An Intellectual Opportunity This is a special time Wealth of new - - PowerPoint PPT Presentation

Understand nuclear processes that

• Power the stars
• Synthesize the elements
• Mediate explosive

phenomena Determine

• Nature of stellar evolution
• Sites of astrophysical

processes

• Properties of universe
• Neutrino properties

For background, see

http://www.nscl.msu.edu/~ austin/ nuclear-astrophysics.pdf

Nuclear Physics in the Cosmos

An Intellectual Opportunity This is a special time

• Wealth of new astronomical observations--require new

nuclear data for a credible interpretation

• New accelerators of radioactive nuclei to provide this

data

• Growing computational power to simulate the

phenomena

The observables: Cosmic abundances, abundances in the solar system and elsewhere, fluxes of gammas and neutrinos. Nature of the nuclear processes involved:

• Reaction rates
• Resonant and non-resonant processes
• Technical details: Gamow peak, S-factor, etc.

The Big Bang and the Nature of the Universe Baryons, dark matter, dark energy Stellar evolution with some digressions

• Quasistatic evolution, solar neutrinos, stellar onion
• Explosive phenomena: supernovae, r-process, neutrinos
• Binary systems: x-ray bursters and x-ray pulsars, the

surface of neutron stars. Outline of the Lectures

Outline-Continued References (www and Google) +

• “Cauldrons in the Cosmos”, Rolfs and Rodney (out of print?)
• “Principles of Stellar Evolution and Nucleosynthesis” D.D. Clayton,
• U. Chicago Press, paperback
• Ann. Revs. of Nuclear and Particle Science; Astronomy and

Astrophysics

• Web pages of the major instruments: WMAP, SNO, Super-

Kamikande, Chandra, HST,…..

• Joint Institute For Nuclear Astrophysics (JINA) web page:

www.jinaweb.org. See there also the link to the Virtual Journal for Nuclear Astrophysics.

Cosmic Abundances (Really solar system, mainly) A qualitative view-Suess-Urey Plot

• Very large range of

abundances

• Names denote

various creation processes

Group Mass Fraction

1,2H

0.71

3,4He

0.27 Li, Be, B 10-8 CNO Ne2x10-2 Na-Sc 2x10-3 A= 50-62 2x10-4 A= 63-100 10-6 A> 100 10-7

Neut ron Capt ures Log Abundance

A

A More Detailed Picture

50 100 150 200 250 mass number 10

• 3

10

• 2

10

• 1

10 10

1

10

2

10

3

abundance

Makes most of Gold and Platinum Makes Uranium

50 100 150 200 250 mass number 10

• 13

10

• 12

10

• 11

10

• 10

10

• 9

10

• 8

10

• 7

10

• 6

10

• 5

10

• 4

10

• 3

10

• 2

10

• 1

10 number fraction

Rapid n-capture(r) process all processes

Solar abundances

mostly s + r

R-Process abundance

Nuclides made by slow (s) and rapid (r) neutron capture Model s process, fairly accurate Subtract from solar r

Populations I, II and III

What about elsewhere?

In the halo of the galaxy find (old) stars (Pop II stars) with small (10-4) abundances

• f metals (A > 4) compared to the solar

system values typical of Pop I stars.

Pop II stars

• Reflect processes in the

early galaxy

• Investigation of Pop II

stars is a hot area of astrophysics

What are Pop III stars?

• Stars that produce the

material from which Pop II are made.

• Probably very large (>

100 Msun) fast evolving stars made from products

• f the Big Bang.

The Stars as Element Factories Interstellar Gas Dust Stars Nuclear Reactions Element Synthesis Condensation Ejection-Supernovae Planetary nebulae Star Forming Region DEM192-LMC Supernova remnant N132D-LMC

10-20 1 1020 1040

TIME AFTER BIG BANG (seconds) TEMPERATURE (K)

Creation of matter Elementary particles

1020 1010 1 10-10 quark/gluon hadron Light elements

Now

Stars

TIME

N U C L E A R P H Y S I C S

3o 1010 years

The Big Bang

Some Milestones

100 sec--Light elements (1,2H,

3,4He, 7Li) made

300 kyear—Atoms form, CMB 200 Myear--First stars form

Assumptions:

• General relativity
• Universe isotropic,

homogeneous

• Tnow= 2.735 K (CBR))

Production of elements

• 10-300 sec after BB
• T ≈1010 K, ρ ≈ 1g/cm3
• Big Bang produces only

lightest elements: 1,2H, 3,4He,

7Li, because there are no

stable mass 5 or 8.

• Yield depends on density ρB
• f baryons

Nucleosynthesis in the Big Bang Reaction network

Need to know noted reactions- = Poorly known reactions

Can we Determine the Baryon Density from the Big Bang? Method

• Find ρB where predicted and
• bserved abundances equal.
• If ρB same for all nuclides,

assume it is the universal density Result OK, EXCEPT for 7Li. Perhaps predicted abundance wrong (poor cross sections) or primordial Li higher (star destroys). New data on the CMB may change our conclusions

Nollett and Burles, PRD 61,123505 (2000)

It’s Close, Why Does It Matter? Cosmic Background Radiation

Surrounds us, Planck distribution (T~ 2.7 K), remnant of early BB Fluctuations (at 10-5 level) give information on total density of Universe and on ρB.

It implies

Universe is just bound Ωtot = 1 Baryon density ρB ~ 0.04 Dark matter density ρD ~ 0.23 perhaps WIMPS, weakly interacting massive particles Dark energy Λ?) ρΛ ~ 0.73

Era of precision cosmology

Far reaching conclusions must be checked

Value of ρB

Best possibility. Need more accurate cross sections for several reactions affecting 7Li.

Supernova Ia

A standard candle to measure rate of expansion. Universe accelerating, measures dark energy.

Aside-Nature of Cosmic Background Radiation

WWW site: http://map.gsfc.nasa.gov/m_mm.html

WMAP: C.L. Bennet, et al, Relative Temperature,

angular resolution 0.3 deg See February 2004 Scientific American

Angular Power Spectrum Analysis (G. Hinshaw, et al.)

Perform angular multipole decomposition

Results

Good agreement with earlier results, summarized in red points Strong peak at l = 200 = > Ωtot= 1 Secondary peak l = 500 = > Baryon density

Resulting Cosmological Parameters Model parameters: from WMAP + earlier CMB measurements

(COBE, CBI, ACBAR) + large scale galactic structure + Lyman forest

Other results: First stars at 200 Myear, Mν < 0.23 eV

Big Bang Nucleosynthesis Revisited Assume: know η (photons/baryons) Predict: BB nucleosynthesis Result: Agrees with observation for

2H, not for 4He, 7Li

To sharpen comparison need

Better cross sections for

3H(α,γ)7Be for 7Li

p(n, γ), d(p,γ), d(d,n) for 2H

3He(d,p) and d(p,γ) for 3He

Better abundance measurements Especially for 7Li and 4He

Do stellar astrophysics with BB?? Details: RH Cyburt et al., PLB 567,

227 (2003) Gray-observation. Black--BB

What energy source powers the stars?

Mass initial Mass final Mass converted = f Mass initial Energy released f Massinitial c2

Reaction

All energy comes from mass Must provide solar luminosity for > 4.6 x 109 yrs L sun = 3.826 x 1033 erg/sec M sun = 1.989 x 1033 g Of the possibilities

f chemical ≈ 1.5 x 10-10 ⇒ 2200 yrs f gravity

⇒ 107 yrs

f nuclear ≈ 0.007 ⇒1011 yrs

Only nuclear sufficient Other evidence

Technetium is seen in stellar

• spectra. BUT the longest lived

isotope is unstable--lifetime of 4 x 106 yrs. Must have been synthesized in the star.

Reaction Rates and Energy Scales Reaction Rate

• Ionized gas (plasma) with Ni

/cm3 of species “i”

• Assume species x moving at

velocity v through species y at

• rest. Rate of reactions rxy is rxy

= NxNyvσxy

• Average over velocity

distribution (Max. Boltz.) Einc Turning point

Environment

• k = 8.6171 x 10-5 eV/K
• T = 107-1010 K ⇒kT= 1-900 keV
• Coulomb barriers MeV range
• Reactions are often far sub-coulomb

rxy = NxNy(1+ δxy)-1< vσxy>

}

# of pairs/cm3

σxy(E) ∝ tunneling probability

for point coulomb charge

Non-Resonant and Resonant Reactions Non-Resonant

Typical case: Direct capture at stellar energies-light nuclei

Resonant Capture

Common for all but lightest nuclei

Nature of Cross Sections S Factor = σE exp(b/E1/2)

Increase Rapidly with Energy Removes penetrability, nearly constant away from resonance

Example –7Be(p,g)8B

What Energies are Important?

E0 = 5.9 keV p + p 27 keV p+ 14N 56 keV α + α 237 keV 16O+ 16O

Cross sections at Eo too small to be measured Gamow Peak:

Maximum in product of MB distribution and penetrability of Coulomb barrier

S contains the nuclear structure information-At what energy do we need

to determine it?

S -- Resonant and Non-Resonant Phenomena

Resonance in Resonance in Gamow Gamow Peak Peak dominates the rate dominates the rate

• Rate ∝ ΓpΓγ/(Γp+Γγ)•exp(-Er/kT)
• Measure: Γs, Er ⇒ Rate.
• Γs may be strong functions of E
• Classic expts. with low-E accelerators:

small σ’s at low-E

• Measure cross sections to low-E,

extrapolate to Eo to extract S-Factor.

• Long used for resonant rates, esp. Er
• Recent emphasis on new techniques to

measure non-resonant rates. No resonance--Rate characterized by slowly varying S factor at low energy. Role of High-E facilities

Screening History of a trajectory

• Large r: Ion sees no charge—nucleus screened by electrons
• Small r: Sees full nuclear charge, but in a reduced potential:
• Barrier: Z1Z2e2/Rnucleus - Z1Z2e2/Ratomic , height reduced, σ larger
• Effect usually small (except at the lowest energies), and is to increase

the cross section. Similar effect in stellar plasma. Need to remove the lab screening, put in the stellar screening. Examples: ∆barrier Ue ∆σ (5 keV) p+ p 29 eV 2.3%

α+ 12C

2.1 keV

16O + 16O

14.7 keV

How it looks! Image of Sun: Goddard Space Flight Center (http://antwrp.gsfc.nasa.gov/apod/i mage/9709/solprom1_eit_big.jpg) How it works! Gravity pushes inward, but the center

• f the sun in heated by nuclear

reactions, making a high pressure that pushes outwards. They balance, and the sun just sits there burning its nuclear fuel. Has gone on for 4.5 billion years and will continue for another 5 billion years. T=

Core: H + He(25%) Density=150g/cm3 Temp = 15 x 106 K Pressure--out Gravity--in

Nature of Stellar Evolution

One Page- One Nobel- 1967

...for his contributions to the theory of nuclear reactions, especially his discoveries concerning energy production in stars.

A Scenario-H.A. Bethe (CNO Cycles)

Physical Review 55, 103(L) 1939.

“Energy Production in Stars”

Low E n.s

The pp Chains and Neutrino Sources

Observing the Center of the Sun with Solar Neutrinos Problem

• Can’t look with telescopes
• Light is absorbed in L, re-

emitted in random direction.

• Drunkard’s walk: Distance

covered = (N)1/2 L

• N number of steps;
• L length of a step.

Result:

For sun, L = 0.1 cm, D = 6.96 x 1010

• cm. Takes: 5 x 104 yr

Sun

SUN D Look at emitted neutrinos

• Made in solar cycle, escape

without hindrance

• Nn ~ T18, measuring flux measures

T at center of sun

But it’s hard

• ν hardly interacts
• Need a huge detector
• SNU, Solar Neutrino Unit, 1

event/sec/1036 detector atoms

Solar Neutrino Spectra-Detector Thresholds

, SNO (p + e + p --> 2H + ν)

3He + e + p ->4He + ν

ν ref.: http://www.sns.ias.edu/~jnb

The Detector

• 100,000 gallons cleaning fluid

(perchlorethylene C2Cl4), Homestake gold mine, S.D.

ν + 37Cl e + 37Ar - Inverse β-

decay

• Collect by bubbling He through

tank (every 30 days)-count radioactive 37Ar

Motivation

“To see into the interior of a star and thus verify directly the hypothesis of nuclear energy generation.”

First experiment-R. Davis (1968)

Nature of the SN Problem

• A one-number problem
• Solution in solar physics?

nuclear physics? particle physics?

• Motivated a search for the

cause: 1968 to present

• Better solar models,

improved input nuclear physics.

Implications of the Davis Experiment Results

Expected 2 37Ar/day, 8 SNU.

Got 0.5/day, 2 SNU-a shocker! “The solar neutrino problem”

New Experiments

Different neutrino energy sensitivities. Gallex/Sage (Ga)--p-p, 7Be νe Super-K(H2O)------8B νe, (νx) Davis( Cl)----------8B, 7Be νe SNO (D2O)---------8B νe, νx Others yet to come, see http://www.sns.ias.edu/~ jnb

New Models and New Experiments New Models

The Super-Kamiokande Detector Japan, US, Korea, Poland Properties

• 50,000 tons H2O, 11200 P.M.s
• 1000m underground, Mozumi

mine, Kamioka Mining Co.

• Observe νe-e scattering (EC)

(mainly)-via Ĉerenkov light

Light multipliers n e Cerenkov light

SNO—The Sudbury Neutrino Detector Unique characteristics

• 1000 tons heavy water (D2O)
• See electron neutrinos and

muon and tau neutrinos

• Charged current (CC)

νe + D p + p + e-

• Neutral current (NC)

νx + D νx + p + n νx + e νx + e (ES)

Location

• 6800 feet under ground,

Creighton mine Sudbury, Ontario.

• Canada, US, UK

A Neutrino Picture of the Sun---Super-Kamiokande

500 Days of Data

Neutrinos vs. angle from the sun Some details

• Brighter colors-more intensity
• Each pixel twice diameter of

sun

• Good angular resolution of the

experiment permits a great reduction in the isotropic background

Solar Neutrino Experiments--Summary Standard solar model vs. Expt.

Bahcall Pinsonnealt, 2000, updated. For solar neutrino details see http://www.sns.ias.edu/~ jnb/

It appears: different

fractions of neutrinos arrive at the detectors

• All of p-p ν’s
• ~ 0.5 of 8B ν’s
• Few of 7Be ν’s

As compared to the standard solar model

Note:

SNO differs from S-K because S-K sensitive to νx at 1/6 level of νe

Neutrino oscillations

• Neutrinos have mass. Oscillate into

another type of neutrino. (νeνµ)

• Detector not sensitive to these

neutrinos

• Probability of survival: P(νe→νe)
• P(νe→νe) = 1-(sin2ΘV) sin2(a∆m2d/E)
• ∆m2 = (m1

2 -m2 2)

• Passage through matter changes

the constraints-resonant conversion.

How might this happen Flaws in the physics input?

Stellar physics--no, details to be

• settled. A check from

Helioseismology Nuclear Physics--no, but better prediction of fluxes needed .

Properties of neutrinos--the answer

Determining ∆m2 and ΘV Survival Probability

Show two cases: large mixing angle (now favored) and small mixing angle. Analysis is complex and subtle Solar ν’s + Kamland

Analysis Results

Take into account results from SuperKamiokande, SNO and Kamland (observation of reactor neutrinos at Kamiokande)

Neutral Currents from SNO First results (SNO--S-K)

Ascribe difference in SNO CC rate (νe) and S-K ES rate (νe + νx) to νx Extract νx –agrees with Standard Solar Model (SSM)

Then from SNO Combine all three SNO

detection modes CC, NC, ES

Results

• φ(νe) = 1.76 x 106 cm-3sec-1; φ(νµ + ντ ) = 3.41 x 106 cm-3sec-1
• Total number of neutrinos in good agreement with SSM- solar neutrino

problem resolved. Prediction

Observation

What Next? Remaining issues Pin down the neutrino mass differences, absolute masses Pin down the neutrino mass differences, absolute masses Sharpen the test of the neutrino fluxes Sharpen the test of the neutrino fluxes-

• is the solar

is the solar neutrino problem really solved? neutrino problem really solved? Required nuclear data Required nuclear data

3He(α,γ)7Be for Borexino, Ga, Cl 7Be(p, γ)8B, (S17)

Is there a fourth neutrino (sterile neutrino)? Evidence of third Evidence of third ∆

∆m

m2

2 of about 1

• f about 1 eV

eV from LSND from LSND expt expt Made improbable by WMAP limit of 0.23 Made improbable by WMAP limit of 0.23 eV eV? ? Takes four neutrinos to explain three differences Takes four neutrinos to explain three differences Not more than 3 interacting neutrinos: width of Z Not more than 3 interacting neutrinos: width of Z boson and Big Bang, = > sterile boson and Big Bang, = > sterile Checked at FNAL in near future (mini Checked at FNAL in near future (mini-

• BOONE).

BOONE). M= 0

∆m atmos ∆m solar

How the Sun Evolves Core hydrogen burning ends

• Consumed central 10% of sun
• No heat source, pressure

decreases, gravity wins

• Core collapses, releases

gravitational energy which heats the core

Core helium burning starts

• Core hot-allows fusion of

two α’s (Z= 2)

• Helium fuses to 12C, 16O
• Hydrogen burns in shell

H-burning shell Non-burning envelop He Burning Core T= 108 K

ρ = 104 g/cm3

What’s next for the Sun? It’s the end of the line

• Helium burning ends after 108 years, C and O core
• Gravitational collapse, BUT, never reach sufficient T to fuse C + C.
• Collapse continues to 107 g/cm3--electron pressure stops collapse
• Shells still burning, unstable, blow off planetary nebula

Star becomes a white dwarf (e.g. Sirius B). Ring nebula in Lyra-NGC 6720—a

planetary nebula Property Earth Sirius B Sun Mass (M sun) 3x10-6 0.94 1.00 Radius(R sun) 0.009 0.008 1.00 Luminosity(L sun) 0.0 0.0028 1.00 Surface T (K) 287 27,000 5770 Mean ρ (g/cm3) 5.5 2.8x106 1.41 Central T (K) 4200 2.2x107 1.6x107 Central ρ (g/cm3) 9.6 3.3x107 160

The Evolutionary Process for Heavy Stars With this background can guess what happens for heavier stars

The Result

Mg

e Stellar On on

→ Non-b Fu

Start like the sun:

He burning core T=108 K

ρ =107 kg/ m3 H burning shell Non-burning envelope

When He is exhausted in the core and the core collapses, it does get hot enough to burn carbon and oxygen in core. The successive stages in the core are H → He, gravity, He

→ C,O, gravity, → C,O → Mg,

Si, gravity, Si →Fe.

Attention!! ! H He C O

Si

Iron (Fe)

He Burning Core T= 108 K

ρ= 107 kg/m3

Heavy Stars--The Stellar Onion

Non-burning H

Evolutionary Stages of a 25 Msun Star Weaver et al., 80 1.2-7.0 0-1-10 sec Explosive 3 x 1014 34.8 Millisec Core Bounce 3 x 109 8.1 Seconds Core collapse 3 x 107 4.1 1 d Si 1 x 107 2.3 0.5 y O 4 x 106 1.7 1 y Ne 2 x 105 0.93 600 y C 700 0.23 5 x 105 y He 5 0.006 7 x 106 y H

ρ (g/cm3)

T(K) x 109 Time Scale Burning Stage

Fe (Iron) is special

Core of our stellar onion is “Fe”, most tightly bound nucleus. Result of fusing two “Fe's” is heavier than two “Fe's”; costs energy to fuse them. No more fusion energy is available.

Core collapses

Collapse continues, until reach ~ nuclear density. Then nuclei repel, outer core bounces.

Outgoing shock wave forms, blows off stellar envelope "Fe" core Collapses Bounce--

Form Shock Wave

Shock moves out

Fe →p's , n's in

• uter part of Fe core

Time Supernovae Core Collapse

What’s Produced in a Supernova Model

• Evolve the Pre-SN star
• Put in a piston that gives the

right energy to the ejecta (Don’t know how explosion really works).

• Calculate what is ejected
• Calculate explosive processes

as hot shock passes.

• Example: Wallace and Weaver,
• Phys. Rep. 227,65(1993)

Find

• Elements, mass 20-50, generally

reproduced at same ratio to solar.

• Modifications by explosive processes

are small

• Ue

12C(α,γ)16O—an Important Reaction

Helium Burning- A two Stage Process

• 3 α : α + α ⇔ 8Be* + α → 12C* → 12C (gs) Rate known to ± 13%
• 12C(α, γ)16O Poorly known (20-30% )
• Ratio of rates affects 12C/ 16O after He burning—important effects Get

good agreement on previous slide) only for tuned value of the rate. Element Synthesis in SN Mass of Pre-SN Cores

Core masses

Fe core size

Updated Models-Heger and Woosley 2001 It’s more complex than the onion: M= 22 Msun. Along the x-axis

sequential episodes of convective carbon, neon (brief), oxygen, and silicon burning.

More From Updated Models—Rauscher et al. Ap.J. 576,323(2002) What’s new

• Improved nuclear physics
• Network of 2500 nuclides
• Mass loss included
• ν process included

Results

• Constant value compared to

solar up to mass 80.

• Less exceptions than before—

for example 19F, 11B (too much), 7Li made in ν process (ν’s breakup heavier nuclei)

• Constant ratio = > 1/7 stars

have such explosions

Some important Nuclear Rates for SN Synthesis

12C(α,γ)16O

Sα12 = 170 ± 20 keV-b (300 keV) describes abundances (tuned) Experiment: 100-200, preference near 150, but uncertain. High priority reaction

12C + 12C for Carbon burning 22Ne(α, γ), 22Ne (α, n)

Production of light slow neutron capture (s-process) nuclide,A= 60- 88 For more details R.Hoffmann et al. UCRL-JC- 146202 and many references at: http://www.ucolick.org/~ alex/ Weak decay rates for gamma line

• emitters. E.g. 60Fe

Charged particle reactions on N= Z nuclei for production of p-process nuclei

What About the Core Collapse—> SN ? We know that

Shock blows off outer layers of star, a supernova 1051 ergs (1foe) visible energy released (total gravitational energy

• f 1053 ergs mostly emitted as

neutrinos).

Theoretically

Spherical SN don’t explode Shock uses its energy dissociating “Fe”, stalls Later, ν’s from proto-neutron star deposit energy, restart the shock. Still no explosion.

1-D model (T. Mezzacappa) Question: Could better microphysics, e.g. better weak interaction rates

lead to explosion? Maybe, only 1% of available energy to explosion.

Is sphericity the problem?

Have 3-D calculations which explode But have only a part of the detailed microphysics. Their stability against such changes is not known. See, e.g. Fryer and M. Warren, Astrophysical Journal, 574:L65- L68 (2002) Find 2-D, 3-D similar

Two views Red upwelling Blue sinking Non-Spherical Calculations

The Question Remains—What do we need for an explosion? SN 1987a in Large Magellanic Cloud

Aside--Nature of Allowed Weak Interactions Allowed Strength-non-relativistic

Fermi

τ± = Σ τ ±(i) = > ∆L = ∆S = 0, ∆T = 0, 1

0+ 0+ (IAS dominates) Sum Rule: Σ β- - Σ β + = N-Z Gamow-Teller σ τ ± = Σ σ(i) τ ±(i) = > ∆L = 0, ∆S= 1, ∆T = 0, 1 0+ 1+ (Giant resonance GTGR) Sum Rule: Σ β- - Σ β+ = 3(N-Z)

IAS

GTGR L = 1

90Zr (p,n)

120 MeV 0 deg

Weak GT strength important for r process

Weak Strength and Supernovae Core Collapse Gamow-Teller (GT) Strength?

Mediates β-decay, electron capture(EC), n induced reactions GT (allowed) Strength S= 1; L = 0, e.g. 0+ → 1+ ; GT+,GT- Lies in giant resonances

Situation

After silicon burning, Tcore ≈3.3 x 109 K, density≈108 g/cm3. e- Fermi energy allows capture into GT+. Reduces e- pressure emits neutrinos. Speeds collapse. At higher T, GT+ thermally populated, β- decays back to ground state. β- ⇔ E.C. (n,p) (p,n)

GT+ dominates process

How Weak Strength Affects the SN Core Core size depends on Ye= < Z/A>

Starts near 0.5 Reduced by electron capture As Ye decreases, β- decay becomes important. Competition of EC and β- stabilizes Ye near 0.45 When EC and β- compete we have the possibility of a cyclic process-the URCA process.

Urca process (named after

former Casino da Urca in Rio de Janeiro that takes your money slowly but surely)

ZA + e- → Z-1A + ν Z-1A → ZA + e- + ν

Net result: production of neutrinos removes energy from the core T, entropy reduced

Effects of Changed Weak Rates-Heger et al. Ap.J. 560 (2001) 307 IPM vs Shell Model

WW standard Wallace-Weaver rates based on independent particle model LMP-from large basis shell model

• calculations. Langanke and

Martinez-Pinedo, NPA 673, 481(00) Compare results of pre-core- collapse calculations Significant differences

100 101 102 103 104 105 106 Time till collapse (s)

10-8 10-7 10-6 10-5 10-4 10-3

⏐ ⏐⎯⎯ dYe dt ⏐ ⏐ (s−1)

LMP-EC LMP-

−

0.44 0.46 0.48 0.50

Ye

WW LMP

2 4 6 8

T (109 K)

T

15 M

Si Ignition Si depleted Core contract Core collapse URCA

• 0.2
• 0.1

0.0 0.1

∆S (kB)

10 15 20 25 30 35 40

Star Mass (M )

• 0.2
• 0.1

0.0 0.1

∆ MFe ( M )

0.005 0.010 0.015 0.020

∆Y

e

Effects

Larger, lower entropy "Fe" pre- collapse core More e-'s (Ye larger), lower T core. Larger homologous core

These changes tend to make explosions easier To improve rates need to know which nuclei are important More Weak Interaction Results

Improving Weak Interaction Strengths—Pre-collapse Most important nuclei-Heger et al.

Generally closer to stability than predicted earlier. Stable and radioactive nuclei important

Can’t rely on exp’t

Need many rates Some transitions are from thermally excited states

Need

Experiments to verify accuracy Measurements for the most crucial cases, if possible Reliable calculations

53Mn 59Ni 56Fe 56Fe 53Cr 55Mn

* * * Dominant Stable

100 101 102 103 104 105 106 0.42 0.44 0.46 0.48 0.50

Ye

55Fe 57Co 57Fe 61Ni 57Fe 53Cr 57Fe 53Cr

WW LMP

15 M

Time till core collapse(sec)

Heavy Nuclei Are Also Important and Unstudied Heavier nuclei are important during collapse (A= 60-120) Previously ignored: In IPM transitions are blocked for Z< 40, N> 40. But in nature: configuration mixing and thermal excitation relieve blocking.

Langanke et al PRL

Results: Langanke , et al. find that nuclei (previously unimportant) dominate electron capture by x10 over protons Changes nature of core (Hix, et al.) : Mass behind shock (-20% ), central density (-10% ), ν luminosity (+ 15% ), average ν energy (+ 1 MeV)

Present Situation

Quite good, some problems Need data for unstable isotopes. And for heavier elements for collapse phase.

: FFN (IPM model)

: data (n,p) (TRIUMF) : Caurier et al. (1999) SM : Caurier et al. folded with experimental resolution

? ? ?

(Caurier, et al NPA 653, 439(99

First expts

Secondary triton beams 106/sec at MSU/NSCL, 120 MeV/A tritons (Daito et al). Resol: 160 keV achieved Preliminary data on 58Ni

Near Future (Zegers, et al.) 107/sec secondary beams on

CH2, 24 Mg, 50 V, 53 Cr 74 Ge, 94 Mo

Comments (t, 3He) Option

90 10 20 30 40 50 60 70 80

.0 MeV 1+ 4.5 MeV 2- 7.7 MeV 1- 160 keV

12C(t, 3He) 12B

Θ

lab=0

• ∼

1.7

• 10

20 30 40 50 60 70 80 90

• 2

2 4 6 8 10 12

0.0 MeV 1+ 4.5 MeV 2- 7.7 MeV 1- 230 keV

12C(t,3He) 12B lab

• ∼

Θ =1.7 3.4o

E(MeV)

Sherrill, et al

C

• u

n t s

Pros: Unique beam-spectrometer(S800), simple analysis, calibration

available from (3He, t) reaction at Osaka.

Con: More beam nice though not essential.

Preliminary results on 58Ni Expt’l detail or two

Et = 350 MeV, I = 106/sec. 24 hour

• run. (Future: > 10x more beam)

∆Et = 3.5 MeV, dispersion matched

spectrometer (S800) gives 0.16 MeV

Comments

The results of preliminary analysis are inconsistent with the TRIUMF (n,p) data. A similar inconsistency has been found in 85 MeV/nucleon (d, 2He) data at KVI

Radioactive Nuclei--General Approach No targets ⇒ Inverse kinematics Problems: Outgoing light particle charged

Has very low energy, won’t escape from practical target. Tritium target.

Solution:

Detect heavy particle in S800 at 0

• deg. A general solution.

3He 56Co 56Ni 3H

S800

Target Focal Plane

Superconducting Magnets Bρ = 4.2 Tm 5% ∆p, 20 msr

Expected resolution 1 MeV

Other Possibilities (d, 2He) in Inverse Kinematics

Pros: Outgoing protons allow thicker targets, maybe pure 2H (cryogenic). Can completely characterize final state. Cons: Complicated detection scheme (large Ω, good resolution); simulations (R. Zegers) look promising, not yet complete. Will be hard to get sub-MeV resolution.

(p,n) Inverse Kinematics

Good resolution with a highly granular neutron detector Well understood reaction model Test models for GT, forbidden strength T< , T0,{ T> ?} states) (7Li, 7Be) Inverse Kinematics. Complex--need to detect decay gamma from 7Be(0.431, 3/2-) to guarantee S = 1?? If so, lose efficiency. First test with 56Ni; results are being analyzed.

The r-Process What is it?

Heavy elements formed by rapid neutron capture on seed nuclei Flow along path near neutron drip line until (n,γ) = (γ,n)

N(Z) ∝ tβ

After explosion, decay back to stable region. Where does it

• ccur?

In hot bubble just inside SN shock? Or in fusion of two neutron stars? This is seeming unlikely. Hot bubble Peak

β− β− β− β− β−

Hot bubble

R Process—A Global View

Sensitivity to Nuclear Properties

Masses Half-lives Delayed Neutron Emission Ratio Waiting Points (n,γ) ⇔ (γ,n) Abundances Timescale Abundances

CERN/ISOLDE Much work with low energy beams formed by the ISOL technique.

Fast beams

A powerful way of preparing stopped nuclides for study—Greater reach toward drip line. Examine this next

R-process-Schematic Calculations Problem:

Nuclear physics or astrophysics model failure?

Sensitivity to nuclear structure. Need to measure the important

• quantities. Understand the nuclear structure:

Is there shell quenching (smaller shell gaps)?

A

relative abundance

100 120 140 160 180 200 220 10

• 4

10

• 3

10

• 2

10

• 1

10 10

1

Pfeiffer & Kratz, Mainz

solar

ETFSIQ (shell quenching) ETFSI1 (no quenching)

• The Problem

ν Processing

Generic r- Process Calculation

• 4 cm x 4 cm active area
• 1 mm thick
• 40-strip pitch in x and y

dimensions ->1600 pixels

Correlation of fragment implants and subsequent beta decays by pixelation

• f BCS

NERO PPAC Si detector (∆E)

Degrader

neutron

β

Si Si Si BCS

The Principle

Hosmer et al. (NSCL, Mainz, Maryland, Notre Dame, PPNL

First r-Process Experiments at the NSCL

117Ru

For some of these nuclei the fractions

• f β-delayed

neutron emission will be obtained-not yet analyzed A = 115 to 124, Walters, Montes et al.

117Ru, 119,120Rh, 121Pd, 124Ag have yielded

lifetimes to date. A= 75-80, Hosmer, et al Half-Lives—Preliminary Analyses

Known half-life r-process Pb Fe NSCL/MSU Reach

Exp 02028 Exp 02032 Exp 01015

Hosmer et al. Montes, et al. Walters, et al.

RIA will extend these measurements to the great majority of nuclei involved in the r-process

Approved expt

Future r-process Studies at the NSCL

New Observations--”r-process stars”

New observations (Details http://www.nhn.ou.edu/~ cowan/)

Very old metal poor stars in galactic halo, extremely rich in r-process

• elements. Same distribution as in the solar system. Implies r-process is

unique? (But distribution different for A< 130, two processes?)

Solar

U-Th clock

r-process NSCL Reach Known Half-lives

RIA Reach

• Next generation exotic beam facility in US
• In planning stage
• Would cover 80% of r-process path up to A=208

for at least a half-life measurement

RIA and the R Process

Explosive Hydrogen Burning-Accreting Binary Systems First X-ray pulsar: Cen X-3 (Giacconi et al. 1971) with UHURU First X-ray burst: 3U 1820-30 (Grindlay et al. 1976) with ANS

Today: ~ 50 Today: ~ 40 Total ~ 230 X-ray binaries known Total ~ 230 X-ray binaries known T~ 5s 10 s

Mass transfer by Roche Lobe Overflow Star expands on main sequence.

Fills its Roche Lobe Mass transfer begins through the L1 Lagrangian point

iew View

Neutron Star Donor Star (“normal” star) Accretion Disk Neutron stars: 1.4 Mo, 10 km radius (average density: ~ 1014 g/cm3)

Typical systems:

• accretion rate 10-8/10-10 Mo/yr (0.5-50 kg/s/cm2)
• orbital periods 0.01-100 days
• orbital separations 0.001-1 AU’s

An Artist’s View

Energy generation: thermonuclear energy Ratio gravitation/thermonuclear ~ 15-30 4H 4He 5 4He + 84 H 104Pd 6.7 MeV/u 0.6 MeV/u 6.9 MeV/u Energy generation: gravitational energy

E = G M mu R = 200 MeV/u

3 4He 12C (“triple alpha”) (rp process) Nuclear vs. Gravitational Energy

Observation of thermonuclear energy

Unstable, explosive burning in bursts (release over short time) Burst energy thermonuclear Persistent flux gravitational energy

Visualizing reaction network solutions

Proton number

27Si

neutron number

13 14 (p,γ) (α,p) (a, γ) (β+) Lines = Flow =

dt dt dY dt dY F

i j j j i i j i

∫

⎥ ⎦ ⎤ ⎢ ⎣ ⎡ − =

> − > − ,

0 1 2 3 4 5 6 7 8 9 10 111213 14 1516 17181920 2122 2324 25262728 2930 3132 33343536 37383940 41 424344 45464748 4950 5152 53 5455 56 5758 59 H (1) He (2) Li (3) Be (4) B (5) C (6) N (7) O (8) F (9) Ne (10) Na (11) Mg (12) Al (13) Si (14) P (15) S (16) Cl (17) Ar (18) K (19) Ca (20) Sc (21) Ti (22) V (23) Cr (24) Mn (25) Fe (26) Co (27) Ni (28) Cu (29) Zn (30) Ga (31) Ge (32) As (33) Se (34) Br (35) Kr (36) Rb (37) Sr (38) Y (39) Zr (40) Nb (41) Mo (42) Tc (43) Ru (44) Rh (45) Pd (46) Ag (47) Cd (48) In (49) Sn (50) Sb (51) Te (52) I (53) Xe (54)

3α reaction

α+α+α

12C

αp process:

14O+α 17F+p 17F+p 18Ne 18Ne+α …

rp process:

41Sc+p 42Ti

+p 43V +p 44Cr

44Cr 44V+e++νe 44V+p …

Most calculations (for example Taam 1996) Wallace and Woosley 1981 Hanawa et al. 1981 Koike et al. 1998

Schatz et al. 2001 (M. Ouellette) Phys. Rev. Lett. 68 (2001) 3471

Models: Typical reaction flows

Schatz et al. 1998

X-ray burst: Importance of waiting points

time (s)

300 400 500 600 10

• 3

10

• 2

10

• 1

10 300 400 500 600 10

• 5

10

• 4

10

• 3

10

• 2

10

• 1

300 400 500 600 0e+ 00 5e+ 16 1e+ 17

luminosity(erg/g/s)

cycle fuel abundance

1H 4He

abundance

56Ni 72Kr 104Sn 64Ge 68Se

• Luminosity:
• Abundances of

waiting points

• H, He abundance

Definition: Flow is hampered by slow decay or weakly bound nuclei

0 1 2 3 4 5 6 7 8 9 10 111213 14 1516 17181920 2122 2324 25262728 2930 3132 33343536 37383940 41 424344 45464748 4950 5152 53 5455 56 5758 59 H (1) He (2) Li (3) Be (4) B (5) C (6) N (7) O (8) F (9) Ne (10) Na (11) Mg (12) Al (13) Si (14) P (15) S (16) Cl (17) Ar (18) K (19) Ca (20) Sc (21) Ti (22) V (23) Cr (24) Mn (25) Fe (26) Co (27) Ni (28) Cu (29) Zn (30) Ga (31) Ge (32) As (33) Se (34) Br (35) Kr (36) Rb (37) Sr (38) Y (39) Zr (40) Nb (41) Mo (42) Tc (43) Ru (44) Rh (45) Pd (46) Ag (47) Cd (48) In (49) Sn (50) Sb (51) Te (52) I (53) Xe (54)

The Sn-Sb-Te cycle

104Sb 105Sb 106 107Sb 103Sn 104Sn 105Sn 106Sn 105Te 106Te 107Te 108Te 102In 103In 104In 105In

(γ,a)

Sb

β+ (p, ) γ

Known ground state α emitter

Endpoint: Limiting factor I – SnSbTe Cycle

Nuclear reactions on accreting neutron stars--H. Schatz

Thermonuclear burning (rp process)

• Why do burst durations vary ? (10s – min)

Neutron Star Surface fuel ashes

• cean

inner crust

• uter

crust H,He atmosphere

Electron captures Pycnonuclear reactions

Galactic nucleosynthesis contribution ?

• Gravitational wave emission ?
• Crust heating ?
• Dissipation of magnetic fields ?

Start composition for deeper processes ?

Deep H, C, … burning

• Origin of Superbursts ? 100X stronger
• What nuclei are made in the explosion ?

Need nuclear physics to answer and to understand observations

56Ni 56Ar

and so on …

(1.5 x 1012 g/cm3)

Electron capture and n-emission Electron capture and n-emission

Pyconuclear fusion

Ni (28) Fe (26) Cr (24) Ti (22) Ca (20) Ar (18) S (16) Si (14) Mg (12) Ne

(2.5 x 1011 g/cm3) (1.5 x 109 g/cm3) (1.5 x 109 g/cm3)

34Ne

Electron capture Electron capture

border of known masses

68Ca

NSCL Reach Crust reactions in accreting neutron stars

Need to measure:

• Masses (Exp 1035 Santi, Ouellette at S800)
• Electron capture rates (Exp 1038 Sherrill at S800)

(Charge exchange in inverse kinematics (7Li,7Be) )

From Haensel & Zdunik 1990

H,He αp/rp process αp/rp process

S-Factors at High Energies Two Topics Briefly

ANCs and S-Factors Coulomb Breakup measurements (more detailed)

ANCs and S-Factors Measure ANC ⇒ S(E = 0) for (p, γ), (a, γ) reactions Principle:

• Low-E (x,γ) reactions occur far

from the nuclear surface

• σ ∝ |ψ(large r)| 2 ∝ ANC2

Important region r

Experiments: Transfer reactions at low energies measure ANC

Detailed work: Texas A&M(Ajhari, Gagliagardi, Mukhamedzhanov, Tribble, et al.) 7Be(p, γ)8B, 13C(p, γ), 16O(p, γ) . Issues: Require accurate OM Potentials, limits accuracy to about 10% ; checked to 10% against 16O(p, γ) Example 10B(7Be,8B)9Be, 14N(7Be,8B)13C at 85 MeV ⇒ S(7Be(p, γ)) to 10%

S factor for 16O(p, γ)17F — A test of the ANC Method

Prediction by ANC

Test case-known from Direct Capture ANC’s for16O(3He,d)17F (C2)gnd = 1.08 ± .10 fm-1 (C2)ex = 6490 ± 680 fm-1 Direct Capture data from Morlock, et al. Agree within the relative errors- the 10% level

Coulomb Breakup Studies for Astrophysics

Obtain σ (α,γ) from σ for breakup by detailed balance

k 2 γ ) ( ) ( 2 b a c k c b a + → + = + → + γ σ γ σ c

Advantages: Much larger cross section, convenient for (n,γ),

experiments are relatively straightforward, not simple

Issues: Require detailed theory to justify approach. Issues (E2, nuclear

contributions) seem under control.

Some Experiments: Motobayashi, et al.: 13N(p,γ)14O, 7Be(p, γ)8B,

breakup of 8B, 14O GSI, NSCL: 7Be(p, γ)8B, 8Li(n, γ)9Li, 3He(α,γ)7B

Results from 8B(g,p)7Be

• B. Davids et al. Phys. Rev. Lett. 86, 2750 (2001)

S17 = 16.8±1.1 eV-barn, PRC 68 045802 (03) E1, E7, θrel ⇒ Erel

The Problem of Extrapolation Structure uncertain ⇒ extrapolation uncertain to 2 eV-barn at the

extremes

Choices for extrapolation

Junghans et al use Descouvement-Baye, Assign 0.6 eV-b error Davids-Typel fit n-7Li and p-7Be scattering lengths 0.9 eV-b error

Davids-Typel PRC 68, 045802 (2003) S error chi2/n-1 S error chi2/n-1 All data 20.3 0.4 3.3 19.4 0.3 3.3 Direct 21.4 0.5 1.2 19.9 0.4 3.0 Indirect 18.0 0.7 0.9 17.5 0.7 0.8 Junghans, et al .Nuc-ex/0308003

Astrophysical Models Nuclear Physics Experiments Astronomical Observations

Associated:

• SciDAC SSC

LLNL, UCSC, LANL

• U. of Arizona
• UC Santa Barbara
• ANL

Nuclear Theory

Core institutions:

• Notre Dame
• MSU/NSCL
• U. of Chicago
• Identify and address the critical open questions and needs of the field
• Form an intellectual center for the field
• Overcome boundaries between astrophysics and nuclear physics

and between theory and experiment

• Attract and educate young people

www.jinaweb.org Joint Institute for Nuclear Astrophysics (JINA)

a new NSF Physics Frontiers Center

Virtual Journal of Nuclear Astrophysics

Articles from 30 journals, appears weekly: www.nscl.msu.edu/~ jina/VJ/

Biochemistry NSCL Chemistry Biomedical & Physical Sciences

National Superconducting Cyclotron Laboratory

RPMS n-Ball S800 20 msr

∆E/E = 10-4

Test 4Pi Scat. Cham. Sweeper Trap

Experimental Areas

A1900 Properties

• Bρ = 6 Tesla-m
• ∆p/p = ±2.5%
• dΩ = 8 msr

Beam Production at NSCL/CCP

Dipole Dipole Properties

• dΩ = 20 msr (7ox10o)
• ∆E/E = 10-4

NSCL S800 Spectrometer

Fast Beams from the NSCL CCP