Nuclear astrophysics A survey in 3 acts log (abundance) Jeff Blackmon, Physics Division, ORNL Where did this come from? Act II - Stellar obituary 4.

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Transcript Nuclear astrophysics A survey in 3 acts log (abundance) Jeff Blackmon, Physics Division, ORNL Where did this come from? Act II - Stellar obituary 4.

Nuclear astrophysics
A survey in 3 acts
log (abundance)
Jeff Blackmon, Physics Division, ORNL
Where did this
come from?
Act II - Stellar obituary
4. Stellar evolution
s process
5. Supernovae
r process
Mass
Stellar Classification
Aldebaran
Betelgeuse
Alnitak
Rigel
Sirius
Arneb
Stellar evolution
Brighter
Globular cluster
Most stars formed at
about the same time
AGB
He burning
Asymptotic Giant Branch Star
Giant branch
H-shell burning
Convective
envelope
Core H2
exhausts
H burn
Cooler
He burn
CO core
He burning & the “Hoyle” state
t1/2(8Be)=9.7x10-17 s
N 8 B
N  


 5 1010
7.367
8Be+
0+ resonance near the
Gamow energy was
predicted by Hoyle
0+
7.654
e+e-
8Be
4.439
2+
Phys Rev 92 (1953) 1095.
Numerous complementary techniques
12C(p,p’)12C*
, 3, e+e13C(3He,)12C*
Largest uncertainty ee~12%
Experiments now at West. Mich. U.
0+
12C
12C(,)16O
- the “holy grail” ?
The 12C(,)16O reaction rate fixes the
ratio of 12C/16O in the core
Ecm
The 12C/16O ratio substantially affects
the subsequent evolution of the star:
Size of Fe core
Supernova?
300 keV
16O
Influence of subthreshold states
substantial uncertainties in
extrapolation
New Stuttgart measurements:
Kunz et al., PRL (2001)
improvement?
= 0.1 fb
12C(,)16O
- via
16N
 decay
Azuma et al. PRC 50 (1994)
Ecm

16N

12C
16O
New WNSL Measurement
France et al., PRC 75 (2007) 065802.
Approach @ ANL (Tang et al.)
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=S1S2DWBA
W (r)
 C
r
12C(,)16O
SubCoulomb  transfer
to subthreshold states
via ANC
Brune et al. PRL 83 (1999)
6.92 (2+)
7.12 (1-)
12C
16O
DWBA
C2(2+)=(1.30.2) x 1010 fm-1
C2(1-)=(4.30.8) x 1028 fm-1
w/ 16N  decay
16
SE 2 (300keV)  4223
keV  b
SE1 (300keV )  101  17keV  b
Neutron sources in AGB Stars
Stars are thermally unstable: mixing, convection, mass loss
12C(p,)13N(n)13C(,n)16O
22Ne(n)25Mg
H envelope
radius
convective
envelope
driven off
Flash
mixing
13C(,n)
13C(,n)
Convective
pocket
He intershell
CO core
(white dwarf)
CO core
time
Synthesis of heavy elements
• s process
~ 80% of isotopes
(n,) rates needed
Branch points crucial
•
r process
~ 70% of isotopes
Far from stability
See supernovae
•
p process
~ 10% of isotopes
Very low abundance
 n,  ~
1
 v ~ constant
v
Secondary process
(s-wave)
Neglected here
Recipe for untangling r & s
abundances
Calculate s process yields and fit to s only isotopes
log (abundance)
Subtract s abundances from solar system to get r abundances
Mass
Stardust in a haystack
Nittler, Earth Planetary Sci Lett (2003)
Tiny grains isolated from meteorites
Unusual grains identified with SIMS
Nguyen & Zinner, Science 303 (2004) 1496.
( Nd/
Relative abundances for isotopes
of a given element from a single
AGB star
144Nd)/(solar)
(XNd/
144
i
Some grains have preserved isotopic
composition from solar environment
Nd)/s olar
Nd Isotope Ratios in SiC Grains
Meteorite data
Stellar model before ORELA data
Stellar model with new ORELA data
2.0
1.5
Solar Nd
1.0
0.5
Guber et al.
0.0
142
143
144
145
146
Mass Number
147
148
(n,) cross sections for the s process
Good data on most stable
isotopes
Spallation n sources
TOF techniques
Good energy resolution
Often high level densities
ORELA
Maxwellians at kT = 8 and 30 keV
8 keV
30 keV
1.0
Resonance Areas
Influence of low-energy
levels on <v> at low temp
Effect of thermal excitations
in stellar environment
Branch point isotopes
Relative Height
Major outstanding issues
0.8
0.6
0.4
0.2
0.0
0
10
20
En (keV)
30
The new frontier
Source
flight path (m)
resoluti on (ns/m)
power (kW)
flux (n/s/ cm2)
FOM (n/s/c m2)
ORELA Lujan n TOF
40
0.2
8
2x104
5x105
20
6.2
64
5x106
6x109
180
0.05
45
3x105
5x108
SNS
20
18
2000
2x108
9x1010
Experiments now
possible with samples of
only ~ 1016 atoms/cm2.
Important s process branch points
High efficiency detector arrays
High segmentation to handle
rate from radioactive sources
DANCE
status
feasible
Synthesis of heavy elements
• s process
~ 80% of isotopes
(n,) rates needed
Branch points crucial
•
r process
~ 70% of isotopes
Far from stability
See supernovae
•
p process
~ 10% of isotopes
Very low abundance
 n,  ~
1
 v ~ constant
v
Secondary process
Neglected here
The r process site
Galactic chemical evolution arguments favor supernovae as the dominant
source for elements early in the history of the Galaxy  an r process
Argast et al., A&A 416 (2004) 997.
Creation of elements in the early Galaxy
Now many observations of unmixed supernova nucleosynthesis in the Galactic halo
Cowan & Sneden, Nature 440 (2006) 1151.
CS22892-052
Fe/H = (8x10-4) solar = very old
r/Fe = 50 solar
Only 2 known in 2000
Now extensive surveys
Z>55 pattern matches solar
e.g. see Frebel et al., ApJ 652 (2006) 1585
SEGUE (Sloan DSS)
Spectra of >2x105 selected halo stars
Expect ~ 1% with Fe/H < 0.001solar
~36 known r process stars
11 with r/Fe > 10 solar
Distribution Fe/H puzzling
Lowest Fe/H stars intriguing
Frebel et al., Nature 434 (2005) 871.
CS22892
Fe/H < 10-5solar
(C&S, Nature 440)
Z<50 abundances vary
Anatomy of a supernovae
Stars > 10 solar masses
Higher gravity
Faster burning stages
Less mass loss
C burning
O burning
Si burning
In rapid succession
• 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
Standing Accretion Shock Instability
History of SN1987a
QuickTime™ and a
Video 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 interaction
Effect 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 distributions
first-forbidden contribution
Fröhlich et al., PRL 96 (2006)
Abundaces relative to solar
with n reactions
without n reaction
Cole et al., PRC 74 (2006) 034333.
Charge exchange reactions with fast beams at the NSCL
Charge exchange reactions
such as (t,3He) are sensitive
probes for GT strength at
100 – 200 MeV/u
Needed for
• core collapse supernova models
• type Ia supernova models
• neutron star crust processes
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 usage
Need to develop technique for inverse kinematics and radioactive beams
nSNS
SNS neutrino spectrum
nm
0.04
0.035
Neutrino Flux
 A proposal has been submitted to
DOE to construct a facility for
neutrino reaction measurements at
the Spallation Neutron Source.
http://www.phy.ornl.gov/nusns
0.03
nm
0.025
ne
0.02
0.015
0.01
0.005
0
GeV protons
BL18
ARCS
0
5
10
15
20
25
Segmented
Accumulator
Detector
Likely initial
program
ne+OF+e- (450 events/yr)
ne+FeCo+e- (1100 events/yr)
ne+AlSi+e- (1100 events/yr)
ne+Pb Bi+e- (4900 events/yr)
35
40
45
50
Proton beam
(RTBT)
Homogeneous Det.
Hg target
30
Energy, MeV
Cartoon r process
Y(A 1) 1 2 2  Sn /(kT )
 
nne
Y(A)
2 mukT 

Large Sn
Small Sn
(n,) >> (,n) >> t1/2
(,n) >> (n,) >> t1/2
 Free parameters nn, kT, t
 Instantaneous freezeout & decay to stability
Only masses, t1/2, and Pn needed
Calculated r process
QuickTime™ and a
None decompressor
are needed to see this picture.
Results of r process calculations
 Many different n densities needed
 Reasonable fits to A=130,190 peaks
 Not so nice reproduction of
intermediate nuclei
Fission? (Qian & Wasserburg)
Evidence for quenching of
the shell gaps? (Kratz et al.)
Astrophysical environment?
Freezeout effects?
Masses?
P. Hosmer et al. PRL 94 (2005) 112501.
NSCL fast beam r-process campaign: the half-life of 78Ni
3He
t1/2(78Ni): 110 +100-60 ms
+ n -> t + p
Effect of new t1/2 on r process abundances
neutron
Si stack
~ 100 MeV/u
NERO
Particle identification in rare isotope beam
Model Calculation: Half-Lives from
Moeller, et al. 97
Abundance (A.U.)
r-process beam
Observed Solar Abundances
1.E+02
Same but with present 78Ni Result
1.E+01
1.E+00
1.E-01
1.E-02
70
120
170
Mass (A)
78Ni
Half-life of 78Ni measured with 11 events.
Shorter 78Ni half-life leads to
greater production of A=190 peak
The properties of neutron-rich nuclei
are crucial for understanding the
site(s) of the r process and the
chemical history of the Galaxy
220
Mass measurements
Large number of isotopes circulate
and are measured in ring
Matos, Ph.D. Univ. Giessen
Yu. Litvinov et al., NPA756 (2005) 3.
2 modes:
Schottky - slow, more precise
isochronous - fast, less precise
Measurements now crossing
into regime of light r process
The Chart of the Nuclides
http://www.nndc.bnl.gov/chart/
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/
r process
= half-life measurements since 2000 (6th ed.)
(neutron-rich nuclei only)
Only a few measurements in r process path
Structure n-rich nuclei and the r process
Masses, half-lives and Pn are crucial
direct impact on r process abundances.
Must rely on theory.
Dillman et al., PRL 91 (2003) 162503.
Properties like level energies and B(E2) values
provide some direct benchmarks.
Radford et al., PRL 88 (2002) 222501.
Varner et al., EPJ 25 (2005) 391.
@ HRIBF
Jones et al.
Understanding the structure of neutron-rich
nuclei is crucial to improving extrapolations
to more neutron-rich (unmeasured nuclei).
2f7/2
2f5/2
HRIBF
3p1/2?
3p3/2
132Sn(d,p)133Sn
EP (channels)
Ex
The HRIBF
CARIBU
CPT measurements of
very neutron-rich nuclei
Intense beams and high
energy will allow
unique structure
studies, e.g. (p,t)
Intense 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
Atomic number (Z)
Next-generation
RIB Facilities
RIBF (RIKEN), FAIR
(GSI), SPIRAL-II
(GANIL), RIA (USA)
Neutron number (Z)
Ground state properties of
nearly all r process nuclei
up to the A=190 peak can
be measured
Atomic number (Z)
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
Neutron number (Z)