Chapter 3
THE INTERSTELLAR MEDIUM
Introduction
The interstellar medium (ISM): the gas and dust
distributed between stars in a galaxy
In the Milky Way:
mass of gas � mass of dust : Mdust ' 0.1Mgas
ISM is generally a small fraction of a galaxy’s lumi-
nous mass:
' 0 % for an elliptical
' 1− 25 % for a spiral (increases from Sa to Sd)
' 15− 50 % for an irregular
Very diffuse: in the plane of the Galaxy, particle
number density ' 103 to 109 atomic nuclei m−3
Mixture of:
• gas remaining from the formation of the galaxy
• gas ejected by stars
• gas accreted from outside (such as infalling dif-
fuse gas or the ISM of accreted galaxies)
1
Important (1) in galaxy evolution - gas promotes
star formation in denser regions – absorption by dust
allows molecular clouds to cool
Important (2) observationally – emission lines from
gas are prominent and can be used to observe dy-
namics
Chemical composition is about 90 % H, 9% He, plus
a trace of heavy elements (expressed by numbers of
nuclei)
Heavy elements in the gas can be depleted into dust
grains
2
Spectroscopy of Interstellar Gas
Gas in the ISM readily emits detectable radiation
The very low density of gas allows detection of
forbidden lines :
– spectral lines not normally observed in the lab
– low transition probabilities
– in the lab the excited states get collisionally
de-excited before they can radiate
– but in the ISM, although collisional times are
� the lifetimes of the forbidden states, the huge
number of atoms in the ISM means that these
are commonly observed
In astronomy we use notation such as HI, HII, HeI
HeII and HeIII where:
I – neutral atom
II – singly ionised positively charged ion
III – doubly ionised positive ion
etc.
So, HI is H0, HII is H+, HeI is He0, HeII is He+,
HeIII is He2+, LiI is Li0, etc.
A negatively charged ion, such as H−, is indicated
only as H−, although few of these are encountered
in astrophysics
Square brackets indicate a forbidden line e.g. [OII]
3
Figure 1: The Milky Way within +/- 10 deg of Galactic plane (360 deg in longitude)in various wavebands. Note how the appearance of the Galaxy varies from image toimage. The dark areas close to the Galactic plane in the optical waveband representthe obscuring effect of dust. However, the same areas are bright in the infrared images,showing that blue light is preferentially attenuated by dust. The near infrared is largelyunaffected by dust and gives a more accurate map of the distribution of stars, includingthe bright central bulge. Radio continuum images indicate hot ionised gas and highenergy radiation from supernovae. Molecular hydrogen maps molecular clouds - theraw materials for star formation, Atomic hydrogen extends over almost the full 360longitude range and is useful for mapping the outer reaches of the Galaxy, particularits rotation curve at large radii [image credit: Jodrell Bank, Leiden Dwingeloo, MaxPlank Institute, IRAS, COBE, A. Mellinger, J. Friedlander, S. Digel, ROSAT, NASAGoddard Flight Center].
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Cold/Warm Gas: the 21 cm Line ofNeutral Atomic Hydrogen
Atomic hydrogen, HI , emits at 21 cm wavelength
(radio) from hyperfine splitting of ground state
– cool/warm ISM – T ∼ 10 to 100K in high density
regions, 103 to 104K in lower density regions
– ‘spin-flip’ transition – coupling of nuclear and
electron spin – forbidden line
Energy difference between the two spin states:
∆E = 9.4× 10−25 J = 5.9× 10−6 eV
producing emission with a rest wavelength:
λ0 = hc/∆E = 21.1061 cm
and rest frequency:
ν0 = ∆E/h = 1420.41 MHz
5
Transition probability, A = 2.87× 10−15 s−1, so lifetime
of an excited state is ' 1/A = 11 million years
21cm transition itself cannot be observed in a labo-
ratory
In the ISM, the 21cm line is observed primarily in
emission, but can also be observed in absorption
against a background radio continuum source
HI observations have many uses: one critically im-
portant application is to measure the orbital mo-
tions of gas to determine rotation curves in our own
Galaxy and in other galaxies
HI observations can map the distribution of gas in
and around galaxies - the 21 cm radio emission pen-
etrates dust
6
Figure 2: An example of 21cm radio emission in M83, The Southern Pinwheel Galaxy,type SABc, showing its extended disk at 21 cm radio wavelength in red on the left,with the UV image superimposed (near-UV in green, far-UV in blue). On the rightis the UV image only showing near-UV in yellow and far-UV in blue. Each image isapproximately 100 kpc by 100 kpc. [Credit: NASA/JPL/Caltech/VLSA/MPIA].
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Cold Gas: Molecules
Molecular hydrogen and other molecules
T ∼ 10 to 100K, relatively high density (cold
dusty molecular clouds)
Molecular hydrogen, H2, is rare and very difficult to
detect directly
– abundance controlled by ‘sticking’ of HI atoms
to dust grains
– no radio lines, so no direct way of tracing H2 in
cold dense gas clouds using radio observations
– but some H2 band absorption in the far UV can
be detected
– H2 can be photo-dissociated by UV radiation into
atomic hydrogen, HI
Other molecules do emit in radio/microwave
– they act as an indirect tracer of cold dense H2 gas
– CO is particularly useful - has strong lines at
1.3 mm and 2.6 mm from rotational transitions
– CO and H2 densities are similar so we can use CO
as a tracer for H2
8
Hot Gas: HII Regions
Hot gas is readily observed in the optical via emis-
sion lines from largely ionised gas
HII regions are regions of partially ionised hydrogen
around hot young stars (O or B type), with T ∼ 104 K
These stars emit strongly in UV
Any UV (Lyman continuum) photons with wave-
lengths λ < 912 A will photoionise hydrogen produc-
ing H+, i.e. HII ions
The ions and electrons recombine to produce excited
hydrogen atoms
The electrons then cascade down energy levels, emit-
ting photons (radiative decay)
Photons emitted in UV, optical, infrared and radio
– free/bound transitions → continuum radiation
– bound/bound transitions → emission lines
Prominent optical lines from transitions down to
first excited level (n = 2) give the Balmer series
Transitions down to ground state (n = 1) give Lyman
series (in UV)
9
Each series is labelled α, β, γ, δ, ..., in order of increas-
ing energy
Transitions from n to n − 1 levels are the strongest
i.e. α lines are the strongest
Lyman series lines of hydrogen are:
Lyα λ = 1216 A (in ultraviolet)
Lyβ 1026 A ( ” )
Lyγ 973 A ( ” )
Balmer series lines of hydrogen are:
Hα λ = 6563 A (in optical)
Hβ 4861 A ( ” )
Hγ 4340 A ( ” )
Hδ 4102 A ( ” )
Hε 3970 A ( ” )
10
Collisional excitation can also occur
– not for H - no levels accessible at collision
energies characteristic of HII regions (T ∼ 104 K)
– but possible for NII, OII, SII, OIII, NeIII
[OIII] lines at 4959A and 5007A are particularly promi-
nent
Some of the most prominent optical lines of HII re-
gions are:
[OII] 3727 A [OIII] 4959 A
[NeIII] 3869 A [OIII] 5007 A
Hε 3970 A HeI 5876 A
Hδ 4102 A [NII] 6548 A
Hγ 4340 A Hα 6563 A
Hβ 4861 A [NII] 6584 A
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In HII regions one Balmer photon is produced for
each Lyman continuum photon from the hot star
– so observations of Balmer lines of nebula gives
number of UV photons from star
This happens because most H atoms are in the ground
state, and are therefore opaque to Lyman photons
but transparent to others
A Lyman continuum photon from star will be ab-
sorbed by a neutral H atom → ionises H atom to
produce a free electron
The free electron is then recaptured (free-bound tran-
sition), emitting a continuum photon depending on
which state it is captured into :
• If captured into the ground state (n = 1) → emits
another Lyman continuum photon – back to where
we started
• If captured into n = 2 → emits Balmer contin-
uum photon in going to n=2 – one UV photon
produces one Balmer photon – then decays to
n = 1 emitting a Lyα line photon which will al-
most certainly be absorbed again
12
• If captured to n > 2 it can then decay to n=2
emitting a Balmer line photon, or directly to
n = 1 – but a transition to n = 1 emits a Lyman
line photon that can excite another ground-state
H atom, so the process repeats, eventually pro-
ducing a Balmer line photon
So each ionising UV photon from star (λ < 912 A)
will produce on Balmer (line or continuum) photon
HII regions and planetary nebulae also produce ther-
mal continuum radiation – free-free emission:
– the free electrons in the HII can interact with
protons without recombination
– electrons are accelerated, producing radiation
The resulting spectrum is not blackbody because the
gas is transparent to free-free photons: there is no
redistribution of the energy of the free-free photons.
In fact the spectrum is quite flat at radio frequencies
Strengths of the emission lines from HII regions can
provide information on temperature, density and
chemical composition of the interstellar gas
13
Colour optical images of HII regions show strong
red/pink and green colours:
– the red and pink is produced mainly by the Hα line
– the green is produced by [OIII] and Hβ
– HII regions are seen prominently in images of
spiral and irregular galaxies
– their emission lines dominate the spectra of
late-type galaxies and are valuable for use in
measuring redshifts
14
Figure 3: The Orion Nebula, M42. The most famous example of an HII region. Thegas fluoresces because of the UV radiation from the hot young stars, recently formedin a dense region of gas [Hubble Space Telescope: NASA, ESA, M. Robberto (SpaceTelescope Science Institute/ESA), Hubble Space Telescope Orion Treasury ProjectTeam.]
Figure 4: The optical spectrum of the Orion Nebula, showing very strong emissionlines from Hα (red/pink), [OII] (blue), and [OIII] and Hβ (both green).
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Hot Gas: Planetary Nebulae
Like other HII regions except:
– compact regions around hot evolved stars
– gas is ejected by star through mass loss
– UV photons from star ionise gas
– gas emits photons like HII regions (similar
emission process)
– relatively luminous, with prominent emission lines
– also observed in other galaxies
– useful for tracing distribution & kinematics of stars
Figure 5: Examples of planetary nebulae: the Ring Nebula (M57), left, and the HelixNebula, right. Gas has been ejected from a hot, evolved star and the ultravioletradiation from the star ionises the gas. [Images from the Hubble Space Telescope.Ring Nebula: produced by the Hubble Heritage Team (AURA/STScI/NASA). HelixNebula: produced by NASA, NOAO, ESA, the Hubble Helix Nebula Team, M. Meixner(STScI), and T.A. Rector (NRAO).]
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Hot Gas: Supernova Remnants(SNRs)
Supernovae eject material at very high velocities
into the interstellar medium
– this gas shocks, heats and disrupts the ISM
– low density components of the ISM can be
significantly affected
– dense molecular clouds are less affected
– hot gas from supernovae can even be ejected out
of the Galactic disc into the halo
Supernova remnants have strong line emission. They
expand into and mix with the ISM
Figure 6: Examples of supernova remnants: the Crab Nebula (M1), left, and part of theVeil Nebula, right. The Crab Nebula is a very young supernova remnant, producedby a supernova observed in the year 1054. The Veil Nebula is an older example.[Crab Nebula image from the Hubble Space Telescope: NASA, ESA, J. Hester andA. Loll (Arizona State University). Veil Nebula image from the 0.9m Burrell SchmidtTelescope at Kitt Peak National Observatory, Arizona: NOAO/AURA/NSF.]
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Hot Gas: Masers
Some very high density HII regions around young
stars or old evolved stars can show maser emission
– density ∼ 1014 m−3
– population inversion between certain energy states
of molecules due to radiative excitation
– transitions are in the radio
– the overpopulated state decays by stimulated
emission → maser emission
– coherent emission - polarised
– very narrow emission lines of high intensity
– OH and H20 masers are observed (e.g. in Orion)
– useful kinematic tracers
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Hot Gas: Synchrotron Radiation
Broad-band non-thermal radiation emitted by elec-
trons moving relativistically in a magnetic field
– observed in both optical and radio
– photons are emitted in the instantaneous direction
of electron motion
– polarised perpendicular to the magnetic field
Spectacular sources of synchrotron emission are sys-
tems with jets → young stellar objects with bipo-
lar outflows, or active galactic nuclei, lobes of radio
galaxies.
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Absorption Line Spectra
If interstellar gas is seen in front of a continuum
source, light from the source is absorbed at certain
wavelengths
A number of interstellar lines and molecular bands
are seen in absorption
Some absorption features are not well-understood
Particularly problematic are diffuse interstellar bands
in the IR – probably caused by carbon molecules,
possibly polycyclic aromatic hydrocarbons (PAHs)
Cold interstellar CN molecules:
– CN has rotational modes which produce radio
lines, like most heteronuclear molecules
– radio lines can be observed directly, but more
interesting are the optical lines that are split
because of the rotational modes
Optical observations of absorption by cold CN in
continuum spectra of background stars show rela-
tive populations of the rotational modes (from line
strengths) and hence the temperature of the CN
– temperature ' 2.7 K, i.e., these cold clouds are in
thermal equilibrium with the CMB
20
Components of the Gaseous ISM
It’s convenient to divide diffuse gas in the ISM into
distinct components – also called phases:
• cold neutral medium – neutral hydrogen (HI)
and molecules at temperatures T ∼ 10 − 100 K
and relatively high densities
• warm neutral medium – neutral hydrogen (HI)
but at temperatures T ∼ 103 − 104 K and lower
densities
• warm ionised medium – ionised gas (HII) at tem-
peratures T ∼ 104 K and lower densities
• hot ionised medium – ionised gas (HII) at very
high temperatures T ∼ 105 − 106 K but very low
densities
The phases are pressure-confined and are stable in
the long term
Ionisation by supernova remnants is an important
mechanism in producing the hot ionised medium
Cold neutral medium makes up ∼ 50% of the ISM’s
mass, but very small fraction by volume
Supernova remnants, planetary nebulae, and giant
molecular clouds not normally included in these phases
because they are not in pressure equilibrium with
the other components
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Interstellar Dust
Consists of particles of silicates or carbon compounds
– relatively small, but broad range in size
– largest ' 0.5 µm in size with ∼ 104 atoms
– some have . 102 atoms – like large molecules
Dust has a profound observational effect – absorbs
and scatters light – extinction
– diminishes light of background sources e.g.dark
nebulae, zone of avoidance for galaxies at
low galactic latitudes
Dust in galaxies is important because:
– it allows dense molecular clouds to cool - facilitates
star formation
– it catalyses the formation of molecules e.g. molec-
ular hydrogen
22
Interstellar Dust - extinction
The absorption and scattering of light by dust is
called extinction
Light of wavelength λ, specific intensity Iλ (i.e. flux)
passing through an element of interstellar space will
experience a change dIλ in intensity Iλ due to extinc-
tion
This is related to the change dτλ in the optical depth
τλ at the wavelength λ that the light experiences
along its journey by
dIλIλ
= − dτλ
Integrating over the line of sight from a light source
to an observer, the observed intensity is
Iλ = Iλ 0 e − τλ
where Iλ 0 is the light intensity at the source and τλis the total optical depth along the line of sight
What is the loss of light in magnitudes?
Magnitude m in some photometric band is related
to the flux F (i.e. intensity) in that band by:
m = C − 2.5 log10(F )
23
where C is a calibration constant (see introductory
material in Lecture 1 for more information about
the magnitude system)
In the presence of extinction, for a particular wave-
length λ we have (substituting for the expression for
Iλ above) :
mλ = C − 2.5 log10( Iλ 0 e − τλ)
= C − 2.5 log10 Iλ 0 − 2.5 log10( e − τλ)
= C − 2.5 log10 Iλ 0 −2.5 ln( e − τλ)
ln(10)
= C − 2.5 log10 Iλ 0 −2.5
ln(10)τλ
= C − 2.5 log10 Iλ 0 + 1.086 τ
So the observed magnitude m is related to intrinsic
magnitude m0 by:
mλ = m0 + Aλ
where the intrinsic magnitude, m0, is the magnitude
that the star would have in the absence of interstel-
lar extinction, and A is the extinction in magnitudes,
given by:
Aλ = − 2.5 log10(e−τ ) = + 1.086 τ
Note that Aλ depends on the photometric band
24
For example, for the V (visual) band (yellow-green,
centred at 5500 A) :
V = V0 + AV
For the B (blue) band (centred at 4400 A):
B = B0 + AB
Aλ is a strong function of wavelength and scales as
Aλ ∼ 1/λ (not as strong as Rayleigh Law ∼ 1/λ4)
There is much stronger absorption in the blue than
in the red → reddening by interstellar dust
25
The interstellar extinctionlaw. The extinction causedby dust is plotted againstwavelength and extends fromthe ultraviolet through to thenear-infrared.[Based on data from Savage& Mathis, Ann. Rev. As-tron. Astrophys., 1979.]
Colour indices, e.g. B − V , are reddened so that the
observed value is:
(B − V ) = (B0 + AB)− (V0 + AV )
= (B0 − V0) + (AB − AV )
≡ (B − V )0 + EB−V
where (B − V )0 is the intrinsic value (no extinction)
and EB−V = AB − AV is the colour excess or redden-
ing, which tells us how reddened a source is, based
on the extinctions in the two magnitudes
For the V photometric band, AV ' 3 EB−V (as shown
in the plot above)
If the intrinsic colour, (B − V )0, can be predicted
from spectrum, then EB−V can be calculated using
EB−V = (B − V ) − (B − V )0
EB−V data can then be used to map the dust distri-
bution in space
26
Extinction gets less severe for λ & 1µm as the wave-
length gets much longer than the grains
For sight lines through the Galaxy at the Galactic
poles: AV ' 0.00 to 0.05 mag
At intermediate galactic latitudes: AV ' 0.05 to 0.2 mag
In Galactic plane, extinction can be many magni-
tudes in V and UV (less in IR). Distribution can be
patchy (e.g. Baade’s Window, in the bulge)
Towards the Galactic Centre: AV � 20 mag
X-rays can pass through dust grains (AK ' 3 mag)
27
Summary of terminology:
B ≡ mB, V ≡ mV etc
B − V is a colour
Extinction is the absorption and scattering of light
by dust
Aλ is the extinction at some wavelength, λ
AV is the extinction in the visual waveband (for ex-
ample)
E(B − V ) ≡ EB−V is the colour excess or ‘reddening’
AVEB−V
is the ‘reddening ratio’ for interstellar dust in
the visual band ∼ 3.0 (see the graph on page 26)
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Interstellar Dust: polarisation
Dust grains are not spherical
Spinning dust grains tend to align with their long
axes perpendicular to the local magnetic field (Davis-
Greenstein Effect)
Preferentially block light perpendicular to the mag-
netic field: extinction produces polarised light
Polarisation will tend to be parallel to the magnetic
field: polarisation measurements of starlight reveal
the direction of the magnetic field
Dust also reflects light, with some polarisation: ob-
servable as reflection nebulae, where faint diffuse
starlight can be seen reflected by dust
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Interstellar Dust: Radiation byDust
Dust absorbs light - warms the dust - re-emits the
light as black-body radiation (approximately)
So dust has diffuse black-body emission superim-
posed on reflected starlight spectrum
Wien’s displacement law states that the maximum
of the Planck function Bλ of a black body at a tem-
perature T is found at a wavelength
λmax =2.898× 10−3
TK m
This predicts that the peak of the black-body spec-
trum for dust at a temperature of T = 10 K will be
at a wavelength λmax = 290 µm
for dust at T = 100 K will be at a wavelength λmax =
29 µm
and for T = 1000 K will be at λmax = 2.9 µm
So radiation emitted by dust will found in the in-
frared, given the expected temperatures of dust
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Star Formation in the ISM
Stars form by collapse of dense regions of the ISM
under their own gravity
i.e. in cores of molecular clouds, where gas is cold (∼10 K) and densities relatively high ≥ 1010 molecules
m−3
Figure 7: A star-forming HII region within M16, The Eagle Nebula. The blue-greencolour from the mostly ionised gas is caused by the light of [OIII] and Hβ emissionlines from neutral hydrogen atoms. The gas is being ionised by ultraviolet radiationfrom hot, young stars off the top of the picture. The dark pillars, in contrast, areregions of cold, dense molecular hydrogen gas in which star formation is occurring.They are dark because the cold molecules emit virtually no light and because of theabsorption of light by dust mixed with the gas. The ultraviolet radiation is ‘burning’away the surface of the cold gas by photoionisation. [Hubble Space Telescope imageproduced by NASA, ESA, STScI, J. Hester and P. Scowen (Arizona State University).]
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A region of cold gas collapses when its gravitational
self-attraction > hydrostatic pressure support
For gas of uniform density ρ, the Jeans length λJ is
the diameter of a region of gas just large enough for
gravitational force to exceed pressure support:
λJ = cs
√π
Gρ
where cs is the speed of sound in the gas
The Jeans mass is the mass of a region that has a
diameter equal to the Jeans length:
MJ =
(4
3π(λJ/2)3
)ρ =
π
6ρ λ 3
J
The free-fall collapse time Tff , is the time taken for
a static cloud to collapse under its own gravity in
the absence of gas pressure
The free-fall collapse time for a spherically symmet-
ric distribution of mass with a total mass M and
initial radius R to collapse from rest is
Tff = π
√R3
8GM=
√3π
32
1
Gρ
where ρ is the mean density before the collapse starts
(see also the introductory material in lecture 1)
Star formation can be self-propagating
Stars form, heat up and ionise cold molecular gas →outward flow of gas compresses gas ahead of it
32
Causes instabilities – locally collapse to form new
stars
The enhanced density in the spiral arms of spiral
galaxies means that star formation occurs preferen-
tially in the arms
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