Jeff Blackmon, Physics Division, ORNL
Nuclear astrophysics
A survey in 3 acts
4. Stellar evolution s process
5. Supernovaer process
Mass
log
(a
bun
dan
ce)
Where did this come from? Act II - Stellar obituary
Stellar Classification
Aldebaran
Betelgeuse
Alnitak
Rigel
SiriusArneb
Stellar evolution
Globular cluster
Most stars formed at about the same time
Brig
hter
Cooler
Core HCore H22
exhaustsexhausts
Giant branch
H burn
He burnCO core
Convective
envelope
H-shell burningH-shell burning
Asymptotic Giant Branch StarHe burningHe burning
AGBAGB
He burning & the “Hoyle” state
e+e-
7.654
4.439
12C
0+
0+
2+
8Be+7.367
t1/2(8Be)=9.7x10-17 s
8Be
€
N 8B( )
N α( )≈ 5 ×10−10
0+ resonance near the Gamow energy was predicted by Hoyle
Phys Rev 92 (1953) 1095.
Numerous complementary techniques12C(p,p’)12C*13C(3He,)12C*
, 3, e+e-
Largest uncertainty ee~12%Experiments now at West. Mich. U.
12C(,)16O - the “holy grail” ?Ecm
16O
The 12C(,)16O reaction rate fixes the ratio of 12C/16O in the core
The 12C/16O ratio substantially affects the subsequent evolution of the star:
Size of Fe coreSupernova?
Influence of subthreshold states substantial uncertainties in extrapolation
= 0.1 fb
300
ke
V
Kunz et al., PRL (2001)
New Stuttgart measurements: improvement?
12C(,)16O - via 16N decayEcm
16O
16N
12C
Azuma et al. PRC 50 (1994)
Approach @ ANL (Tang et al.)France et al., PRC 75 (2007) 065802.
New WNSL Measurement
12C(,)16O via ANC
A nucleon or “cluster” of nucleons (no internal degrees of freedom) is transferred from one nucleus to another.
The core nuclei are unperturbed.
exp=S1S2DWBA
€
ψ → CW (r)
r
16O
12C
6.92 (2+) 7.12 (1-)
Brune et al. PRL 83 (1999)
C2(2+)=(1.30.2) x 1010 fm-1
C2(1-)=(4.30.8) x 1028 fm-1 DWBADWBA
€
SE 2(300keV ) = 42−23+16keV ⋅b
€
SE1(300keV ) =101±17keV ⋅b
12C(,)16O via ANCSubCoulomb transfer to subthreshold states
w/ 16N decay
Neutron sources in AGB Stars
H envelope
He intershell
FlashFlash
mixing
CO coreCO core
radi
us
time
12C(p,)13N()13C(,n)16O
13C(,n)
13C(,n)
2222Ne(Ne(n)n)2525MgMg
Stars are thermally unstable: mixing, convection, mass loss
convective envelope driven off
CO core(white dwarf)
Convective
Synthesis of heavy elements
• s process
~ 80% of isotopes
(n,) rates needed
Branch points crucial
• p process
~ 10% of isotopes
Very low abundance
Secondary process
Neglected here
• r process
~ 70% of isotopes
Far from stability
See supernovae
€
n,γ( ) ~1
v⇒ σv ~ constant (s-wave)
Recipe for untangling r & s abundances
Mass
log
(a
bun
dan
ce)
Calculate s process yields and fit to s only isotopes
Subtract s abundances from solar system to get r abundances
Stardust in a haystack
142 143 144 145 146 147 148
Mass Number
0.0
0.5
1.0
1.5
2.0
Meteorite dataStellar model before ORELA dataStellar model with new ORELA data
Solar Nd
Nd Isotope Ratios in SiC Grains
(XN
d/14
4 Nd
)/(s
ola
r)
Tiny grains isolated from meteorites
Unusual grains identified with SIMS
Nguyen & Zinner, Science 303 (2004) 1496.Nguyen & Zinner, Science 303 (2004) 1496.
Nittler, Earth Planetary Sci Lett (2003)
Guber et al.
Some grains have preserved isotopic composition from solar environment
Relative abundances for isotopes of a given element from a single AGB star
(n,) cross sections for the s process
0 10 20 30En (keV)
0.0
0.2
0.4
0.6
0.8
1.08 keV 30 keV
Resonance Areas
Maxwellians at kT = 8 and 30 keVORELA
Good data on most stable isotopes
Spallation n sources
TOF techniques
Good energy resolution
Often high level densities
Major outstanding issues
Influence of low-energy levels on <v> at low temp
Effect of thermal excitations in stellar environment
Branch point isotopes
The new frontierSource ORELA Lujan n TOF SNSflight path (m) 40 20 180 20resolution (ns/m) 0.2 6.2 0.05 18power (kW) 8 64 45 2000flux (n/s/cm2) 2x104 5x106 3x105 2x108
FOM (n/s/cm2) 5x105 6x109 5x108 9x1010
Experiments now possible with samples of only ~ 1016 atoms/cm2.
DANCE
Important s process branch points status feasibleHigh efficiency detector arrays
High segmentation to handle rate from radioactive sources
Synthesis of heavy elements
• s process
~ 80% of isotopes
(n,) rates needed
Branch points crucial
• p process
~ 10% of isotopes
Very low abundance
Secondary process
Neglected here
• r process
~ 70% of isotopes
Far from stability
See supernovae
€
n,γ( ) ~1
v⇒ σv ~ constant
The r process site
Argast et al., A&A 416 (2004) 997.
Galactic chemical evolution arguments favor supernovae as the dominant source for elements early in the history of the Galaxy an r process
Creation of elements in the early Galaxy
CS22892-052Fe/H = (8x10-4) solar = very oldr/Fe = 50 solar
Only 2 known in 2000Now extensive surveys
e.g. see Frebel et al., ApJ 652 (2006) 1585SEGUE (Sloan DSS)Spectra of >2x105 selected halo stars Expect ~ 1% with Fe/H < 0.001solar
~36 known r process stars11 with r/Fe > 10 solarDistribution Fe/H puzzlingLowest Fe/H stars intriguing
Now many observations of unmixed supernova nucleosynthesis in the Galactic halo
Cowan & Sneden, Nature 440 (2006) 1151.
CS22892
Z<50 abundances vary
Z>55 pattern matches solar
(C&S, Nature 440)
Frebel et al., Nature 434 (2005) 871.
Fe/H < 10-5solar
Anatomy of a supernovae• Fermi degeneracy initially supports core• Shell Si burning increases core size of• Electron capture on nuclei in core begins
to reduce pressure support• Core undergoes runaway collapse• Reaches supernuclear densities & shock
rebounds -- EOS important• Mechanism involves interplay of
hydrodynamics and nuclear physics• Spherical models fail to explode• Multidimensional effects are critical
Stars > 10 solar massesHigher gravityFaster burning stagesLess mass loss
C burningO burningSi burning
In rapid succession
Standing Accretion Shock Instability
Lecture 3
History of SN1987a
QuickTime™ and aVideo decompressor
are needed to see this picture.
Nucleosynthesis sites in supernovae
Fe group nuclei produced from nuclear statistical equilibrium
Environment above neutron star is likely site for the r process
Influence of weak interactionEffect of e-capture rates on
formation of the shock Electron capture rates affect the
formation of the shock wave.
Neutrino interactions play a role in driving the explosion.
Neutrino induced reactions alter nucleosynthesis.
Weak rates in this mass region are not well understood:
GT strength distributionsfirst-forbidden contribution
Fröhlich et al., PRL 96 (2006)
Abundaces relative to solar
with n reactions
without n reaction
Abundaces relative to solar
with n reactions
without n reaction
Special case or systematic issue? Need systematic measurements for entire relevant range(especially beyond fp shell where nuclear models become much simpler)can help decide which theoretical model to use and can help to improve theoretical models for supernova usageNeed to develop technique for inverse kinematics and radioactive beams
Special case or systematic issue? Need systematic measurements for entire relevant range(especially beyond fp shell where nuclear models become much simpler)can help decide which theoretical model to use and can help to improve theoretical models for supernova usageNeed to develop technique for inverse kinematics and radioactive beams
Charge exchange reactionssuch as (t,3He) are sensitiveprobes for GT strength at100 – 200 MeV/u
Needed for• core collapse supernova models• type Ia supernova models• neutron star crust processes
Charge exchange reactions with fast beams at the NSCL
Cole et al., PRC 74 (2006) 034333.
SNS
BL18ARCS
Proton beam (RTBT)
Homogeneous Det.
A proposal has been submitted to DOE to construct a facility for neutrino reaction measurements at the Spallation Neutron Source.
Segmented Detector
GeV protonsAccumulator
Hg target
0
0.005
0.01
0.015
0.02
0.025
0.03
0.035
0.04
0 5 10 15 20 25 30 35 40 45 50
Energy, MeV
Neutrino Flux
SNS neutrino spectrum
e
e+OF+e- (450 events/yr)
e+FeCo+e- (1100 events/yr)
e+AlSi+e- (1100 events/yr)
e+Pb Bi+e- (4900 events/yr)
Likely initial program
http://www.phy.ornl.gov/nusns
Cartoon r process
Free parameters nn, kT, t Instantaneous freezeout & decay to stability
Large Sn
(n,) >> (,n) >> t1/2
Small Sn
(,n) >> (n,) >> t1/2
€
Y (A +1)
Y (A)≈
1
2
2πh2
mukT
⎛
⎝ ⎜
⎞
⎠ ⎟nne
Sn /(kT )
Only masses, t1/2, and Pn needed
Calculated r process
QuickTime™ and aNone decompressor
are needed to see this picture.
Many different n densities needed
Reasonable fits to A=130,190 peaks
Not so nice reproduction of
intermediate nuclei
Evidence for quenching of the shell gaps? (Kratz et al.)
Masses?
Freezeout effects?
Fission? (Qian & Wasserburg)
Astrophysical environment?
Results of r process calculations
1.E-02
1.E-01
1.E+00
1.E+01
1.E+02
70 120 170 220
Mass (A)
Abundance (A.U.)
Observed Solar Abundances
Model Calculation: Half-Lives fromMoeller, et a l. 97
Same but with present 78Ni Result
1.E-02
1.E-01
1.E+00
1.E+01
1.E+02
70 120 170 220
Mass (A)
Abundance (A.U.)
Observed Solar Abundances
Model Calculation: Half-Lives fromMoeller, et a l. 97
Same but with present 78Ni Result
Effect of new t1/2 on r process abundances
Particle identification in rare isotope beam
78Ni
t1/2(78Ni): 110 +100-60 ms
NSCL fast beam r-process campaign: the half-life of 78Ni
P. Hosmer et al. PRL 94 (2005) 112501.
r-process beam
neutron
~ 100 MeV/uSi stack
3He + n -> t + p
NERO
Half-life of 78Ni measured with 11 events.Half-life of 78Ni measured with 11 events.
The properties of neutron-rich nuclei are crucial for understanding the site(s) of the r process and the chemical history of the Galaxy
Shorter 78Ni half-life leads to greater production of A=190 peak
Mass measurements
2 modes:Schottky - slow, more preciseisochronous - fast, less precise
Yu. Litvinov et al., NPA756 (2005) 3.
Measurements now crossing into regime of light r process
Matos, Ph.D. Univ. Giessen
Large number of isotopes circulate and are measured in ring
The Chart of the Nuclides
http://www.nndc.bnl.gov/chart/
= half-life measurements since 2000 (6th ed.)(neutron-rich nuclei only)
The Chart of the Nuclides
http://www.nndc.bnl.gov/chart/
= half-life measurements since 2000 (6th ed.)(neutron-rich nuclei only)
Only a few measurements in r process path
r process
EP (channels)
Ex
2f7
/2
3p3
/2
2f5
/2
3p1
/2?
Jones et al.
132Sn(d,p)133Sn @ HRIBF
Preliminary
Structure n-rich nuclei and the r process
Radford et al., PRL 88 (2002) 222501.Varner et al., EPJ 25 (2005) 391.
HRIBF
Masses, half-lives and Pn are crucial direct impact on r process abundances.
Dillman et al., PRL 91 (2003) 162503.
Must rely on theory.
Understanding the structure of neutron-rich nuclei is crucial to improving extrapolations to more neutron-rich (unmeasured nuclei).
Properties like level energies and B(E2) values provide some direct benchmarks.
The HRIBF
CARIBUIntense 252Cf fission source
under construction at ATLAS
Gas stopping technology
Neutron-rich RIBs will push the boundaries of our knowledge
Different region on nuclei complementary to HRIBF
CPT measurements of very neutron-rich nuclei
Intense beams and high energy will allow unique structure studies, e.g. (p,t)
Next-generation RIB Facilities
Ato
mic
nu
mb
er
(Z)
Neutron number (Z)
Ato
mic
nu
mb
er
(Z)
Neutron number (Z)
Ground state properties of nearly all r process nuclei up to the A=190 peak can be measured
Nuclear structure studies far from stability will greatly improve our ability to extrapolate to the unknown
Understanding observations of the oldest stars and the origin of the heavy elements in our Galaxy
RIBF (RIKEN), FAIR (GSI), SPIRAL-II
(GANIL), RIA (USA)