Jeff Blackmon, Physics Division, ORNL

Nuclear astrophysicsA survey in 3 acts

4. Stellar evolution s process

5. Supernovaer process

Mass

log

(abu

ndan

ce)

Where did this come from? Act II - Stellar obituary

Stellar Classification

Aldebaran

Betelgeuse

Alnitak

Rigel

SiriusArneb

Stellar evolutionGlobular cluster

Most stars formed at about the same time

Brig

hter

Cooler

Core HCore H22 exhaustsexhausts

Giant branch

H burn

He burn CO core

Convectiveenvelope

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 s8Be

€

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

keV

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)

Convectivepocket

Synthesis of heavy elements• s process

~ 80% of isotopes(n,) rates neededBranch points crucial

• p process~ 10% of isotopesVery low abundanceSecondary processNeglected here

• r process~ 70% of isotopesFar from stabilitySee supernovae

€

n,γ( ) ~ 1v⇒ σv ~ constant (s-wave)

Recipe for untangling r & s abundances

Mass

log

(abu

ndan

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 148Mass 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

(X Nd/

144 N

d)/(s

olar

)

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 issuesInfluence 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 neededBranch points crucial

• p process~ 10% of isotopesVery low abundanceSecondary processNeglected here

• r process~ 70% of isotopesFar from stabilitySee supernovae

€

n,γ( ) ~ 1v⇒ 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 supernovaeFe 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 solarwith n reactionswithout 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

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 NSCLCole 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

00.005

0.010.015

0.020.025

0.030.035

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)

≈ 12

2πh2

mukT

⎛ ⎝ ⎜

⎞ ⎠ ⎟nneSn /(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 220Mass (A)

Abundance (A.U.)

Observed Solar Abundances

Model Calculation: Half-Lives fromMoeller, e t 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 beamneutron

~ 100 MeV/uSi stack

3He + n -> t + p

NERO

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/

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 technologyNeutron-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

num

ber (

Z)

Neutron number (Z)

Ato

mic

num

ber (

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)