The Evolution and Explosion of Massive Stars Nuclear Physics Issues S. E. Woosley, A. Heger, T. Rauscher, and R. Hoffman

We study nuclear astrophysics because: The origin of the elements is an interesting problem Nuclear transmutation (and gravity) are the origin of all stellar energy generation. Nuclear physics determines stellar structure. We can use that understanding as a diagnostic... of the Big Bang of stellar evolution of nova and supernova explosions of x-ray and  -ray bursts of particle physics of the evolution of galaxies and the universe

Stars are gravitationally confined thermonuclear reactors. Each time one runs out of one fuel, contraction and heating ensue, unless degeneracy is encountered. For a star over 8 solar masses the contraction and heating continue until an iron core is made that collapses. What is a massive star?

The advanced burning stages are characterized by multiple phases of core and shell burning. The nature and number of such phases varies with the mass of the star. Each shell burning episode affects the distribution of entropy inside the helium core and the final state of the star (e.g., iron core mass) can be non-monotonic and, to some extent, chaotic. Neutrino losses are higher and the central carbon abundance lower in stars of higher mass.

Iron core collapse triggers a catastrophe. The star at death is typically a red supergiant with a highly evolved, compact core of heavy elements.

Burrows, Hayes, and Fryxell, (1995), ApJ, 450, Solar masses – exploded with an energy of order erg. see also Janka and Mueller, (1996), A&A, 306, 167 Paper: Thursday - Janka

First three-dimensional calculation of a core-collapse 15 solar mass supernova. This figure shows the iso-velocity contours (1000 km/s) 60 ms after core bounce in a collapsing massive star. Calculated by Fryer and Warren at LANL using SPH (300,000 particles). The box is 1000 km across. 300,000 particles 1.15 Msun remnant 2.9 foe 1,000,000 “ 1.15 “ 2.8 foe – 600,000 particles in convection zone 3,000,000 “ in progress Fryer and Warren (2002)

Explosive Reprocessing

Rauscher, Heger, Woosley, and Hoffman (2002) nb. 62 Ni Papers: Tuesday: Heger Limongi Maeda Thursday: Nomoto

Rauscher, Heger, Woosley, & Hoffman (2002)

``There is something fascinating about science. One gets such a wholesale return of conjecture out of such a trifling investment of fact.” Mark Twain in Life on the Mississippi As cited at the beginning of Fowler, Caughlan, and Zimmerman, ARAA, 13, 69, (1975)

PROBLEMS PARTICULAR TO NUCLEAR ASTROPHYSICS Both product and target nuclei are frequently radioactive Targets exist in a thermal distribution of excited states There are a lot of nuclei and reactions (tens of thousands) Need weak interaction rates at extreme values of temperature and density Papers: Tuesday: Motobayashi Thielemann Wednesday Kaeppeler Thursday Schatz Goriely Kajino Friday Smith Rauscher

Specific Nuclear Uncertainties: 12 C(  ) 16 O 22 Ne( ,n) 25 Mg 12 C(n,  ) 13 C, 16 O(n,  ) 17 O and other 30 keV (n,  ) cross sections Neutrino spallation of 4 He, 12 C, 16 O, 20 Ne, La, Ta Weak rates for the iron group Rates for the rp-process in proton- rich winds of young neutron stars Hauser-Feshbach rates for A > 28 Photodisintegration rates for heavy nuclei for the  process – Mohr, Utsunomiya Mass excesses and half lives for the r-process Reaction rates affecting the nucleosynthesis of radioactive nuclei: 22 Na, 26 Al, 44 Ti, 56,57 Ni, 60 Co -Diehl The nuclear EOS for core collapse supernovae – Session 11 Electron capture rates at high densities (  ~ – ) for very heavy nuclei in core collapse (A up to several hundred)- Langanke (massive stars only)

12 C(a,  ) 16 O Papers: Tuesday Fey Posters: A18 Fynbo A32 Makii A47 Sagara A62 Tsentalovich

* Buchmann (1996) Heger, Woosley, & Boyse (2002)

current uncertainty Heger, Woosley, & Boyse (2002)

uncertainty Heger, Woosley, & Boyse (2002)

CF88

Papers: Monday Sneden Aoki Wednesday Kaeppeler Galino Posters: A64 – Zhang B02 – Tomyo B03 – Tomyo B09 – Sonnabend

Kaeppeler et al. 1994, ApJ, 437, 396

Jaeger et al. 2001, PRL, 87, Ne(a,n) 25 Mg

62 Ni (n,  ) 63 Ni bigger is better.... Needs measuring. s-wave extrapolation is bad. Are there others? 40 K(n,  ) 41 K (and 40 K(n,p) 40 Ar)

Rauscher, Heger, Woosley, and Hoffman (2002) nb. 62 Ni

12 C (n,  ) 13 C 16 O(n,  17 O 58,59,60 Fe(n,  ) 59,60,61 Fe Important for producing 60 Fe.

Solar Metallicity

Papers: Tuesday Thielemann Friday Rauscher

Hauser-Feshbach applicable for essentially all A>28 except near closed shells.

In general, variation of the Hauser-Feshbach rates results in approximately less than a factor of two variation in the nucleosynthesis of A < 70, but there are exceptions. The agreement will not be nearly so good for A > 70 since these nuclei are made by processes that are out of equilibrium. Hoffman et al., 1999, ApJ, 521, 735

nb. Both sets of calculations used experimental rates below A = 28 and both sets employed (n  ) rates that had been normalized, at 30 keV, to Bao and Kaepeller (1987).

(n  ) Cross Sections at 30 keV

The -Process (possibly sensitive to flavor mixing) Papers: Tuesday Langanke Heger Thielemann Wednesday Boyd Thursday Janka Poster A41 – Martinez-Pinedo

Kolbe & Langanke (2002) vs Haxton (1990) Heger, Langanke, & Woosley (2002)

T- and  -dependent weak interaction rates affect both nucleosynthesis and presupernova structure. Papers: Tuesday Langanke Posters: A34 – Sampaio A38 – Messner B18 - Borzov

conv Si burning These rates should still be regarded as very uncertain

Different choices of rates can give quite different results for key quantities at iron core collapse. Most of the difference here comes from WW95 using beta decay rates that were way too small. Need to know rates on nuclei heavier than mass 60 at higher temperature and density than

The r-Process Papers: Monday Sneden Aoki Wednesday Nishimura Thursday Goriely Kajino Sumiyoshi Friday Ryan Takahashi Wanajo Posters: A52 Ishiyama A53 Ishikawa B36 – Ishimaru B38 - Honda B30 – Tamamura B31, B32 – Terasawa B33-Panov B39 - Otsuki

Need: Binding energies (neutron-separation energies) along the r-process path Temperature-dependent beta-decay half-lives along the r-process path May need neutron-induced fission cross sections May need -induced decay rates and neutral current spallation cross sections

Nucleonic wind, seconds Anti-neutrinos are "hotter" than the neutrinos, thus weak equilibrium implies an appreciable neutron excess, typically 60% neutrons, 40% protons * favored r-Process Site #1: The Neutrino-powered Wind sensitive to the density (entropy)

Nucleonic disk 0.50 Z = N Radius Electron Mole Number Neutron-rich 1 Entropy Radius The disk responsible for rapidly feeding a black hole, e.g., in a collapsed star, may dissipate some of its angular momentum and energy in a wind. Closer to the hole, the disk is a plasma of nucleons with an increasing neutron excess. r-Process Site #2: Accretion Disk Wind

Reactions governing the assembly to carbon are critical: (e.g., Terasawa et al (2001)) Also important for the very short time scale r-process (Meyer 2001) are reactions governing the reassembly of neutrons and protons to alphas (like a neutron-rich Big Bang).

Neutrino flavor mixing and the r-process Qian et al. (1995); Qian & Fuller (1995)

Neutrino-powered wind – p-nuclei Hoffman, Woosley, Fuller, & Meyer, ApJ, 460, 478, (1996) In addition to being a possible site for the r-process, the neutrino- powered wind also produces interesting nucleosynthesis of “p-process” nuclei above the iron-group, especially 64 Zn, 70 Ge, 74 Se, 78 Kr, 84 Sr, 90,92 Zr, and 92 Mo. Reaction rate information in this mass range is non-existant.

Qian & Woosley (1996) A proton-rich wind??

Flavor mixing (e.g., Schirato and Fuller 2002) For the sun, (  m) 2 = 3.7 x eV 2 and sin 2 2   large mixing angle solution) For the (controversial) LSND result, (  m) 2 is larger, perhaps of order 1 eV 2 and the mixing angle is small (S&F02 adopt sin 2 2  = 3.5 x ). In some cases it may be possible to get a wind with Y e > 0.5 Thusday - Schatz

Specific Nuclear Uncertainties: 12 C(  ) 16 O 22 Ne( ,n) 25 Mg 12 C(n,  ) 13 C, 16 O(n,  ) 17 O and other 30 keV (n,  ) cross sections Neutrino spallation of 4 He, 12 C, 16 O, 20 Ne, La, Ta Weak rates for the iron group Rates for the rp-process in proton- rich winds of young neutron stars Hauser-Feshbach rates for A > 28 Photodisintegration rates for heavy nuclei for the  process – Mohr, Utsunomiya Mass excesses and half lives for the r-process Reaction rates affecting the nucleosynthesis of radioactive nuclei: 22 Na, 26 Al, 44 Ti, 56,57 Ni, 60 Co -Diehl The nuclear EOS for core collapse supernovae – Session 11 Electron capture rates at high densities (  ~ – ) for very heavy nuclei in core collapse (A up to several hundred)- Langanke (massive stars only)