nuclear physics (and astrophysics) of core collapse...

W.R. Hix (ORNL/ U. Tenn.) MSU Astrophysics Seminar, March 2008

Nuclear Physics (and Astrophysics) of Core Collapse Supernovae

CollaboratorsM. Baird (UTK), E. Lentz (UTK), O.E. B. Messer (ORNL), A. Mezzacappa (ORNL)K. Langanke(GSI), G. Martínez-Pinedo (GSI), A. Juodogalvis(Vilnius), J. Sampaio(Lisbon)H.-T. Janka (MPA), B. Müller (MPA), A. Marek (MPA)S.W. Bruenn (FAU), J. Blondin(NCSU), C.-T. Lee(UTK)


Microscopic ⇔ MacroscopicCore Collapse Supernovae present an interesting Nuclear Physics problem in that physics at the femtometer scale drives (and is driven by) astrophysical processes at the kilometer scale,

leading to explosions that spread debris over exameter scales.


Gain Radius

Heating

Cooling

ν-LuminosityMatter Flow

Proto-NeutronStar

ν-Spheres

νe + n ← p + e-

νe + p ← n + e+_

νe + n → p + e-

νe + p → n + e+_

Shock

Textbook Supernova


Supernova Nucleosynthesis


Spherically symmetric collapse, bounce and shock stall simulations using AGILE-BOLTZTRAN Fully implicit, multi-group, 4-flavor Boltzmann neutrino transport

• Most modern neutrino physics, LMSH electron capture rates (Hix et al. 2004, Langanke et al. 2003) or Bruenn 85 electron capture rates (Bruenn 1985)

• Implicit, spherically symmetric hydrodynamics with an adaptive mesh.

• Both general relativistic (including gravitational redshift)and Newtonian gravity.

• Modular architecture allows use of multiple realistic equations of state.

Examining Microphysics Effects

Liebendörfer, Messer, Mezzacappa, Bruenn, Cardall & Thielemann (2004)


History of Captures on NucleiEntropy of iron core is low (S/k ~1) so few free nucleons are present. Thus e- and ν capture on heavy nuclei via 1f7/2⇔1f5/2 GT

transition dominates. (Bethe,Brown, Applegate & Lattimer 1979)

During collapse, average mass of nuclei increases, quenching e- capture (at N=40).

Thermal unblocking and first forbidden were considered but rates too small. (Fuller 1982, Cooperstein & Wambach 1984)

Implemented using average nucleus. Bruenn (1985)


New e-/ν Capture Rates

Shell Model calculations are currently limited to A~65.

Langanke et al (2003) employed a hybrid of shell model (SMMC) and RPA to calculate a scattering of rates for A<110.

Langanke, …, Hix, … (2003)


New e-/ν Capture Rates

Shell Model calculations are currently limited to A~65.

Langanke et al (2003) employed a hybrid of shell model (SMMC) and RPA to calculate a scattering of rates for A<110.

Electron/neutrino capture on heavy nuclei remains important throughout collapse.

Langanke, …, Hix, … (2003)


The impact of EC with stellar mass

Higher mass cores have higher initial entropy.Effects of nuclear electron capture are reduced

but comparable (1/2 to 2/3).

15 M 25 M


PNS Convection

Fluid instabilities which drive convection result from complete neutrino radiation-hydrodynamic problem including nuclear interactions.

Hix, Messer, Mezzacappa, … 2003

Updated nuclear e-/ν capture restricts PNS convection to smaller, deeper region.


Changes in Neutrino Emission

νe burst slightly delayed and prolonged.

Other luminosities minimally affected (~1%).



Changes in Neutrino Emission

νe burst slightly delayed and prolonged.

Other luminosities minimally affected (~1%).

Mean ν Energy altered:

1-2 MeV during collapse

~1 MeV up to 50ms after bounce

~.3 MeV at late time



1010 1011 1012 1013 1014

Density (g/cm3)

0.25

0.3

0.35

0.4

Yl,Y

e

Np*Nh=0.1Np*Nh=1.0Np*Nh=10.

0.05 Mo.

Ye

Yl

What Rates are needed?Change in lepton abundance (Yl=Ye+Yν) occurs gradually up to ~3x1012 g/cm3.

By 3x1012, rate of electron capture is determined largely by blocking.


1010 1011 1012 1013 1014

Density (g/cm3)

0.25

0.3

0.35

0.4

Yl,Y

e


0.05 Mo.

Ye

Yl

1010 1011 1012 1013 1014

density [g/cm3]

0

50

100

150

200

A, Z

A

Z



Average Nuclear Mass by 1012 is 100 or more with many nuclei contributing.


1010 1011 1012 1013 1014

Density (g/cm3)

0.25

0.3

0.35

0.4

Yl,Y

e


0.05 Mo.

Ye

Yl

1010 1011 1012 1013 1014

density [g/cm3]

0

50

100

150

200

A, Z

A

Z



Need theory for the large number of reactions of interest and experiments to constrain this theory.

Average Nuclear Mass by 1012 is 100 or more with many nuclei contributing.


What can RIBs say about e-/ν Capture?Charge Exchange Reactions, e.g. (n,p),(d,2He),(t,3He),also sample GT+ strength distribution, providing strong constraints on structure models.

Baümer et al. PRC 68, 031303 (2003)

For A=80-100 nuclei of interest are 2-6 neutrons richer than stability.

Current Experiments, on stable nuclei, agree well with shell model calculations for A<60.

Should be achievable with NextGen RIBs.


Advertisement New tabulation of nuclear electron capture coming soon.

Improvements include:✦ Expanded range of nuclei in NSE✦ Partition functions exceeding 10 GK (Rauscher 2000)✦ Updated screening prescription in NSE✦ Screening of weak reactions✦ Expanded set of SMMC/RPA rates (250)✦ Finely tabulated spectra (.5 MeV)

Sampaio, Juodogalvis, Langanke, Martínez-Pinedo & Hix (2008) ADNDT, in prep.


Inelastic Neutrino-Nucleus Scattering As with electron/neutrino capture, advances in nuclear structure physics are improving our understanding of other neutrino-nucleus interactions.

While coherent, elastic ν-nucleus scattering has long been considered, inelastic ν-nucleus scattering (INNS) has often been ignored. The exception is Bruenn & Haxton (1991) which used INNS rate calculated for 56Fe at T=0.

Recently, Juodogalvis, Langake, Martinez-Pinedo, Hix, Dean & Sampaio (2005) calculated INNS for 40 isotopes of Mn, Fe, Co & Ni at finite temperature using a combination of shell model and RPA. A tabulation of NSE-averaged rates has been produced for use in supernova simulations.


Dynamical Effects of INNS?

After bounce, heating rate just above shock is boosted (2-3x). However, heating of supersonically infalling matter is ineffective, so dynamics are little effected.

During collapse, INNS works like NES to equilibrate the neutrino distribution. However there is little effect from this addition.

Müller, Janka (2008), priv. comm.


Observable Effects of INNS

Higher energy tail of ν spectra are strongly suppressed during this interval.

Langanke, Martínez-Pinedo, Müller, Janka, Marek, Hix, Juodagalvis & Sampaio (2007)

Has a surprisingly large effect on potential terrestrial ν detectors.


Summary for Nuclear Opacities Electron and neutrino captures on heavy nuclei play an important role during core collapse, affecting✦ the initial launching point of the shock by up to

20% in mass.✦ the rate of collapse of the outer iron core✦ material gradients in the proto-neutron star.

Inelastic scattering of neutrinos off heavy nuclei can have a large impact on the local heating rate just above the shock, however, the dynamic impact is small.A much larger impact is felt by the high energy tail of the supernova neutrino spectra, potentially impacting terrestrial detectors.


Role of the Equation of State In general, the equation of state closes the system of hydrodynamic equations by relating the pressure to the internal energy (or temperature or entropy), density and composition.

For supernovae, it includes contributions from photons, degenerate electrons & positrons, and nuclei or nuclear matter.

In regions where Nuclear Statistical Equilibrium can be assumed, the supernova EOS also provides the nuclear composition, usually in the form of mass fractions for free nucleons, alpha particles and a representative heavy nucleus.


Examining Equations of State ✦ LS EoS (Lattimer-Swesty 1991)

• Liquid drop model• Compressibility – 180 MeV• Complete EoS - Includes baryon, lepton and photon contributions• Alpha abundance error

✦ STOS EOS (Shen, Toki, Oyamatsu & Sumiyoshi 1998)• Relativistic mean field theory• Compressibility – 281 MeV• Tabular Nuclear EoS - Includes only nuclear contributions

✦ Wilson EOS (Mayle & Wilson 1991, McAbee & Wilson 1994)• Empirical Relation of Baron, Cooperstein & Kahana (1985), constrained

by relativistic Brueckner-Hartree-Fock calculations of Muther, Prakash & Ainsworth (1987).

• Compressibility – 200 MeV • Includes Pions at high density constrained by comparison of simulations

of Pion production in heavy ion collisions with experimental data.• Tabular Complete EoS


Differences limited to nuclear matter. At lower density, electrons dominate pressure.

Pressure

for fixed thermodynamic conditions near bounce

Baird, Lentz, Hix, Messer & Mezzacappa, in prep.


Moderate differences, with STOS having higher internal energies in the region of forming shock.

Internal energy




Entropy

Noticable differences at center, even larger differences at lower density.Entropy is dominated by the nuclear contribution.




Composition of unshocked matter shows large differences.

Composition



Composition after bounce

Shocked compositions similar except Wilson near phase transition



Self-consistent Simulations

Wilson: Shock launches strongly, peaking near 150 ms with largest shock stall radius.

0 50 100 150 200101

102

Time After Bounce [s]

Radi

us [k

m]

Examined the impact of the EOS using Agile-Boltztran GR simulations of 15 M progenitor.

Differences persist in spherically symmetric models for more than 200 ms after bounce.

STOS: Shock launches slowly and peaks near 50 ms with smallest shock stall radius.

LS

STOS

Wilson



0.25

0.35

0.45

0.55Y

e (

Ele

ctro

n F

ract

ion

)

1

2

3

4

5

En

tro

py

pe

r b

ary

on

108

1010

1012

1014

De

nsi

ty (

g c

m!3)

STOSLSWilson

0 0.2 0.4 0.6 0.8 1 1.2 1.4 1.6!1

0

1x 10

5

Ve

loci

ty (

km/s

)

Enclosed Mass

EOS Comparison w/Nuclear EC

At bounce, Wilson shows larger core and faster collapsing outer core.



0.25

0.35

0.45

0.55Y

e (

Ele

ctro

n F

ract

ion

)

1

2

3

4

5

En

tro

py

pe

r b

ary

on

108

1010

1012

1014

De

nsi

ty (

g c

m!3)

STOSLSWilson

0 0.2 0.4 0.6 0.8 1 1.2 1.4 1.6!1

0

1x 10

5

Ve

loci

ty (

km/s

)

Enclosed Mass

0.1

0.2

0.3

0.4

0.5

0.6Y

e (

Ele

ctro

n F

ract

ion

)

02468

101214

En

tro

py

pe

r b

ary

on

108

1010

1012

1014

De

nsi

ty (

g c

m!3)

L!S

STOS

Wilson

0 0.2 0.4 0.6 0.8 1 1.2 1.4 1.6!1

0

1x 10

5

Ve

loci

ty (

km/s

)

Enclosed Mass

100 ms after bounce

EOS Comparison w/Nuclear EC

At bounce, Wilson shows larger core and faster collapsing outer core.After bounce, Wilson and STOS show different Ye gradients.



0.2

0.3

0.4

0.5

Ye (

Ele

ctro

n F

ract

ion)

0

2

4

6

Ent

ropy

per

bar

yon

108

1010

1012

1014

Den

sity

(g

cm!3)

STOSLSWilson

0 0.2 0.4 0.6 0.8 1 1.2 1.4 1.6!10

!5

0

5x 10

4

Vel

ocity

(km

/s)

Enclosed Mass

EOS Comparison w/proton EC

At bounce, shock location similar, but core Ye gradient stronger.



0.2

0.3

0.4

0.5

Ye (

Ele

ctro

n F

ract

ion)

0

2

4

6

Ent

ropy

per

bar

yon

108

1010

1012

1014

Den

sity

(g

cm!3)

STOSLSWilson

0 0.2 0.4 0.6 0.8 1 1.2 1.4 1.6!10

!5

0

5x 10

4

Vel

ocity

(km

/s)

Enclosed Mass

EOS Comparison w/proton EC

At bounce, shock location similar, but core Ye gradient stronger.After bounce, Wilson and STOS still show different Ye gradients.

0.1

0.2

0.3

0.4

0.5

0.6

Ye (

Ele

ctro

n F

ract

ion)

0

2

4

6

8

10

12

Ent

ropy

per

bar

yon

108

1010

1012

1014

Den

sity

(g

cm!3)

L!S

STOS

WIlson

0 0.2 0.4 0.6 0.8 1 1.2 1.4 1.6!1

0

1x 10

5

Vel

ocity

(km

/s)

Enclosed Mass

100 ms after bounce



A Closer look at bounce with Wilson EoS…Lower Pressure results in higher bounce density.

Ye decreases rapidly in the central core as Pions become important.

Pions reduce proton chemical potential shifting detailed balance away from electrons and toward neutrinos. Baird, Lentz, Hix, Messer & Mezzacappa, in prep.


Interplay of EOS and Nuclear Electron Capture

Composition provided by the EOS impacts relative strength of opacities. Thus statements about importance of opacities must be re-examined in light of progress on the EOS.



EOS effects Equations of State play a large role in collapse and drives the physics involved in post bounce evolution.

In spite of relatively small differences in the nuclear matter regime:•Initial proto-neutron star size different by

almost 20%•20% differences in shock location after 200 ms•Altered lepton and entropy gradients

Effects of EOS not limited to pressure but also include composition and therefore opacities.


CHIMERA Building Blocks• “RbR-Plus” MGFLD Neutrino Transport

• PPM Hydrodynamics • Lattimer-Swesty EOS• Nuclear (α) Network• Newtonian Gravity with Spherical GR Corrections

• Neutrino Emissivities/OpacitiesStandard (Bruenn 1985)Elastic Scattering on NucleonsNucleon–Nucleon Bremsstrahlung

• Run for postbounce times > 400 ms.• Run on a 180 degree grid.

Multi-Physics Supernova Models

Ray-by-Ray ApproximationRadial transport allowed.

Lateral transport suppressed.

Cast:Bruenn, Blondin, Hix, Marronetti, Messer, Mezzacappa, Lentz, Budiardja, Lee, Yakunin, …


Instability behind the shock distorts and expands the shock, increasing neutrino heating. Explosion takes several hundred ms after bounce to develop.

20 M GRResolutionr,θ: 256, 256ν energy: 20

2D Results


Anatomy of Explosion


Time (s)

Avg Atomic Mass

Radi

us (

km)



Si Layer

O Layer

IronCore


Time (s)

Avg Atomic Mass

Radi

us (

km)



ShockSi Layer

O Layer

IronCore


Time (s)

Avg Atomic Mass

Radi

us (

km)


Effects of Nuclear Burning?



Effects of Nuclear Burning?

Heat near the shock from O/Si burning helps boosts shock.



Blondin, Mezzacappa & DeMarino (2003)

Shock wave unstable to nonradial perturbations.

Stationary Accretion Shock Instability (SASI)

Need for 3D?

• SASI has axisymmetric and nonaxisymmetric modes that are both linearly unstable!

– Blondin and Mezzacappa, Ap.J. 642, 401 (2006)– Blondin and Shaw, Ap.J. 656, 366 (2007)


Blondin, Mezzacappa & DeMarino (2003)

Power in l=1 mode.

Power in m=1 mode.

Shock wave unstable to nonradial perturbations.

Stationary Accretion Shock Instability (SASI)

Need for 3D?

• SASI has axisymmetric and nonaxisymmetric modes that are both linearly unstable!

– Blondin and Mezzacappa, Ap.J. 642, 401 (2006)– Blondin and Shaw, Ap.J. 656, 366 (2007)


Summary for RBR models Multi-group transport, increased run times, improved neutrino emissivities & opacities, and the inclusion of nuclear burning(?) are all critical.

With self-consistent explosions, we can examine use nucleosynthesis to further constrain models.

3D may be necessary as hydro studies suggest that the outcomes in 3D may be fundamentally different than 2D.

Impact of magnetic fields remains to be seen.


Summary 1) Nuclear Physics provides vital ingredients, in

the form of neutrino opacities, the equation of state and reaction rates relevant for nucleosynthesis, for supernova simulations.

2) The interrelation of these microphysical ingredients, both with each other, and the macroscopic astrophysical environment makes ongoing examinations important.

3) We are entering a very interesting period in core collapse supernova simulation.

nuclear physics (and astrophysics) of core collapse...

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