The Death of Massive Stars and the Birth of the Elements Stan Woosley Cornell, June 2, 2006

The Death of Massive Stars and the Birth of the Elements

Stan Woosley

Cornell, June 2, 2006

Stars are gravitationally confined thermonuclear reactors.

3T

Fe

Si, S, Ar, CaO, Mg, Ne

He, C

H, HeHe

900 R 1 R

0.1 R 0.01 R

25 M Presupernova Star

With each progressive burningstage the central entropy decreases.Red giant formation leads to an increased entropy in the outer hydrogen envelope.

Entropy

S/NAk

The key role of enropy…

Fe

O, Ne, Mg, Si

HeH

SiO

S Ne

Mg

C

C

C

collapses to a neutron star

bounce = 5.5 x 1014 g cm-3

20 Solar Masses

Mayle and Wilson (1988)

1 0 ms2 0.53 1.04 3.05 100 ms6 230 ms

Some of the contributions of Hans Bethe to our understandingof core collapse supernovae:

• The idea of a cold, low entropy bounce at nuclear density – Bethe, Brown, Applegate and Lattimer (Nucl. Phys. A , 324, 487, (1979))

2

( )

( ) ~ / 9

G exp akT

E akT a A

Fowler, Englebrecht, and Woosley(1978) not withstanding

• The prompt shock model and its failure - Baron, Brown, Coopersteinand Kahana (PRL, 59, 736, (1987))

works if iron core mass is < 1.1 solar masses (but it isn’t)

simplified models for the nuclear equation of state

• Delayed neutrino-powered explosions –

• Analytic Models for Supernovae.

“entropy” “gain radius”“net ram”“foe” – now the “Bethe” = B

Bethe, RMP, 62, 801, (1990) ApJ, 412, 192, (1993) ApJ, 419, 197, (1993) ApJ, 449, 714 (1995) ApJ, 469, 192 (1996) ApJ, 473, 343 (1996) Nuc. Phys A, 606, 195 (1996) ApJ, 490, 765, (1997)

Bethe and Wilson, ApJ, 295, 14, (1985) - see also Wilson (1982)

Wilson discovered in 1982 that neutrino energy deposition on a time scale longer than previously suspected could re-energize the shock. Hans joined with Jim in the first refereed publication to state this..

Infall

Accretion Shock;Ram pressure

radius

Velocity

13000 skm

gain radiusN

eutr

inos

pher

e

Inside the shock, matter is in approximate hydrostatic equilibrium.Inside the gain radius there is net energy loss to neutrinos. Outsidethere is net energy gain from neutrino deposition. At any one time thereis about 0.1 solar masses in the gain region absorbing a few percentof the neutrino luminosity.

Energy deposition here drives convectionBethe, (1990), RMP, 62, 801

(see also Burrows, Arnett, Wilson, Epstein, ...)

NeutronStar

TSI/ORNL

LANLMPI

Univ.Arizona

No one has yet done a 3-dimensional simulation of the full stellar core includingneutrino transport that the community would agree is “good”.

But they are getting there and the attemptsare providing insights.

Scheck, Janka, et al (2006)

Scheck et al. (2004)

Stationary Accretion Shock Instability (SASI)

Supernova shock wave will become unstable.Instability will

1. help drive explosion,2. lead to gross asphericities.

New ingredient in the explosion mechanism.

Confirmed by: Scheck et al. 2004 Janka et al. 2005 Ohnishi et al. 2006 Burrows et al. 2006

Blondin, Mezzacappa, and DeMarino (2003)

Buras et al. (2003) Physics

Parameterized neutrino heating/cooling.

Livne et al. (2004) Physics

see also Foglizzo (2001,2002)

SASI induced flow is remarkably self similar, with an aspect ratio ~2 that is consistent with supernova spectropolarimetry data.

Blondin, Mezzacappa, and DeMarino (2003)

Burrows et al. (2006) find considerableenergy input from neutron star vibrations – enough even to explode the star, andsurely enough to influence the r-process

And so – maybe – most massive stars blow up the way Hans and others talked about:

Rotation and magnetic fields unimportant in the explosion (but might be important after an explosion is launched)

Kicks and polarization from “spontaenous symmetry breaking” in conditions that started spherical.

Just need better codes on bigger faster computers to see it all work.

But ….

Dana Berry (Skyworks) and SEW

Need iron core rotation at death to correspond to a pulsar of < 5 ms period if rotation and B-fields are to matter at all. Need a period of ~ 1 ms to make GRBs. This is much faster than observed in common pulsars.

52 2 2ro

12 5

t

2 2 -1

E ~ 2 10 (1 ms/P) (R/10 km) erg

j

Total rotational kinetic energy for a neutr

~ (1ms/P) (R/10 km)

on st

cm s at M 1.4M

a

0

r

6.3 1R

For the last stable orbit around a black hole in the collapsar model (i.e., the minimum j to make a disk)

1 2 -1

2

6

11 -6

2 3 / / 3 cm s non-rotatin4.6 10

1

g

2 / 3 / / 3 cm s Kerr a = .5 110

LSO BH

LSO BH

j GM c M M

j GM c M M

Heger, Woosley, & Spruit (2004)using magnetic torques as derived inSpruit (2002)

Stellar evolution including approximate magnetic torques gives slow rotation for common supernova progenitors.

Still faster rotation at death ispossible for stars born with unusually fast rotation – Woosley & Heger (2006)Yoon & Langer (2005)

The spin rates calculated for the lighter (more common)supernovae are consistent with what is estimated for young pulsars

So, one could put together a consistent picture …

ROTATION

Ordinary SN IIp - SN Ib/c - GRB

“slow” pulsar 90%(?)

fast pulsar/magnetar? 10%(?)

millisecond magnetar oraccreting black hole (< 1%)

Cas-A - ChandraGRB – Dana Berry - Skyworks

Nucleosynthesisin Massive Stars:

Work with Alex Heger (LANL) and Rob Hoffman (LLNL)

Survey - Solar metallicity:

• Composition – Lodders (2003); Asplund, Grevesse, & Sauval (2004)

• 32 stars of mass 12, 13, 14, 15, 16, 17, 18, 19, 20, 21, 22 23, 24, 25, 26, 27, 28, 29, 30, 31, 32, 33 35, 40, 45, 50, 55, 60, 70, 80, 100, 120 solar masses. More to follow.

• Evolved from main sequence through explosion with two choices of mass cut (S/NAkT = 4 and Fe-core) and two explosion energies (1.2 B, 2.4 B) – 128 supernova models

• Use the Kepler implicit hydrodynamics code (1D)

• Use best current information on nuclear reaction rates and opacities

• Include best current estimates of mass loss at all stages of the evolution

(Weaver, Zimmerman, and Woosley 1978; see RMP, 74, 1015, (2002) for description of physics)

• Use recently revised solar abundances

Lodders, ApJ, 591, 1220 (2003)Asplund, Grevesse, & Sauval, ASP Conf Series, (2004)

Mayle and Wilson (1988)

The explosion can be characterized by a pistonwhose location and speedare free parameters.

The piston location is constrained by:

• Nucleosynthesis• Neutron star masses

The piston energy isconstrained by:• Light curves• Fall back

e

15M

Electron mole number

Y

• The Explosion Model

3 4

Except near the "mass cut", the shock temperature to which the

explosive nucleosynthesis is most sensitive is given very well by

4

Explosion ene

3

y

rg

R aT

51 10 erg 1 B

The edge of the iron core sets a lower bound to the mass cut. Otherwise, too many neutron-rich isotopes …

The location where the entropyS/NAkT = 4, typically at the baseof the oxygen shell sets an upperlimit. Stars that explode in realsimulations typically develop their mass cut here. A larger valuegives neutron stars that are too massive.

Density (g cm

-3)

Ent

rop

y pe

r ba

ryon

(S

/NAkT

)

20M presupernova

Density

Entropy

(after fall back)

1.2 B explosions; mass cut at Fe core

Above 35 M

black holes form

in Z=0 stars

1.2 B explosions

Thorsett and Chakrabarty, (1999), ApJ, 512, 288

Vertical line is at

1.35 0.04 M

Ransom et al., Science, 307, 892, (2005) find compelling evidence for a 1.68 solar mass neutron star in Terzian 5

If in the models the mass cut is taken at the edge of the iron corethe average gravitational mass for for stars in the 10 – 21 solar massrange is (12 models; above this blackholes start to form by fall back):

1.38 0.16 M

1.45 0.18M

If one instead uses the S = 4criterion, the average from 10 –21 solar masses is

From 10 to 27 solar masses theaverage is

1.53 0.22 M

15 M

2.4 B

1.2 B

2.4 B

1.2 B

25 M

1.2 B of kinetic energy at infinity gives good light curvesin agreement with observations.

2.4 B gives too bright a supernova making Type IIalmost as brilliant as Type Ia.

Though not shown here 0.6 B would give quite faintsupernovae, usually with very weak “tails”.

Isotopic yields for 31 stars averaged over a Salpeter IMF, = -1.35

Intermediate mass elements(23< A < 60) and s-process (A = 60 – 90) well produced.

Carbon and Oxygen over-produced.

p-process deficient by a factor of ~2 for A > 130 and absent for A < 130

Conclusions

• Overall good agreement with solar abundances – see also WW95.

• Lightest neutron star 1.16 solar masses; average 1.4 solar masses. Black holes a likely product for some current generation stars in the 30 – 50 solar mass range (more black holes at metallicities lower than the sun)

• Overproduction of C and O suggests that current estimates of Wolf-Rayet mass loss rates may be too large (and/or Lodders (2003) abundances for C and O too small).

Two Mysteries

• The nature of the r-process site

• The origin of the p-process 90 < A < 130

Anti-neutrinos are "hotter" thanthe neutrinos, thus weak equilibriumimplies an appreciable neutron excess,typically 60% neutrons, 40% protons

favoredat late times

r-Process Site: The Neutrino-powered Wind *

Duncan, Shapiro, & Wasserman (1986), ApJ, 309, 141Woosley et al. (1994), ApJ, 433, 229

T9 = 5 – 10T9 = 3 - 5

T9 = 1 - 2

Results sensitive to the (radiation)entropy, T3/and thereforeto aT4/, the energy density

Nucleonic wind1 – 10 seconds

- , . .,

,

- ( ) 3

5

The r process is favored by high entropy i e low

density at a given temperature because the reactions that

assemble particles to heavy seed nuclei between

and billion K increase rapidly wit

7 12 2

( , ) ( , )

-

.

h density

n Be n C

Keeping the density low thus keeps most of the mass

in particles and thus increases the ratio of free neutrons

to heavy seed

Integrated abundances in the late time wind resemble ther-process abundance pattern.

But,

The entropy in these calculations by Wilson was not replicatedin subsequent analyses which gave s/kB about 4 times smallers/kB ~ 80 not 300.

Woosley et al (1994)

Throughout nucleosynthesis epoch

Early on (t about 1/2 second)

Late (t >1 s)

L L

Wilson (1994)

With nearly equal fluxes of neutrinos,each having about the same energy,and with the lifting of degeneracythe neutron-proton mass differencefavors protons in weak equilibrium.Later, the neutrino energy difference favors neutrons.

QIAN AND WOOSLEY (1996)

t ~ 10 s

t < ½ s

data from Wilson (1994)

Janka, Buras, and Rampp (2003) 15 solar mass star – 20 angle averaged trajectories

Froehlich et al (2005)Pruett, Hoffman, and Woosley (2005)

The neutrino-assisted rp-process

9 5

0.36 s

2.05 0.2

7

/ 77 0.562

b e

T

s k Y

13

0.124

0.844

10

p

n

X

X

X

Unmodified trajectory number 6 from Janka et al.

Entropy times two

Summary – Neutrino Wind

• Very rich site for nucleosynthesis – about half of all the isotopes in nature are made here

• Can produce both (part of) the p-process and the r-process nearly simultaneously in one site.

• In both cases, the nucleosynthesis suggests a higher entropy (and outflow with more internal energy per baryon) than traditional models have provided

•Nuclear physics very uncertain – need RIA

• This has important implications for the explosion model

Energy input from: (Qian & Woosley 1996) vibrations - Burrows et al (2006) Alfven waves or reconnection – Suzuki and Nagataki (2005) magnetic confinement – Thompson (2003)