volume elements of space and momentum. With respect to an orthonormal
tetrad, volume elements transform as they do in special relativity. With respect to an orthonormal tetrad,
a Lorentz-invariant spatial 3-volume element can be constructed as
d4x
dτ=dt
dτd3x = E d3x . (24.1)
Lorentz invariance of Ed3x is evident because the tetrad-frame 4-volume element d4x is a scalar, and likewise
the interval dτ of proper time is a scalar. Similarly, with respect to an orthonormal tetrad, a Lorentz-invariant
momentum-space 3-volume element can be constructed as∫
E>0
2 δD(E2 − p2 −m2) d4p =
∫
E>0
2 δD(E2 − p2 −m2) dE d3p =d3p
E. (24.2)
Lorentz-invariance of d3p/E is evident because d4p is a Lorentz-invariant momentum 4-volume, and the
argument E2 − p2 −m2 of the delta-function is a scalar. From the Lorentz invariance of Ed3x and d3p/E it
follows that the product
d3xd3p (24.3)
of spatial and momentum 3-volumes is Lorentz-invariant.
24.2 Occupation numbers
WARNING: d3x IS TETRAD-FRAME VOLUME, BUT x IS A COORDINATE. Each species of energy-
momentum is described by a dimensionless occupation number, or phase-space distribution, a function
f(η,x,p) of conformal time η, comoving position x, and tetrad-frame momentum p, which describes the
number dN of particles in a tetrad-frame element d3xd3p/(2π~)3 of phase-space
dN(η,x,p) = f(η,x,p)g d3xd3p
(2π~)3, (24.4)
with g being the number of spin states of the particle. The tetrad-frame phase-space element d3xd3p/(2π~)3
is dimensionless and Lorentz invariant, and the occupation number f is likewise dimensionless and Lorentz
invariant. The tetrad-frame energy-momentum 4-vector pm of a particle is
pm ≡ emµdxµ
dλ= E,p = E, pi , (24.5)
where λ is the affine parameter, related to proper time τ along the worldline of the particle by dλ ≡ dτ/m,
which remains well-defined in the limit of massless particles, m = 0. The tetrad-frame energy E and
momentum p ≡ |p| for a particle of rest mass m are related by
E2 − p2 = m2 . (24.6)
24.3 Occupation numbers in thermodynamic equilibrium 389
The tetrad-frame components of the energy-momentum tensor Tmn of any species are integrals over its
occupation number f weighted by the product pmpn of 4-momenta:
Tmn =
∫
f pmpn g d3p
E(2π~)3. (24.7)
The energy-momentum tensor Tmn defined by equation (24.7) is manifestly a tetrad-frame tensor, thanks
to the Lorentz-invariance of the momentum-space 3-volume element d3p/E.
24.3 Occupation numbers in thermodynamic equilibrium
Frequent collisions tend to drive a system towards thermodynamic equilibrium. Electron-photon scatter-
ing keeps photons in near equilibrium with electrons, while Coulomb scattering keeps electrons in near
equilibrium with ions, primarily hydrogen ions (protons) and helium nuclei. Thus photons and baryons
can be treated as having unperturbed distributions in mutual thermodynamic equilibrium, and perturbed
distributions that are small departures from thermodynamic equilibrium.
In thermodynamic equilibrium at temperature T , the occupation numbers of fermions, which obey an
exclusion principle, and of bosons, which obey an anti-exclusion principle, are
f =
1
e(E−µ)/T + 1fermion ,
1
e(E−µ)/T − 1boson ,
(24.8)
where µ is the chemical potential of the species. In the limit of small occupation numbers, f ≪ 1, equivalent
to large negative chemical potential, µ → −large, both fermion and boson distributions go over to the
Boltzmann distribution
f = e(−E+µ)/T Boltzmann . (24.9)
Chemical potential is the thermodynamic potential assocatiated with conservation of number. There is a
distinct potential for each conserved species. For example, photoionization and radiative recombination of
hydrogen,
H + γ ↔ p+ e , (24.10)
separately preserves proton and electron number, hydrogen being composed of one proton and one electron.
In thermodynamic equilibrium, the chemical potential µH of hydrogen is the sum of the chemical potentials
µp and µe of protons and electrons
µH = µp + µe . (24.11)
Photon number is not conserved, so photons have zero chemical potential,
µγ = 0 , (24.12)
which is closely associated with the fact that photons are their own antiparticles.
390 ∗Cosmological perturbations: a more careful treatment of photons and baryons
24.4 Boltzmann equation
The evolution of each species is described by the general relativistic Boltzmann equation
df
dλ= C[f ] , (24.13)
where C[f ] is a collision term. The derivative with respect to affine parameter λ on the left hand side of
the Boltzmann equation (24.13) is a Lagrangian derivative along the (timelike or lightlike) worldline of a
particle in the fluid. Since both the occupation number f and the affine parameter λ are Lorentz scalars,
the collision term C[f ] is a Lorentz scalar. In the absence of collisions, the collisionless Boltzmann equation
df/dλ = 0 expresses conservation of particle number: a particle is neither created nor destroyed as it moves
along its wordline.
The left hand side of the Boltzmann equation (24.13) is
df
dλ= pm∂mf +
dpi
dλ
∂f
∂pi= E∂0f + pi∂if +
dp
dλ· ∂f∂p
+dp
dλ
∂f
∂p. (24.14)
Both dp/dλ and ∂f/∂p vanish in the unperturbed background, so dp/dλ ·∂f/∂p is of second order, and can
be neglected to linear order, so that
df
dλ= E∂0f + pi∂if +
dp
dλ
∂f
∂p. (24.15)
The expression (24.15) for the left hand side df/dλ of the Boltzmann equation involves dp/dλ, which in
free-fall is determined by the usual geodesic equation
dpk
dλ+ Γk
mn pmpn = 0 . (24.16)
Since E2−p2 = m2, it follows that the equation of motion for the magnitude p of the tetrad-frame momentum
is related to the equation of motion for the tetrad-frame energy E by
pdp
dλ= E
dE
dλ. (24.17)
The equation of motion for the tetrad-frame energy E ≡ p0 is
dE
dλ= −Γ0
mn pmpn = Γ0i0 p
iE + Γ0ij pipj . (24.18)
From this it follows that
d ln p
dλ=E
p2
dE
dλ= E
(
Epi
pΓ0i0 + pipjΓ0ij
)
= E
(
− a
a2+Epi
pΓ0i0 + pipj
1
Γ0ij
)
, (24.19)
where in the last expression the tetrad connection Γ0ij , equation (22.14b), has been separated into its
unperturbed and perturbed parts −(a/a2)δij and1
Γ0ij .
24.4 Boltzmann equation 391
In practice, the integration variable used to evolve equations is the conformal time η, not the affine
parameter λ. The relation between conformal time η and affine parameter λ is
dη
dλ= pη = em
ηpm = (δnm + ϕm
n)0
enηpm =
1
a
[
E(1− ϕ00)− piϕi0
]
, (24.20)
whose reciprocal is to linear order
dλ
dη=
a
E
(
1 + ϕ00 +pi
Eϕi0
)
. (24.21)
With conformal time η as the integration variable, the equation of motion (24.19) for the magnitude p of
the tetrad-frame momentum becomes, to linear order,
d ln p
dη= − a
a
(
1 + ϕ00 +pi
Eϕi0
)
+Epi
paΓ0i0 + pipja
1
Γ0ij . (24.22)
With respect to conformal time η, the Boltzmann equation (24.13) is
df
dη=∂f
∂η+ vi∇if +
d ln p
dη
∂f
∂ ln p=dλ
dηC[f ] , (24.23)
with dλ/dη from equation (24.21), and d ln p/dη from equation (24.22). Expressions for dλ/dη and d ln p/dη
in terms of the vierbein perturbations in a general gauge are left as Exercise 24.1. In conformal Newtonian
gauge, the factor dλ/dη, equation (24.21), is
dλ
dη=
a
E(1 + Ψ) . (24.24)
In conformal Newtonian gauge, and including only scalar fluctuations, the factor d ln p/dη, equation (24.22),
is
d ln p
dη= − a
a+ Φ− Epi
p∇iΨ . (24.25)
To unperturbed order, the Boltzmann equation (24.23) is
d0
f
dη=∂
0
f
∂η− a
a
∂0
f
∂ ln p=
a
EC[
0
f ] , (24.26)
where C[0
f ] is the unperturbed collision term, the factor a/E coming from dλ/dη = a/E to unperturbed
order, equation (24.21). The second term in the middle expression of equation (24.26) simply reflects the
fact that the tetrad-frame momentum p redshifts as p ∝ 1/a as the Universe expands, a statement that is
true for both massive and massless particles.
Subtracting off the unperturbed part (24.26) of the Boltzmann equation (24.23) gives the perturbation of
the Boltzmann equation
d1
f
dη=∂
1
f
∂η+ vi∇i
1
f − a
a
∂1
f
∂ ln p+d ln(ap)
dη
∂0
f
∂ ln p=
a
EC[
1
f ] +
1
dλ
dηC[
0
f ] . (24.27)
392 ∗Cosmological perturbations: a more careful treatment of photons and baryons
In conformal Newtonian gauge, the perturbed part of dλ/dη is
1
dλ
dη=
a
EΨ . (24.28)
In conformal Newtonian gauge, and including only scalar fluctuations, d ln(ap)/dη is
d ln(ap)
dη= Φ− Epi
p∇iΨ . (24.29)
Exercise 24.1 Boltzmann equation factors in a general gauge. Show that in a general gauge, and
including not just scalar but also vector and tensor fluctuations, equation (24.21) is
dλ
dη=
a
E
[
1 + ψ +pi
E(∇iw + wi)
]
, (24.30)
while equation (24.22) is, with only scalar fluctuations included,
d ln p
dη= − a
a+ φ+
Epi
p
[
−∇iψ +
(
∂
∂η+a
a
m2
E2
)
(∇iw + wi)
]
+ pipj[
−∇i∇j(w − h)− 12 (∇iWj +∇jWi) +∇jwi + hij
]
. (24.31)
24.5 Non-baryonic cold dark matter
Non-baryonic cold dark matter, subscripted c, is by assumption non-relativistic and collisionless. The un-
perturbed mean density is ρc, which evolves with cosmic scale factor a as
ρc ∝ a−3 . (24.32)
Since dark matter particles are non-relativistic, the energy of a dark matter particle is its rest-mass energy,
Ec = mc, and its momentum is the non-relativistic momentum pic = mcv
ic.
The energy-momentum tensor Tmnc of the dark matter is obtained from integrals over the dark matter
phase-space distribution fc, equation (24.7). The energy and momentum moments of the distribution define
the dark matter overdensity δc and bulk velocity vvvc, while the pressure is of order v2c , and can be neglected
to linear order,
T 00c ≡
∫
fcmcgc d
3pc
(2π~)3≡ ρc(1 + δc) , (24.33a)
T 0ic ≡
∫
fcmc vic
gc d3pc
(2π~)3≡ ρcv
ic , (24.33b)
T ijc ≡
∫
fcmc vic v
jc
gc d3pc
(2π~)3= 0 . (24.33c)
24.5 Non-baryonic cold dark matter 393
Non-baryonic cold dark matter is collisionless, so the collision term in the Boltzmann equation is zero,
C[fc] = 0, and the dark matter satisfies the collisionless Boltzmann equation
dfc
dη= 0 . (24.34)
The energy and momentum moments of the Boltzmann equation (24.23) yield equations for the overdensity
δc and bulk velocity vc, which in the conformal Newtonian gauge are
0 =
∫
dfc
dηmc
gc d3pc
(2π~)3=
∂
∂η
∫
fcmcgc d
3pc
(2π~)3+∇i
∫
fcmcvic
gc d3pc
(2π~)3−∫ (
a
a− Φ
)
∂f
∂ ln pmc
gc d3pc
(2π~)3
=∂ρc(1 + δc)
∂η+∇i(ρcv
ic) + 3
(
a
a− Φ
)
ρc , (24.35a)
0 =
∫
dfc
dηmcv
ic
gc d3pc
(2π~)3=
∂
∂η
∫
fcmcvic
gc d3pc
(2π~)3+∇j
∫
fcmcvicv
jc
gc d3pc
(2π~)3
−∫ (
a
a− Φ +
Epj
p∇jΨ
)
∂f
∂ ln pmcv
i gc d3pc
(2π~)3
=∂ρcv
ic
∂η+ 4
(
a
a− Φ
)
ρcvic + ρc∇iΨ . (24.35b)
The Φρcvic term on the last line of equation (24.35b) can be dropped, since the potential Φ and the velocity
v ic are both of first order, so their product is of second order. Subtracting the unperturbed part from
equations (24.35a) and (24.35b) gives equations for the dark matter overdensity δc and velocity vvvc,
δc + ∇ · vvvc − 3Φ = 0 , (24.36a)
vvvc +a
avvvc + ∇Ψ = 0 . (24.36b)
Transform into Fourier space, and decompose the velocity 3-vector vvvc into scalar vc and vector vvvc,⊥ parts
vvvc = −ikvc + vvvc,⊥ . (24.37)
For the scalar modes under consideration, only the scalar part of the dark matter equations (24.36) is
relevant:
δc − kvc − 3Φ = 0 , (24.38a)
vc +a
avc + kΨ = 0 . (24.38b)
Equations (24.38) reproduce the equations (23.26) derived previously from conservation of energy and mo-
mentum.
Exercise 24.2 Moments of the non-baryonic cold dark matter Boltzmann equation. Confirm
equations (24.35).
394 ∗Cosmological perturbations: a more careful treatment of photons and baryons
24.6 The left hand side of the Boltzmann equation for photons
In the unperturbed background, the photons have a blackbody distribution with temperature T (η). Define
Θ to be the photon temperature fluctuation
Θ(η,x,pγ) ≡ δT (η,x,pγ)
T (η). (24.39)
In the unperturbed background, the photon occupation number is
0
fγ =1
epγ/T − 1. (24.40)
Since pγ ∝ T ∝ 1/a, the unperturbed occupation number is constant as a function of pγ/T . The definition
Θ ≡ δT/T = δ lnT of the photon perturbation is to be interpreted as meaning that the perturbation to the
occupation number of photons is (the partial derivative with respect to temperature ∂/∂ lnT is at constant
photon momentum pγ)
1
fγ =∂
0
fγ
∂ lnTδ lnT =
∂0
fγ
∂ lnTΘ , (24.41)
in which it follows from equation (24.40) that
∂0
fγ
∂ lnT=
0
fγ(1 +0
fγ)pγ
T. (24.42)
The photon Boltzmann equation in terms of the occupation number fγ can be recast as a Boltzmann
equation for the temperature fluctuation Θ through
d1
fγ
dη=
d
dη
∂0
fγ
∂ lnTΘ
=d
dη
∂0
fγ
∂ lnT
Θ +∂
0
fγ
∂ lnT
dΘ
dη=
∂0
fγ
∂ lnT
dΘ
dη, (24.43)
in which the first term of the penultimate expression vanishes because ∂0
fγ/∂ lnT is a function of pγ/T only,
and pγ/T is, to unperturbed order (which is all that is needed since the term is multiplied by Θ, which is
already of first order), independent of time, d(pγ/T )/dη = 0, since pγ ∝ T ∝ a−1:
d
dη
∂0
fγ
∂ lnT= − d ln(pγ/T )
dη
∂20
fγ
∂ lnT 2= 0 . (24.44)
In terms of the temperature fluctuation Θ, the perturbed photon Boltzmann equation (24.27) is
dΘ
dη=∂Θ
∂η+ pi
γ ∇iΘ−a
a
∂Θ
∂ ln pγ− d ln(apγ)
dη=
(
a
pγC[
1
fγ ] +
1
dλ
dηC[
0
fγ ]
)/
∂0
fγ
∂ lnT, (24.45)
where the d ln(apγ)/dη term gets a minus sign from ∂0
f/∂ ln pγ = −∂0
f/∂ lnT .
The unperturbed photon distribution is in thermodynamic equilibrium, so the unperturbed collision term
24.7 Spherical harmonics of the photon distribution 395
in the photon Boltzmann equation (24.45) vanishes, C[0
fγ ] = 0, as found in Exercise 24.3 below. The
photon distribution is modified by photon-electron (Thomson) scattering, §24.10. Since the electrons are
non-relativistic, to linear order collisions change the photon momentum but not the photon energy. As a
consequence, the temperature fluctuation is a function Θ(η,x, pγ) only of the direction pγ of the photon
momentum pγ , not of its magnitude, the energy pγ . This is shown more carefully below, equation (24.78).
Hence the derivative ∂Θ/∂ ln pγ in the photon Boltzmann equation (24.45) vanishes to linear order. Thus
the photon Boltzmann equation (24.45) reduces to
dΘ
dη=∂Θ
∂η+ pi
γ ∇iΘ−d ln(apγ)
dη=
a
pγC[
1
fγ ]/ ∂
0
fγ∂ lnT
. (24.46)
In conformal Newtonian gauge, and including only scalar fluctuations, d ln(apγ)/dη is given by equa-
tion (24.29), so the photon Boltzmann equation (24.45) becomes
dΘ
dη=∂Θ
∂η+ pi
γ ∇iΘ− Φ + piγ∇iΨ =
a
pγC[
1
fγ ]/ ∂
0
fγ
∂ lnT. (24.47)
24.7 Spherical harmonics of the photon distribution
It is natural to expand the temperature fluctuation Θ in spherical harmonics. As seen below, equa-
tions (24.55), the various components of the photon energy-momentum tensor Tmnγ are determined by the
monopole, dipole, and quadrupole harmonics of the photon distribution. Scalar fluctuations are those that
are rotationally symmetric about the wavevector direction k, which correspond to spherical harmonics with
zero azimuthal quantum number, m = 0. Expanded in spherical harmonics, and with only scalar terms
retained, the temperature fluctuation Θ can be written
Θ(η,k, pγ) =
∞∑
ℓ=0
(−i)ℓ(2ℓ+ 1)Θℓ(η,k)Pℓ(k · pγ) , (24.48)
where Pℓ are Legendre polynomials, §24.21. The scalar harmonics Θℓ are angular integrals of the temperature
fluctuation Θ over photon directions pγ :
Θℓ(η,k) = iℓ∫
Θ(η,k, pγ)Pℓ(k · pγ)dopγ
4π. (24.49)
Expanded into the scalar harmonics Θℓ(η,k), the left hand side of the photon Boltzmann equation (24.46)
396 ∗Cosmological perturbations: a more careful treatment of photons and baryons
is, in conformal Newtonian gauge,
dΘ0
dη= Θ0 − kΘ1 − Φ , (24.50)
dΘ1
dη= Θ1 +
k
3
(
Θ0 − 2k2Θ2
)
+k
3Ψ , (24.51)
dΘℓ
dη= Θℓ +
k
2ℓ+ 1[ℓΘℓ−1 − (ℓ + 1)Θℓ+1] (ℓ ≥ 2) . (24.52)
24.8 Energy-momentum tensor for photons
Perturbations to the photon energy-momentum tensor involve integrals (24.7) over the perturbed occupation
number of the form, where F (p) is some arbitrary function of the momentum direction p,
∫
1
fγ p2γ F (pγ)
2 d3pγ
pγ(2π~)3=
∫
∂0
fγ
∂ lnTp2
γ
2 4πp2γdpγ
pγ(2π~)3
∫
ΘF (pγ)dopγ
4π= 4ργ
∫
ΘF (pγ)dopγ
4π, (24.53)
in which the last expression is true because
∫
∂0
fγ
∂ lnTp2
γ
2 4πp2γdpγ
pγ(2π~)3= 4
∫
0
fγ pγ
2 4πp2γdpγ
(2π~)3= 4ργ , (24.54)
which follows from ∂0
fγ/∂ lnT = − ∂0
fγ/∂ ln pγ and an integration by parts. The perturbation of the photon
energy density, energy flux, monopole pressure, and quadrupole pressure are
1
T 00γ = 4 ργ Θ0 , (24.55a)
ki
1
T 0iγ = i 4 ργ Θ1 , (24.55b)
13 δij
1
T ijγ = 4
3 ργ Θ0 , (24.55c)(
32 kikj − 1
2 δij
)
1
T ijγ = − 4 ργ Θ2 . (24.55d)
24.9 Collisions
For a 2-body collision of the form
1 + 2 ↔ 3 + 4 , (24.56)
24.9 Collisions 397
the rate per unit time and volume at which particles of type 1 leave and enter an interval d3p1 of momentum
space is, in units c = 1,
C[f1]g1 d
3p1
E1(2π~)3=
∫
|M|2[
− f1f2(1∓ f3)(1 ∓ f4) + f3f4(1 ∓ f1)(1∓ f2)]
(2π~)4δ4D(p1 + p2 + p3 + p4)g1 d
3p1
2E1(2π~)3g2 d
3p2
2E2(2π~)3g3 d
3p3
2E3(2π~)3g4 d
3p4
2E4(2π~)3. (24.57)
All factors in equation (24.57) are Lorentz scalars. On the left hand side, the collision term C[f1] and
the momentum 3-volume element d3p1/E1 are both Lorentz scalars. On the right hand side, the squared
amplitude |M|2, the various occupation numbers fa, the energy-momentum conserving 4-dimensional Dirac
delta-function δ4D(p1 + p2 + p3 + p4), and each of the four momentum 3-volume elements d3pa/Ea, are all
Lorentz scalars.
The first ingredient in the integrand on the right hand side of the expression (24.57) is the Lorentz-invariant
scattering amplitude squared |M|2, calculated using quantum field theory (F. Halzen & A. D. Martin 1984
Quarks and Leptons, Wiley, New York, p. 91).
The second ingredient in the integrand on the right hand side of expression (24.57) is the combination of
rate factors
rate(1 + 2 → 3 + 4) ∝ f1f2(1∓ f3)(1 ∓ f4) , (24.58a)
rate(1 + 2 ← 3 + 4) ∝ f3f4(1∓ f1)(1 ∓ f1) , (24.58b)
where the 1 ∓ f factors are blocking or stimulation factors, the choice of ∓ sign depending on whether the
species in question is fermionic or bosonic:
1− f = Fermi-Dirac blocking factor , (24.59a)
1 + f = Bose-Einstein stimulation factor . (24.59b)
The first rate factor (24.58a) expresses the fact that the rate to lose particles from 1 + 2→ 3 + 4 collisions
is proportional to the occupancy f1f2 of the initial states, modulated by the blocking/stimulation factors
(1 ∓ f3)(1 ∓ f4) of the final states. Likewise the second rate factor (24.58b) expresses the fact that the
rate to gain particles from 1 + 2 ← 3 + 4 collisions is proportional to the occupancy f3f4 of the initial
states, modulated by the blocking/stimulation factors (1∓ f1)(1∓ f2) of the final states. In thermodynamic
equilibrium, the rates (24.58) balance, Exercise 24.3, a property that is called detailed balance, or microscopic
reversibility. Microscopic reversibility is a consequence of time reversal symmetry.
The final ingredient in the integrand on the right hand side of expression (24.57) is the 4-dimensional
Dirac delta-function, which imposes energy-momentum conservation on the process 1 + 2 ↔ 3 + 4. The
4-dimensional delta-function is a product of a 1-dimensional delta-function expressing energy conservation,
and a 3-dimensional delta-function expressing momentum conservation:
(2π~)4δ4D(p1 + p2 + p3 + p4) = 2π~ δD(E1 + E2 + E3 + E4) (2π~)3δ3D(p1 + p2 + p3 + p4) . (24.60)
398 ∗Cosmological perturbations: a more careful treatment of photons and baryons
Exercise 24.3 Detailed balance. Show that the rates balance in thermodynamic equilibrium,
f1f2(1 ∓ f3)(1∓ f4) = f3f4(1∓ f1)(1∓ f2) . (24.61)
Solution. Equation (24.61) is true if and only if
f11∓ f1
f21∓ f2
=f3
1∓ f3f4
1∓ f4. (24.62)
But
f
1∓ f = e(−E+µ)/T , (24.63)
so (24.62) is true if and only if
−E1 + µ1
T+−E2 + µ2
T=−E3 + µ3
T+−E4 + µ4
T, (24.64)
which is true in thermodynamic equilibrium because
E1 + E2 = E3 + E4 , µ1 + µ2 = µ3 + µ4 . (24.65)
24.10 Electron-photon scattering
The dominant process that couples photons and baryons is electron-photon scattering
e+ γ ↔ e′ + γ′ . (24.66)
The Lorentz-invariant transition probability for unpolarized non-relativistic electron-photon (Thomson) scat-
tering is
|M|2 = 12πm2eσT~
2[
1 + (pγ · pγ′)2]
= 16πm2eσT~
2[
1 + 12P2(pγ · pγ′)
]
, (24.67)
where P2(µ) is the quadrupole Legendre polynomial, §24.21, and
σT =8π
3
(
e2
mec2
)2
(24.68)
is the Thomson cross-section.
24.11 The photon collision term for electron-photon scattering 399
24.11 The photon collision term for electron-photon scattering
Electron-photon scattering keeps electrons and photons close to mutual thermodynamic equilibrium, and
their unperturbed distributions can be taken to be in thermodynamic equilibrium. The unperturbed photon
collision term for electron-photon scattering therefore vanishes, because of detailed balance, Exercise 24.3,
C[0
fγ ] = 0 . (24.69)
Thanks to detailed balance, the combination of rates in the collision integral (24.57) almost cancels, so can
be treated as being of linear order in perturbation theory. This allows other factors in the collision integral
to be approximated by their unperturbed values.
The photon collision term for electron-photon scattering follows from the general expression (24.57). To
unperturbed order, the energies of the electrons, which are non-relativistic, may be set equal to their rest
masses, Ee = me. Since photons are massless, their energies are just equal to their momenta, Eγ = pγ . The
electron occupation number is small, fe ≪ 1, so the Fermi blocking factors for electrons may be neglected,
1 − fe = 0. The number of spins of the incoming electron is two, ge = 2, because photons scatter off both
spins of electrons. On the other hand, the number of spins of the scattered electron and photon are one,
ge′ = gγ′ = 1, because non-relativistic electron-photon scattering leaves the spins of the electron and photon
unchanged. These considerations bring the photon collision term for electron-photon scattering to, from the
general expression (24.57),
C[1
fγ ] =1
16
∫
|M|2[
− fefγ(1+fγ′)+fe′fγ′(1+fγ)]
(2π~)4δ4D(pe+pγ−pe′−pγ′)2 d3pe
me(2π~)3d3pe′
me(2π~)3d3pγ′
pγ′(2π~)3.
(24.70)
The various integrations over momenta are most conveniently carried out as follows. The energy-conserving
integral is best done over the energy of the scattered photon γ′, which is scattered into an interval doγ′ of
solid angle:∫
2π~ δD(Ee + Eγ − Ee′ − Eγ′)d3pγ′
Eγ′(2π~)3= pγ′
doγ′
(2π~)2≈ pγ
doγ′
(2π~)2. (24.71)
The approximation in the last step of equation (24.71), replacing the energy pγ′ of the scattered photon by
the energy pγ of the incoming photon, is valid because, thanks to the smallness of the combination of rates
in the collision integral (24.70), it suffices to treat the photon energy to unperturbed order. As seen below,
equation (24.76) the energy difference pγ−pγ′ between the incoming and scattered photons is of linear order
in electron velocities.
The momentum-conserving integral is best done over the momentum of the scattered electron, which is e′
for outgoing scatterings e+ γ → e′ + γ′, and e for incoming scatterings e+ γ ← e′ + γ′. In the former case
(e+ γ → e′ + γ′),∫
(2π~)3δ3D(pe + pγ − pe′ − pγ′)d3pe′
Ee′ (2π~)3=
1
me, (24.72)
and the result is the same, 1/me, in the latter case (e + γ ← e′ + γ′). The energy- and momentum-
conserving integrals having been done, the electron e′ in the latter case may be relabelled e. So relabelled,
400 ∗Cosmological perturbations: a more careful treatment of photons and baryons
the combination of rate factors in the collision integral (24.70) becomes
− fefγ(1 + fγ′) + fefγ′(1 + fγ) = fe(− fγ + fγ′) . (24.73)
Notice that the stimulated terms cancel. The energy- and momentum-conserving integrations (24.71)
and (24.72) bring the photon collision term (24.70) to
C[1
fγ ] =pγ
16πm2e~2
∫
|M|2fe(− fγ + fγ′)2 d3pe
(2π~)3doγ′
4π. (24.74)
The collision integral (24.74) involves the difference − fγ + fγ′ between the occupancy of the initial and
final photon states. To linear order, the difference is
− fγ + fγ′ = −0
f(pγ) +0
f(pγ′)−1
f(pγ) +1
f(pγ′) =∂
0
fγ
∂ lnT
[
pγ − pγ′
pγ−Θ(pγ) + Θ(pγ′)
]
. (24.75)
The first term (pγ − pγ′)/pγ arises because the incoming and scattered photon energies differ slightly. The
difference in photon energies is given by energy conservation:
pγ − pγ′ = Ee′ − Ee
=
(
me +p′2e2me
)
−(
me +p2
e
2me
)
=(pe + pγ − pγ′)2 − p2
e
2me
=(pγ − pγ′) · (2pe + pγ − pγ′)
2me
≈ (pγ − pγ′) · ve , (24.76)
the last line of which follows because the photon momentum is small compared to the electron momentum,
pγ ∼ T ∼ mev2e = peve ≪ pe . (24.77)
Because the photon energy difference is of first order, and the temperature fluctuation is already of first
order, it suffices to regard the temperature fluctuation Θ as being a function only of the direction pγ of the
photon momentum, not of its energy:
Θ(pγ) ≈ Θ(pγ) . (24.78)
The linear approximations (24.76) and (24.78) bring the difference (24.75) between the initial and final
photon occupancies to
− fγ + fγ′ =∂
0
fγ
∂ lnT[(pγ − pγ′) · ve −Θ(pγ) + Θ(pγ′)] . (24.79)
Inserting this difference in occupancies into the collision integral (24.74) yields
C[1
fγ ] =pγ
16πm2e~2
∂0
fγ
∂ lnT
∫
|M|2fe [(pγ − pγ′) · ve −Θ(pγ) + Θ(pγ′)]2 d3pe
(2π~)3doγ′
4π. (24.80)
24.11 The photon collision term for electron-photon scattering 401
The transition probability |M|2, equation (24.67), is independent of electron momenta, so the integration
over electron momentum in the collision integral (24.80) is straightforward. The unperturbed electron density
ne and the electron bulk velocity vvve are defined by
ne ≡∫
0
fe
2 d3pe
(2π~)3, nevvve ≡
∫
0
fe ve2 d3pe
(2π~)3. (24.81)
Coulomb scattering keeps electrons and ions tightly coupled, so the electron bulk velocity vvve equals the
baryon bulk velocity vvvb,
vvve = vvvb . (24.82)
Integration over the electron momemtum brings the collision integral (24.80) to
C[1
fγ ] =nepγ
16πm2e~2
∂0
fγ
∂ lnT
∫
|M|2 [(pγ − pγ′) · vvvb −Θ(pγ) + Θ(pγ′)]doγ′
4π. (24.83)
Finally, the collision integral (24.83) must be integrated over the direction pγ′ of the scattered photon. In-
serting the electron-photon scattering transition probability |M|2 given by equation (24.67) into the collision
integral (24.83) brings it to
C[1
fγ ] = neσTpγ
∂0
fγ
∂ lnT
∫
[
1 + 12P2(pγ · pγ′)
]
[(pγ − pγ′) · vvvb −Θ(pγ) + Θ(pγ′)]doγ′
4π. (24.84)
The pγ′ · vvvb term in the integrand of (24.84) is odd, and vanishes on angular integration:∫
[
1 + 12P2(pγ · pγ′)
]
pγ′doγ′
4π= 0 . (24.85)
Similarly, the angular integral over the quadrupole of quantities independent of pγ′ vanishes:∫
P2(pγ · pγ′) [pγ · vvvb −Θ(pγ)]doγ′
4π= 0 . (24.86)
The collision integral (24.84) thus reduces to
C[1
fγ(x, pγ)] = neσTpγ
∂0
fγ
∂ lnT
pγ · vvvb(x)−Θ(x, pγ) +
∫
[
1 + 12P2(pγ · pγ′)
]
Θ(x, pγ′)doγ′
4π
, (24.87)
where the dependence of various quantities on comoving position x has been made explicit. Now transform
to Fourier space (in effect, replace comoving position x by comoving wavevector k). Replace the baryon bulk
velocity by its scalar part, vvvb = ikvb. To perform the remaining angular integral over the photon direction
pγ′ , expand the temperature fluctuation Θ(k, pγ′) in scalar multipole moments according to equation (24.48),
and invoke the orthogonality relation (24.134). With µγ defined by
µγ ≡ k · pγ , (24.88)
402 ∗Cosmological perturbations: a more careful treatment of photons and baryons
these manipulations bring the photon collision integral (24.87) at last to
C[1
fγ(k, pγ)] = neσTpγ
∂0
fγ
∂ lnT
[
− iµγvb(k)−Θ(k, µγ) + Θ0(k)− 12Θ2(k)P2(µγ)
]
. (24.89)
24.12 Boltzmann equation for photons
Inserting the collision term (24.89) into equation (24.47) yields the photon Boltzmann equation for scalar
fluctuations in conformal Newtonian gauge,
dΘ
dη=∂Θ
∂η− ikµγΘ− Φ− ikµγΨ = neσTa
[
− iµγvb −Θ + Θ0 − 12Θ2P2(µγ)
]
. (24.90)
Expanded into the scalar harmonics Θℓ(η,k), the photon Boltzmann equation (24.46) yields the hierarchy
of photon multipole equations
Θ0 − kΘ1 − Φ = 0 , (24.91a)
Θ1 +k
3(Θ0 − 2Θ2) +
k
3Ψ =
1
3neσTa (vb − 3Θ1) , (24.91b)
Θ2 +k
5(2Θ1 − 3Θ3) = − 9
10neσTaΘ2 , (24.91c)
Θℓ +k
2ℓ+ 1[ℓΘℓ−1 − (ℓ + 1)Θℓ+1] = −neσTaΘℓ (ℓ ≥ 3) . (24.91d)
The Boltzmann hierarchy (24.91) shows that all the photon multipoles are affected by electron-photon
scattering, but only the photon dipole Θ1 depends directly on one of the baryon variables, the baryon bulk
velocity vb. The dependence on the baryon velocity vb reflects the fact that, to linear order, there is a
transfer of momentum between photons and baryons, but no transfer of number or of energy.
For the dipole, ℓ = 1, electron-photon scattering drives the electron and photon bulk velocities into near
equality,
vb ≈ 3Θ1 , (24.92)
so that the right hand side side of the dipole equation (24.91b) is modest despite the large scattering factor
neσTa. The approximation in which the bulk velocities of electrons and photons are exactly equal, vb = 3Θ1,
and all the higher multipoles vanish, Θℓ = 0 for ℓ ≥ 2, is called the tight-coupling approximation. The tight-
coupling approximation was already invoked in the “simple” model of Chapter 23.
For multipoles ℓ ≥ 2, the electron-photon scattering term on the right hand side of the Boltzmann
hierarchy (24.91) acts as a damping term that tends to drive the multipole exponentially into equilibrium
(the solution to the homogeneous equation Θℓ + neσTaΘℓ = 0 is a decaying exponential). As seen in
Chapter 23, in the tight-coupling approximation the monopole and dipole oscillate with a natural frequency
of ω = csk, where cs is the sound speed. These oscillations provide a source that propagates upward
to higher harmonic number ℓ. For scales much larger than a mean free path, k/(neσTa) ≪ 1, the time
24.13 Diffusive (Silk) damping 403
derivative is small compared to the scattering term, |Θ| ∼ csk|Θ| ≪ neσTa|Θ|, reflecting the near-equilibrium
response of the higher harmonics. For multipoles ℓ ≥ 2, the dominant term on the left hand side of the
Boltzmann hierarchy (24.91) is the lowest order multipole, which acts as a driver. Solution of the Boltzmann
equations (24.91) then requires that
Θℓ+1 ∼k
neσTaΘℓ for ℓ ≥ 2 . (24.93)
The relation (24.93) implies that higher order photon multipoles are successively smaller than lower orders,
|Θℓ+1| ≪ |Θℓ|, for scales much larger than a mean free path, k/(neσTa)≪ 1. This accords with the physical
expectation that electron-photon scattering drives the photon distribution to near isotropy.
24.13 Diffusive (Silk) damping
The tight coupling between photons and baryons is not perfect, because the mean free path for electron-
photon scattering is finite, not zero. The imperfect coupling causes sound waves to dissipate at scales
comparable to the mean free path. To lowest order, the dissipation can be taken into account by including
the photon quadrupole Θ2 in the system (24.91) of photon multipole equations, but still neglecting the higher
multipoles, Θℓ = 0 for ℓ ≥ 3. According to the estimate (24.93), this approximation is valid for scales much
larger than a mean free path, k/(neσTa)≪ 1. The approximation is equivalent to a diffusion approximation.
In the diffusion approximation, the photon quadrupole equation (24.91c) reduces to
Θ2 = − 4k
9neσTaΘ1 . (24.94)
Substituted into the photon momentum equation (24.91b), the photon quadrupole Θ2 (24.94) acts as a source
of friction on the photon dipole Θ1.
24.14 Baryons
The equations governing baryonic matter are similar to those governing non-baryonic cold dark matter,
§24.5, except that baryons are collisional. Coulomb scattering between electrons and ions keep baryons
tightly coupled to each other. Electron-photon scattering couples baryons to the photons.
Since the unperturbed distribution of baryons is in thermodynamic equilibrium, the unperturbed collision
term vanishes for each species of baryonic matter, as it did for photons, equation (24.95),
C[0
fb] = 0 . (24.95)
For the perturbed baryon distribution, only the first and second moments of the phase-space distribution
are important, since these govern the baryon overdensity δb and bulk velocity vvvb. The relevant collision term
404 ∗Cosmological perturbations: a more careful treatment of photons and baryons
is the electron collision term associated with electron-photon scattering. Since electron-photon scattering
neither creates nor destroys electrons, the zeroth moment of the electron collision term vanishes,
∫
C[1
fe]2d3pe
me(2π~)3= 0 . (24.96)
The first moment of the electron collision term is most easily determined from the fact that electron-photon
collisions must conserve the total momentum of electron and photons:
∫
C[1
fe]meve2 d3pe
me(2π~)3+
∫
C[1
fγ ] pγ2 d3pγ
pγ(2π~)3= 0 . (24.97)
Substituting the expression (24.87) for the photon collision integral into equation (24.97), separating out
factors depending on the absolute magnitude pγ and direction pγ of the photon momentum, and taking
into consideration that the integral terms in equation (24.87), when multiplied by pγ , are odd in pγ , and
therefore vanish on integration over directions pγ , gives
∫
C[1
fe]meve2 d3pe
me(2π~)3= neσT
∫
∂0
fγ
∂ lnTp2
γ
2 4πp2γdpγ
pγ(2π~)3
∫
[− pγ · vvvb + Θ(pγ)] pγdoγ
4π. (24.98)
The integral over the magnitude pγ of the photon momentum in equation (24.98) yields 4ρ, in accordance
with equation (24.54). Transformed into Fourier space, and with only scalar terms retained, the collision
integral (24.98) becomes, with µγ ≡ k · pγ ,
k ·∫
C[1
fe]meve2 d3pe
me(2π~)3= 4ργneσT
∫
[iµγvb + Θ] µγdoγ
4π
=4
3iργneσT (vb − 3Θ1) . (24.99)
The result is that the equations governing the baryon overdensity δb and scalar bulk velocity vb look like
those (24.38) governing non-baryonic cold dark matter, except that the velocity equation has an additional
source (24.99) arising from momentum transfer with photons through electron-photon scattering:
δb − kvb − 3Φ = 0 , (24.100a)
vb +a
avb + kΨ = − neσTa
R(vb − 3Θ1) , (24.100b)
where R is 34 the baryon-to-photon density ratio,
R ≡ 3ρb
4ργ= Ra
a
aeq, Ra =
3gρΩb
8Ωm≈ 0.21 , (24.101)
with gρ = 2 + 6 78
(
411
)4/3= 3.36 being the energy-weighted effective number of relativistic particle species
at around the time of recombination (Exercise 10.18). The parameter R plays an important role in that it
modulates the sound speed in the photon-baryon fluid, equation (24.110) below.
24.15 Viscous baryon drag damping 405
24.15 Viscous baryon drag damping
A second source of damping of sound waves, distinct from the diffusive damping of §24.13, arises from the
viscous drag on photons that results from a small difference vb − 3Θ1 between the baryon and photon bulk
velocities. This damping is associated with the finite mass density of baryons, and vanishes in the limit of
small baryon density. However, for realistic values of the baryon density, viscous baryon drag damping is
comparable to diffusive damping.
Equation (24.100b) for the baryon bulk velocity may be written
vb − 3Θ1 = − R
neσTa
(
vb +a
avb + kΨ
)
. (24.102)
The right hand side of this equation is small because the scattering rate neσT is large, so to lowest order
vb = 3Θ1, the tight-coupling approximation. To ascertain the effect of a small difference in the baryon
and photon bulk velocities, take equation (24.102) to next order. In the circumstances where damping is
important, which is small scales well inside the sound horizon, kηs ≫ 1, the dominant term on the right hand
side of equation (24.100b) is the time derivative vb. With only this term kept, equation (24.102) reduces to
vb − 3Θ1 ≈ −R
neσTavb ≈ −
3R
neσTaΘ1 . (24.103)
Substituting the approximation (24.103) into the left hand side of the baryon velocity equation (24.100b)
gives
vb +a
avb + kΨ = 3
(
Θ1 +a
aΘ1 +
k
3Ψ
)
− 3R
neσTaΘ1 , (24.104)
in which the last term on the right hand side is the small correction from the non-vanishing baryon-photon
velocity difference, equation (24.103). As in the approximation (24.103), only the dominant term, the one
arising from the time derivative vb, has been retained in the correction, and in differentiating the right
hand side of equation (24.103), the time derivative of 3R/(neσTa) has been neglected compared to the time
derivative of Θ1. The final simplification is to replace the second time derivative of the photon dipole in
the correction term by its unperturbed value, Θ1 ≈ −c2sk2Θ1, an approximation that is valid because the
correction term is already small. Here cs is the sound speed, found in the next section to be given by
equation (24.110). The result is
vb +a
avb + kΨ = 3
(
Θ1 +a
aΘ1 +
k
3Ψ
)
+3Rc2sk
2
neσTaΘ1 . (24.105)
This equation (24.105) is used to develop the photon-baryon wave equation in the next section, §24.16.
24.16 Photon-baryon wave equation
Combining the photon monopole and dipole equations (24.91a) and (24.91b) with the baryon momentum
equation (24.100b), and making the diffusion approximation (24.94) for the photon quadrupole, and the
406 ∗Cosmological perturbations: a more careful treatment of photons and baryons
approximation (24.105) for the baryon bulk velocity, yields a second order differential equation (24.113) that
captures accurately the behaviour of the photon distribution. The equation is that for a damped harmonic
oscillator, forced by the gravitational potential terms on its right hand side.
Adding the momentum equations (24.91b) and (24.100b) for photons and baryons yields an equation that
expresses momentum conservation for the combined photon-baryon fluid,
Θ1 +k
3(Θ0 − 2Θ2) +
k
3Ψ +
R
3
(
vb +a
avb + kΨ
)
= 0 . (24.106)
Setting the photon quadrupole Θ2 equal to its diffusive value (24.94), and substituting the baryon velocity
equation (24.105), brings the photon-baryon momentum conservation equation (24.106) to
Θ1 +k
3Θ0 +
8k2
27neσTaΘ1 +
k
3Ψ +R
(
Θ1 +a
aΘ1 +
k
3Ψ
)
+R2c2sk
2
neσTaΘ1 = 0 . (24.107)
Equation (24.107) rearranges to[
d
dη+
R
1 +R
a
a+
k2
3neσTa(1 +R)
(
8
9+
R2
1 +R
)]
Θ1 +k
3(1 +R)Θ0 +
k
3Ψ = 0 , (24.108)
where equation (24.110) has been used to replace the sound speed cs. Finally, eliminating the dipole Θ1 in
favour of the monopole Θ0 using the photon monopole equation (24.91a) yields a second order differential
equation for Θ0 − Φ:
d2
dη2+
[
R
1 +R
a
a+
k2
3neσTa(1 +R)
(
8
9+
R2
1 +R
)]
d
dη+
k2
3(1 +R)
(Θ0−Φ) = − k2
3(1 +R)[(1 +R)Ψ + Φ] .
(24.109)
Equation (24.109) is a wave equation for a damped, driven oscillator with sound speed
cs =
√
1
3(1 +R), (24.110)
which is the adiabatic sound speed cs for a fluid in which photons provide all the pressure, but both
photons and baryons contribute to the mass density. The term proportional to a/a on the left hand side of
equation (24.109) can be expressed as, since R ∝ a,R
1 +R
a
a= − 2
cscs
. (24.111)
Define the conformal sound time ηs by
dηs ≡ csdη , (24.112)
with respect to which sound waves move at unit velocity, unit comoving distance per unit conformal time,
dx/dηs = 1. Recast in terms of the conformal sound time ηs, the wave equation (24.109) becomes
d2
dη2s
+
[
− c′s
cs+
k2csneσTa
(
8
9+
R2
1 +R
)]
d
dηs+ k2
(Θ0 − Φ) = − k2 [(1 +R)Ψ + Φ] , (24.113)
24.17 Damping of photon-baryon sound waves 407
where prime ′ denotes derivative with respect to conformal sound time, c′s = dcs/dηs.
The “simple” photon wave equation (23.45) derived in Chapter 23 is obtained from the wave equa-
tion (24.113) in the limit of negligible baryon-to-photon density, R ≈ 0.
24.17 Damping of photon-baryon sound waves
The terms proportional to the linear derivative d/dηs in the wave equation (24.113) are damping terms,
the first being an adiabatic damping term associated with variation of the sound speed, and the others
being dissipative damping terms associated respectively with the finite mean free path of electron-photon
scattering, and with viscous baryon drag. Lump these terms into a damping parameter κ defined by
κ ≡ 1
2
[
− c′s
cs+
k2csneσTa
(
8
9+
R2
1 +R
)]
. (24.114)
The damping parameter κ varies slowly compared to the frequency of the sound wave, so κ can be treated
as approximately constant. The homogeneous wave equation (equation (24.113) with zero on the right hand
side) can then be solved by introducing a frequency ω defined by
Θ0 − Φ ∝ eR
ω dηs . (24.115)
The homogeneous wave equation (24.113) is then equivalent to
ω′ + ω2 + 2κω + k2 = 0 . (24.116)
Since the damping parameter κ is approximately constant (and the comoving wavevector k is by definition
constant), ω′ is small compared to the other terms in equation (24.116). With ω′ set to zero, the solution of
equation (24.116) is
ω = − κ ± i√
k2 − κ2 ≈ −κ ± ik , (24.117)
where the last approximation is valid since the damping rate is small compared to the frequency, κ ≪ k.
Thus the homogeneous solutions of the wave equation (24.113) are
Θ0 − Φ ∝ e−R
κ dηs ± ikηs . (24.118)
In the present case, the first of the sources (24.114) of damping is the adiabatic damping term
κa ≡ −1
2
c′scs
. (24.119)
The integral of the adiabatic damping term is∫
κa dηs = − 12 ln cs, whose exponential is
e−R
κa dηs =√cs . (24.120)
This shows that, as the sound speed decreases thanks to the increasing baryon-to-photon density in the
expanding Universe, the amplitude of a sound wave decreases as the square root of the sound speed.
408 ∗Cosmological perturbations: a more careful treatment of photons and baryons
The remaining damping terms are the dissipative terms
κd ≡k2cs
2neσTa
(
8
9+
R2
1 +R
)
. (24.121)
The integral of the dissipative damping terms is∫
κd dηs =k2
k2d
, (24.122)
where kd is the damping scale defined by
1
k2d
≡∫
cs2neσTa
(
8
9+
R2
1 +R
)
dηs =
∫
1
6neσTa(1 +R)
(
8
9+
R2
1 +R
)
dη . (24.123)
The resulting damping factor is
e−R
κd dηs = e−k2/k2d . (24.124)
Thus the effect of dissipation is to damp temperature fluctuations exponentially at scales smaller than the
diffusion scale kd.
With adiabatic and diffusion damping included, the homogeneous solutions to the wave equation (24.113)
are approximately
Θ0 − Φ ∝ √cs e−k2/k2d e±ikηs . (24.125)
The driving potential on the right hand side of the wave equation (24.113) causes Θ0 − Φ to oscillate not
around zero, but rather around the offset − [(1 +R)Ψ + Φ]. At the high frequencies where damping is
important, this driving potential also varies slowly compared to the wave frequency. To the extent that the
driving potential is slowly varying, the complete solution of the inhomogeneous wave equation (24.113) is
Θ0 + (1 +R)Ψ ∝ √cs e−k2/k2d e±ikηs . (24.126)
As will be seen in Chapter 25, the monopole contribution to CMB fluctuations is not the photon monopole
Θ0 by itself, but rather Θ0+Ψ, which is the monople redshifted by the potential Ψ. This redshifted monopole
is
Θ0 + Ψ = −RΨ +A√cs e
−k2/k2d e±ikηs . (24.127)
Exercise 24.4 Diffusion scale. Show that the damping scale kd defined by (24.123) is given by, with a
normalized to aeq = 1,
H2eq
k2d
=8√
2πGmp
9cσT(1−Xn)Heq
Ωm
Ωb
∫ a
0
a2
Xe
√1 + a(1 +R)
(
8
9+
R2
1 +R
)
da , (24.128)
the Hubble parameter Heq at matter-radiation equality being related to the present-day Hubble parameter
H0 by equation (23.41). If baryons are taken to be fully ionized, Xe = 1, which ceases to be a good
24.18 Ionization and recombination 409
approximation near recombination, then the integral can be done analytically. At times well after matter-
radiation equality,
f(a) ≡∫ a
0
a2
√1 + a(1 +R)
(
8
9+
R2
1 +R
)
da → 2
5a5/2 for a≫ 1 , (24.129)
independent of the value of the constant Ra ≡ R/a. Conclude that, neglecting the effect of recombination
on the electron fraction Xe,
H2eq
k2d
=6.821 h−1
1−Xn
H0
Heq
Ωm
Ωbf(a/aeq) = 4.9× 10−4f(a/aeq) ≈ 0.0016
(
a
a∗
)5/2
, (24.130)
where a∗ is the cosmic scale factor at recombination.
24.18 Ionization and recombination
24.19 Neutrinos
Before electron-positron annihilation at temperature T ≈ 1 MeV, weak interactions were fast enough that
scattering between neutrinos, antineutrinos, electrons, and positrons kept neutrinos and antineutrinos in
thermodynamic equilibrium with baryons. After ee annihilation, neutrinos and antineutrinos decoupled,
rather like photons decoupled at recombination. After decoupling, neutrinos streamed freely.
24.20 Summary of equations
Non-baryonic cold dark matter, baryons, photons, neutrinos:
δc − kvc = 3 Φ , (24.131a)
vc +a
avc = −kΨ , (24.131b)
δb − kvb = 3Φ , (24.131c)
vb +a
avb = − kΨ− neσTa
R(vb − 3Θ1) , (24.131d)
Θ0 − kΘ1 = Φ , (24.131e)
Θ1 +k
3(Θ0 − 2Θ2) = − k
3Ψ +
1
3neσTa (vb − 3Θ1) , (24.131f)
2k
5Θ1 = − 9
10neσTaΘ2 , (24.131g)
410 ∗Cosmological perturbations: a more careful treatment of photons and baryons
Einstein energy and quadrupole pressure equations:
− k2Φ− 3a
aF = 4πGa2(ρcδc + ρbδb + 4ργΘ0 + 4ρνN0) , (24.132a)
k2(Ψ− Φ) = − 32πGa2(ρrΘ2 + ρνN2) , (24.132b)
24.21 Legendre polynomials
The Legendre polynomials Pℓ(µ) satisfy the orthogonality relations∫ 1
−1
Pℓ(µ)Pℓ′(µ) dµ =2
2ℓ+ 1δℓℓ′ (24.133)
and∫
Pℓ(z · a)Pℓ′(z · b) doz =4π
2ℓ+ 1Pℓ(a · b) δℓℓ′ , (24.134)
the recurrence relation
µPℓ(µ) =1
2ℓ+ 1[ℓPℓ−1(µ) + (ℓ+ 1)Pℓ+1(µ)] , (24.135)
and the derivative relationdPℓ(µ)
dµ=
ℓ+ 1
1− µ2[µPℓ−1(µ)− Pℓ+1(µ)] . (24.136)
The first few Legendre polynomials are
P0(µ) = 1 , P1(µ) = µ , P2(µ) = − 1
2+
3
2µ2 . (24.137)
25
Fluctuations in the Cosmic MicrowaveBackground
25.1 Primordial power spectrum
Inflation generically predicts gaussian initial fluctuations with a scale-free power spectrum, in which the
variance of the potential is the same on all scales,
〈Φ(x′)Φ(x)〉 ≡ ξΦ(|x′ − x|) = constant , (25.1)
independent of separation |x′ − x|. A scale-free primordial power spectrum was originally proposed as
a natural initial condition by Harrison and Zel’dovich (Harrison, 1970, PRD 1, 2726; Zel’dovich, 1972,
MNRAS, 160, 1P), before the idea of inflation was conceived. During inflation, vacuum fluctuations generate
fluctuations in the potential which become frozen in as they fly over the horizon. The amplitude of these
fluctuations remains constant as inflation continues, producing a scale-free power spectrum.
The power spectrum PΦ(k) of potential fluctuations is defined by
〈Φ(k′)Φ(k)〉 ≡ (2π)3δD(k′ + k)PΦ(k) , (25.2)
the power spectrum PΦ(k) being related to the correlation function ξΦ(x) by (with the standard convention
in cosmology for the choice of signs and factors of 2π)
PΦ(k) =
∫
eik·xξΦ(x) d3x , ξΦ(x) =
∫
e−ik·xPΦ(k)d3k
(2π)3. (25.3)
The scale-free character means that the dimensionless power spectrum ∆2Φ(k) defined by
∆2Φ(k) ≡ PΦ(k)
4πk3
(2π)3(25.4)
is constant.
Actually, the power spectrum generated by inflation is not precisely scale-free, because inflation comes to
an end, which breaks scale-invariance. The departure from scale-invariance is conventionally characterized
by a scalar spectral index, the tilt n, such that
∆2Φ(k) ∝ kn−1 . (25.5)
412 Fluctuations in the Cosmic Microwave Background
Thus a scale-invariant power spectrum has
n = 1 (scale-invariant) . (25.6)
Different inflationary models predict different tilts, mostly close to but slightly less than 1.
25.2 Normalization of the power spectrum
It is convenient to normalize the power spectrum of potential fluctuations to its amplitude at the recombi-
nation distance today, k(η0 − η∗) = 1, and not to the initial potential Φ(0) (which vanishes for isocurvature
initial conditions), but rather to the late-time matter-dominated potential Φ(late) at superhorizon scales,
∆2Φ(late)(k) = A2
late [k(η0 − η∗)]n−1 . (25.7)
The relation between the superhorizon late-time matter-dominated potential Φ(late) and the primordial
potential Φ(0) is, equations (23.61),
Φ(late) =
910Φ(0) adiabatic ,
85Φ′(0) isocurvature .
(25.8)
25.3 CMB power spectrum
The power spectrum of temperature fluctuations in the CMB is defined by
Cℓ(η0) ≡ 4π
∫ ∞
0
|Θℓ(η0,k)|2 d3k
(2π)3. (25.9)
As seen in Chapter §23, during linear evolution, scalar modes of given comoving wavevector k evolve with
amplitude proportional to the initial value Φ(0,k) of the scalar potential (or of its derivative Φ′(0,k), for
isocurvature initial conditions). The evolution of the amplitude may be encapsulated in a transfer function
Tℓ(η, k) defined by
Tℓ(η, k) ≡Θℓ(η,k)
Φ(late,k), (25.10)
where Φ(late) is the superhorizon late-time matter-dominated potential. By isotropy, the transfer function
Tℓ(η, k) is a function only of the magnitude k of the wavevector k. The power spectrum of the CMB observed
on the sky today is related to the primordial power spectrum PΦ(late)(k) or ∆2Φ(late)(k) by
Cℓ(η0) = 4π
∫ ∞
0
|Tℓ(η0, k)|2 PΦ(late)(k)4πk2dk
(2π)3= 4π
∫ ∞
0
|Tℓ(η0, k)|2 ∆2Φ(late)(k)
dk
k. (25.11)
25.4 Matter power spectrum 413
25.4 Matter power spectrum
The matter power spectrum Pm(k) is defined by
〈δm(k′)δm(k)〉 ≡ (2π)3δD(k′ + k)Pm(k) . (25.12)
At times well after recombination, the matter power spectrum Pm(η, k) at conformal time η is related to the
potential power spectrum (25.2) by, equation (23.117), in units a0 = 1,
Pm(η, k) =
(
2a
3ΩmH20
)2
k4 PΦ(η, k) =
(
2a
3ΩmH20
)2(2π)3
4πk∆2
Φ(k, η) . (25.13)
At superhorizon scales, the potential Φ(η, k) at conformal time η is related to the late-time matter-dominated
potential by
Φ(η) = g(a)Φ(late) for kη ≪ 1 , (25.14)
where g(a) is the grown factor defined by equation (23.116). For a power-law primordial spectrum (25.5),
the matter power spectrum at the largest scales goes as
Pm(η, k) ∝ kn , (25.15)
which explains the origin of the scalar index n.
25.5 Radiative transfer of CMB photons
To determine the harmonics Θℓ(η0,k) of the CMB photon distribution at the present time, return to the
Boltzmann equation (24.90) for the photon distribution Θ(η,k, µ), where µ ≡ k · p:
Θ− ikµγΘ + neσTaΘ = Φ + ikµγΨ + neσTa[
− iµγvb + Θ0 − 12Θ2P2(µγ)
]
. (25.16)
This equation is also called the radiation transfer equation. The terms on the right are the source terms.
Define the electron-photon (Thomson) scattering optical depth τ by
dτ
dη≡ −neσTa , (25.17)
starting from zero, τ0 = 0, at zero redshift, and increasing going backwards in time η to higher redshift. The
photon Boltzmann equation (25.16) can be written
eikµγη+τ d
dη
(
e−ikµγη−τ Θ)
= S0 − iµγS1 + (−iµγ)2S2 , (25.18)
where Si are source terms
S0 ≡ Φ− τ(
Θ0 + 14Θ2
)
, (25.19a)
S1 ≡ −kΨ− τvb = −kΨ− 3τΘ1 , (25.19b)
S2 ≡ −τ 34Θ2 , (25.19c)
414 Fluctuations in the Cosmic Microwave Background
where in S1 the tight-coupling approximation vb = 3Θ1 has been used to replace the baryon bulk velocity
vb with the photon dipole Θ1. Thus a solution for the photon distribution Θ(η0,k, µγ) today is, at least
formally, an integral over the line of sight from the Big Bang to the present time,
Θ(η0,k, µγ) =
∫ η0
0
[
S0 − iµγS1 + (−iµγ)2S2
]
e−ikµγ (η−η0)−τ dη . (25.20)
The −iµγ dependence of the source terms inside the integral can be accomodated through
(−iµγ)ne−ikµγ (η−η0) =
(
1
k
∂
∂η
)n
e−ikµγ (η−η0) , (25.21)
which brings the formal solution (25.20) for Θ(η0,k, µγ) to
Θ(η0,k, µγ) =
∫ η0
0
e−τ
(
S0 + S11
k
∂
∂η+ S2
1
k2
∂2
∂η2
)
e−ikµγ (η−η0) dη . (25.22)
The dipole source term S1 contains a −Ψ term which it is helpful to integrate by parts:∫ η0
0
−e−τΨ∂
∂ηe−ikµγ (η−η0) dη =
[
−e−τΨe−ikµγ(η−η0)]η0
0+
∫ η0
0
∂
∂η
(
e−τΨ)
e−ikµγ (η−η0) dη
= −Ψ(η0) +
∫ η0
0
e−τ(
Ψ− τΨ)
e−ikµγ (η−η0) dη . (25.23)
With the source terms written out explicitly, the present-day photon distribution Θ(η0,k, µγ), equation (25.20),
is
Θ(η0,k, µγ) + Ψ(η0,k)
=
∫ η0
0
[
e−τ(
Ψ + Φ)
− τ e−τ
(
Θ0 + Ψ +1
4Θ2 + 3Θ1
1
k
∂
∂η+
3
4Θ2
1
k2
∂2
∂η2
)]
e−ikµγ (η−η0) dη . (25.24)
The spherical harmonics Θℓ(η,k) of the photon distribution have been defined by equation (24.48). The
e−ikµγ (η−η0) factor in the integral in equation (25.24) can be expanded in spherical harmonics through the
general formula
eik·x =∞∑
ℓ=0
iℓ(2ℓ+ 1)jℓ(kx)Pℓ(k · x) , (25.25)
where jℓ(z) ≡√
π/(2z)Jℓ+1/2(z) are spherical Bessel functions. Resolved into harmonics, equation (25.24)
becomes
[Θℓ(η0,k) + δℓ0Ψ(η0,k)]
=
∫ η0
0
[
e−τ(
Ψ + Φ)
− τ e−τ
(
Θ0 + Ψ +1
4Θ2 + 3Θ1
1
k
∂
∂η+
3
4Θ2
1
k2
∂2
∂η2
)]
jℓ [k(η0 − η)] dη . (25.26)
Introduce a visibility function g(η) defined by
g(η) ≡ −τ e−τ , (25.27)
25.6 Integrals over spherical Bessel functions 415
whose integral is one,∫ η0
0
g(η) dη =
∫ 0
∞−e−τ dτ =
[
e−τ]0
∞ = 1 . (25.28)
The visibility function is fairly narrowly peaked around recombination at η = η∗. In this approximation of
instantaneous recombination,
[Θℓ(η0,k) + δℓ0Ψ(η0,k)] =
∫ η0
0
e−τ(
Ψ(η,k) + Φ(η,k))
jℓ [k(η0 − η)] dη ISW
+ [Θ0(η∗,k) + Ψ(η∗,k)] jℓ [k(η0 − η∗)] monopole
− 3 Θ1(η∗,k) j′ℓ [k(η0 − η∗)] dipole
+ Θ2(η∗,k)
14jℓ [k(η0 − η∗)] + 3
4j′′ℓ [k(η0 − η∗)]
quadrupole .
(25.29)
Here prime ′ on jℓ denotes a total derivative, j′ℓ(z) = djℓ(z)/dz. The first and second derivatives of the
spherical Bessel functions are
j′ℓ(z) =ℓ
zjℓ(z)− jℓ+1(z) , j′′ℓ (z) =
ℓ(ℓ− 1)− z2
z2jℓ(z) +
2
zjℓ+1(z) . (25.30)
Notice that the monopole term (on both sides of equation (25.29)) is not Θ0 but rather Θ0 + Ψ, which
is the temperature fluctuation redshifted by the potential Ψ. On the left hand side, the δℓ0 term arises
from the redshift at our position today, but this monopole perturbation just adds to the mean unperturbed
termperature, and is not observable.
25.6 Integrals over spherical Bessel functions
Computing the photon harmonics Θl by integration of equation (25.26), or the CMB power spectrum Cℓ
in the instantaneous recombination approximation by integration of equation (25.11) with (25.29), involves
evaluating integrals of the form∫ ∞
0
f(z)g(qz)dz
z, (25.31)
with
g(z) ≡ jℓ(z)zn−1 or g(z) ≡ jℓ(z)jℓ′(z)zn−1 . (25.32)
Such integrals present a challenge because of the oscillatory character of the functions g(z). This section
presents a method to evaluate such integrals reliably. Integrals with different ℓ are related by recursion
relations (25.38) that permit rapid evaluation over many ℓ. Details are left to Exercise 25.2.
The approach is to recast the integral (25.31) into Fourier space with respect to ln z, and to apply a
Fast Fourier Transform of f(z) over a logarithmic interval. This involves replacing the true f(z) with a
function that is periodic in ln z over a logarithmic interval [−L/2, L/2] of width L centred on ln z = 0. The
416 Fluctuations in the Cosmic Microwave Background
approximation works because the functions g(z) given by equation (25.32) tend to zero at z → 0 and z →∞,
so spurious periodic duplications at small and large z contribute negligibly to the integral (25.31) provided
that the logarithmic interval L is chosen sufficiently broad. The logarithmic interval L can be broadened to
whatever extent is necessary by extrapolating the function f(z) to smaller and larger z. If f(z) is periodic
in ln z over a logarithmic interval L, then f(z) is a sum of discrete Fourier modes e2πim ln(z)/L in which m is
integral. If f(z) is a smooth function, then f(z) may be adequately approximated by a finite number N of
discrete modes m = −[(N − 1)/2], ..., [N/2], where [N/2] denotes the largest integer greater than or equal to
N/2. Under these circumstances, the function f(z) is given by a discrete Fourier expansion whose Fourier
components fm are related to the values f(zn) at N logarithmically spaced points zn = enL/N ,
f(z) =
[N/2]∑
m=−[(N−1)/2]
fme2πim ln(z)/L , fm =
1
N
[N/2]∑
n=−[(N−1)/2]
f(zn)e−2πimn/N . (25.33)
Since f(z) is real, the negative frequency modes are the complex conjugates of the positive frequency modes,
f−m = f∗m , (25.34)
and for even N (the usual choice) the outermost (Nyquist) frequency mode fN/2 is real. For a function f(z)
given by the Fourier sum (25.33), the integral (25.31) is
∫ ∞
0
f(z)g(qz)dz
z=
[N/2]∑
m=−[(N−1)/2]
fmq−2πim/L
∫ ∞
0
g(z)z2πim/L dz
z. (25.35)
For functions g(z) given by equation (25.32), the integrals over g(z) on the right hand side of equation (25.35)
can be done analytically:
∫ ∞
0
f(z)g(qz)dz
z=
[N/2]∑
m=−[(N−1)/2]
fmq−2πim/LU
(
n− 1 +2πim
L
)
, (25.36)
where for g(z) = jℓ(z)zn−1 or g(z) = jℓ(z)jℓ′(z)z
n−1 the function U(x) is respectively
Uℓ(x) ≡∫ ∞
0
jℓ(z)zx dz
z=
2x−2√πΓ[
12 (ℓ+ x)
]
Γ[
12 (ℓ− x+ 3)
] , (25.37a)
Uℓℓ′(x) ≡∫ ∞
0
jℓ(z)jℓ′(z)zx dz
z=
2x−3πΓ(2 − x)Γ[
12 (ℓ+ ℓ′ + x)
]
Γ[
12 (ℓ + ℓ′ − x+ 4)
]
Γ[
12 (ℓ − ℓ′ − x+ 3)
]
Γ[
12 (ℓ′ − ℓ− x+ 3)
] . (25.37b)
25.7 Large-scale CMB fluctuations 417
Recurrence relations such as
Uℓ(x) = (ℓ + x− 2)Uℓ−1(x− 1) (25.38a)
=ℓ+ x− 2
ℓ− x+ 1Uℓ−2(x) , (25.38b)
Uℓℓ′(x) =ℓ+ ℓ′ + x− 2
ℓ+ ℓ′ − x+ 2Uℓ−1,ℓ′−1(x) (25.38c)
=(ℓ+ ℓ′ + x− 2)(ℓ− ℓ′ + x− 3)
2(x− 2)Uℓ−1,ℓ′(x − 1) (25.38d)
=(ℓ+ ℓ′ + x− 2)(ℓ− ℓ′ + x− 3)
(ℓ+ ℓ′ − x+ 2)(ℓ′ − ℓ+ x− 1)Uℓ−2,ℓ′(x) , (25.38e)
permit rapid evaluation of the functions Uℓ(x) or Uℓℓ′(x) as a function of ℓ and ℓ′, starting from small ℓ, ℓ′.As written, the right hand side of equation (25.36) has a small imaginary part arising from the contribution
of the outermost (Nyquist) mode, m = [N/2], but this imaginary part should be dropped since it cancels
when averaged with the contribution of its negative frequency partner m = −[N/2].
25.7 Large-scale CMB fluctuations
The behaviour of the CMB power spectrum at the largest angular scales was first predicted by R. K. Sachs
& A. M. Wolfe (1967, Astrophys. J., 147, 73), and is therefore called the “Sachs-Wolfe effect,” though why
it should be called an effect is mysterious. The Sachs-Wolfe (SW) effect is distinct from, but modulated by,
the Integrated Sachs-Wolfe (ISW) effect. The ISW effect, ignored in this section, is considered in §25.9.
At scales much larger than the sound horizon at recombination, kηs,∗ ≪ 1, the redshifted monopole
fluctuation Θ0(η∗,k) + Ψ(η∗,k) at recombination is much larger than the dipole Θ1(η∗,k) or quadrupole
Θ2(η∗,k), so only the monopole contributes materially to the temperature multipoles Θℓ(η0,k) today. The
redshifted monopole contribution to the temperature multipoles Θℓ(η0,k) today is, from equation (25.29),
Θℓ(η0,k) = [Θ0(η∗,k) + Ψ(η∗,k)] jℓ [k(η0 − η∗)] . (25.39)
At the very largest scales, kηeq ≪ 1, the solution for the redshifted radiation monopole Θ0 + Ψ at the time
η∗ of recombination is, from equation (23.62),
Θ0(η∗,k) + Ψ(η∗,k) = 2Φsuper(η∗,k)− 3
2Φ(0)
≡ ASW(η∗)
AlateΦsuper(late,k) , (25.40)
where the last expression defines the Sachs-Wolfe amplitude ASW(η∗) at recombination. In the approximation
that recombination happens well into the matter-dominated regime, so that Φsuper(η∗,k) ≈ Φsuper(late,k),
418 Fluctuations in the Cosmic Microwave Background
from equations (23.61),
ASW(η∗)
Alate≈ ASW(late)
Alate= 2− 3Φ(0)
2Φ(late)=
13 adiabatic ,
2 isocurvature .(25.41)
In reality, recombination occurs only somewhat into the matter-dominated regime, and the solutions for
the potential Φsuper(η∗,k) from §23.9 should be used in place of the approximation (25.41). Putting equa-
tions (25.39) and (25.40) together shows that the transfer function Tℓ(η0,k), equation (25.10) that goes into
the present-day CMB angular power spectrum Cℓ(η0), equation (25.11), is
Tℓ(η0,k) =ASW(η∗)
Alatejℓ [k(η0 − η∗)] . (25.42)
If the primordial power spectrum is a power law with tilt n, equation (25.7), then the resulting CMB angular
power spectrum is, with z = k(η0 − η∗),
Cℓ(η0) = 4πASW(η∗)2
∫ ∞
0
jℓ(z)2zn−1dz
z= 4πASW(η∗)
2Uℓ,ℓ(n− 1) , (25.43)
where Uℓ,ℓ(x) is given by equation (25.37b). For the particular case of a scale-invariant primordial power
spectrum, n = 1, the CMB power spectrum Cℓ at large scales today is given by
ℓ(ℓ+ 1)Cℓ(η0) = 2πASW(η∗)2 if n = 1 . (25.44)
Thus the characteristic feature of a scale-invariant primordial power spectrum, n = 1, is that ℓ(ℓ + 1)Cℓ
should be approximately constant at the largest angular scales, ℓ ≪ η0/η∗. This is a primary reason why
CMB folk routinely plot ℓ(ℓ+ 1)Cℓ, rather than Cℓ.
25.8 Monopole, dipole, and quadrupole contributions to Cℓ
At smaller scales, kη∗ >∼ 1, not only the photon monopole Θ0(η∗, k), but also the dipole Θ1(η∗, k), and to
a small extent the quadrupole Θ2(η∗, k), contribute to the temperature multipoles Θℓ(η0, k) today, equa-
tion (25.29). The dipole is related to the monopole by the evolution equation (??) for the monopole.
kΘ1 = Θ0 − Φ . (25.45)
25.9 Integrated Sachs-Wolfe (ISW) effect
Concept question 25.1 Cosmic Neutrino Background. Just as photons decoupled at recombina-
tion, so also neutrinos decoupled at electron-positron annihilation. Compare qualitatively the expected
fluctuations in the CνB to those in the CMB.
25.9 Integrated Sachs-Wolfe (ISW) effect 419
Exercise 25.2 Numerical integration of sequences of integrals over Bessel functions. Write code
that solves integrals (25.31) numerically for g(z) given by equation (25.32), using a Fast Fourier Transform,
equation (25.36), amd recurrence relations appropriate for the monopole, dipole, and quadrupole contribu-
tions to equations (25.26) or (25.29). To compute Uℓ(x) or Uℓℓ′(x) for the initial ℓ, ℓ′, you will need to find
code that implements the complex Gamma function. Note that most FFT codes store input and output
periodic sequences shifted by [N/2] compared to the convention (25.33). That is, an FFT code typically
takes an input sequence ordered as f(z0), f(z1), ..., f(z[N/2]), ..., f(z−2), f(z−1), with periodic identification
f(zn) = f(zn+N), and evaluates Fourier coefficients fm as
fm =1
N
N−1∑
m=0
f(zn)e−2πimn/N , (25.46)
which yields the same Fourier components fm as (25.33) but in the order f0, f1, ..., f[N/2], ..., f−2, f−1, with
periodic identification fm = fm+N .
Various recurrence relations. For g(z) = jℓ(z)zn−1, the dipole and quadrupole integrals U
(1)ℓ (x) and
U(2)ℓ (x) are related to the monopole integral Uℓ(x) (25.37a) by
U(1)ℓ (x) ≡
∫ ∞
0
j′ℓ(z)zx dz
z= (1− x)Uℓ(x− 1) , (25.47a)
U(2)ℓ (x) ≡
∫ ∞
0
[
14jℓ(z) + 3
4j′′ℓ (z)
]
zx dz
z=
ℓ(ℓ+ 1) + 2x(x− 2)
(ℓ+ x− 2)(ℓ− x+ 3)Uℓ(x) . (25.47b)
The dipole, and quadrupole integrals satisfy the recurrence relations
U(1)ℓ (x) =
ℓ+ x− 3
ℓ− x+ 2U
(1)ℓ−2(x) , (25.48a)
U(2)ℓ (x) =
(ℓ + x− 4) [l(l + 1) + 2x(x− 2)]
(ℓ − x+ 3) [l(l − 3) + 2(x− 1)2]U
(2)ℓ−2(x) . (25.48b)
For g(z) = jℓ(z)jℓ′(z)zn−1, the relations get a bit ugly — it’s a good idea to use Mathematica or a similar
program to generate the relations automatically. The various multipole integrals of interest are related to
420 Fluctuations in the Cosmic Microwave Background
the integral Uℓℓ(x) (25.37b) by
U(0,1)ℓℓ (x) ≡
∫ ∞
0
jℓ(z)j′ℓ(z)z
x dz
z=
1− x2
Uℓℓ(x− 1) , (25.49a)
U(1,1)ℓℓ (x) ≡
∫ ∞
0
[j′ℓ(z)]2zx dz
z=
4ℓ(ℓ+ 1)− x(x− 2)(x− 3)
(3 − x)(2ℓ+ x− 2)(2ℓ− x+ 4)Uℓℓ(x) , (25.49b)
U(0,2)ℓℓ (x) ≡
∫ ∞
0
jℓ(z)[
14jℓ(z) + 3
4j′′ℓ (z)
]
zx dz
z=
x [2ℓ(ℓ+ 1) + (x− 1)(x− 2)]
2(x− 3)(2ℓ+ x− 2)(2ℓ− x+ 4)Uℓℓ(x) , (25.49c)
U(1,2)ℓℓ (x) ≡
∫ ∞
0
j′ℓ(z)[
14jℓ(z) + 3
4j′′ℓ (z)
]
zx dz
z(25.49d)
=(1− x) [2ℓ(ℓ+ 1)(x− 7) + (x+ 1)(x− 3)(x− 4)]
4(x− 4)(2ℓ+ x− 3)(2ℓ− x+ 5)Uℓℓ(x− 1) ,
U(2,2)ℓℓ (x) ≡
∫ ∞
0
[
14jℓ(z) + 3
4j′′ℓ (z)
]2zx dz
z= (25.49e)
4(ℓ− 1)ℓ(ℓ+ 1)(ℓ+ 2) [12 + x(x − 2)] + x(x − 2)(x− 5) [4ℓ(ℓ+ 1)(x− 6) + (x+ 2)(x− 3)(x− 4)]
4(x− 3)(x− 5)(2ℓ+ x− 2)(2ℓ+ x− 4)(2ℓ− x+ 4)(2ℓ− x+ 6)Uℓℓ(x) .
The various recurrence relations of interest are
U(0,1)ℓℓ (x) =
(2ℓ+ x− 3)
(2ℓ− x+ 3)U
(0,1)ℓ−1,ℓ−1(x) , (25.50a)
U(1,1)ℓℓ (x) =
(2ℓ+ x− 4) [4ℓ(ℓ+ 1)− x(x − 2)(x− 3)]
(2ℓ− x+ 4) [4ℓ(ℓ− 1)− x(x − 2)(x− 3)]U
(1,1)ℓ−1,ℓ−1(x) , (25.50b)
U(0,2)ℓℓ (x) =
(2ℓ+ x− 4) [2ℓ(ℓ+ 1) + (x− 1)(x− 2)]
(2ℓ− x+ 4) [2ℓ(ℓ− 1) + (x− 1)(x− 2)]U
(0,2)ℓ−1,ℓ−1(x) , (25.50c)
U(1,2)ℓℓ (x) =
(2ℓ+ x− 5) [2ℓ(ℓ+ 1)(x− 7) + (x+ 1)(x− 3)(x− 4)]
(2ℓ− x+ 5) [2ℓ(ℓ− 1)(x− 7) + (x+ 1)(x− 3)(x− 4)]U
(1,2)ℓ−1,ℓ−1(x) , (25.50d)
U(2,2)ℓℓ (x) =
(2ℓ+ x− 6)
(2ℓ− x+ 6)U
(2,2)ℓ−1,ℓ−1(x) (25.50e)
× 4(ℓ− 1)ℓ(ℓ+ 1)(ℓ + 2) [12 + x(x− 2)] + x(x− 2)(x− 5) [4ℓ(ℓ+ 1)(x− 6) + (x+ 2)(x− 3)(x− 4)]
4(ℓ− 1)ℓ(ℓ+ 1)(ℓ − 2) [12 + x(x− 2)] + x(x− 2)(x− 5) [4ℓ(ℓ− 1)(x− 6) + (x+ 2)(x− 3)(x− 4)].