NUCLEOSYNTHESIS IN STELLAR EVOLUTION AND EXPLOSIONS ABUNDANCE YIELDS

NUCLEOSYNTHESIS IN STELLAR EVOLUTION AND EXPLOSIONS: ABUNDANCE YIELDS FOR CHEMICAL EVOLUTION. MASSIVE STARS Marco Limongi INAF – Osservatorio Astronomico di Roma, ITALY and Centre for Stellar and Planetary Astrophysics Monash University – AUSTRALIA Email: marco@oa-roma. inaf. it Work with: Alessandro Chieffi

Massive Stars, those massive enough to explode as supernovae, play a key role in many fields of astrophysics: Evolution of Galaxies: Light up regions of stellar birth induce star formation Production of most of the elements (those necessary to life) Mixing (winds and radiation) of the ISM Production of neutron stars and black holes Cosmology (Pop. III): Reionization of the Universe at z>5 Massive Remnants (Black Holes) AGN progenitors Pregalactic Chemical Enrichment High Energy Astrophysics: Production of long-lived radioactive isotopes: (26 Al, 56 Co, 57 Co, 44 Ti, 60 Fe) GRB progenitors The understanding of these stars, is crucial for the interpretation of many astrophysical objects

Outline Basic Pre. SN Evolutionary Properties of Massive Stars and Their Uncertainties Explosive Nucleosynthesis and its uncertainties Present Status of the presupernova and explosion modelling of Massive Stars Comparison among available yields Strategies for improvements

H burning g H Conv. core g g CNO Cycle g g g Mmin(O) = 14 M t(O)/t(H burning): 0. 15 (14 M ) – 0. 79 (120 M ) MASS LOSS

Hs=0. 695 Hes=0. 285 Cs=3. 18 10 -3 Ns=1. 16 10 -3 Os=1. 00 10 -2 t=6. 8 106 yr t=2 107 yr 1 H 1 H 4 He CNO 13 C, 14 N, 17 O Ne. Na, Mg. Al 23 Na, 26 Al Hs=0. 566 Hes=0. 414 Cs=8. 42 10 -5 Ns=1. 30 10 -2 Os=7. 18 10 -4 4 He 26 Al s=2 Hs=0. 194 Hes=0. 786 10 -6 WIND t=3. 6 106 yr Cs=1. 18 10 -4 Ns=1. 34 10 -2 Os=1. 59 10 -4 26 Al s=7 10 -6 WIND WNL t=2. 7 106 yr 1 H 4 He CNO 13 C, 14 N, 17 O Ne. Na, Mg. Al 23 Na, 26 Al

Major Uncertainties in the computation of core H burning models: Extension of the Convective Core (Overshooting, Semiconvection) Mass Loss Both influence the size of the He core that drives the following evolution

He burning The properties of core He burning mainly depend on the size of the He core M ≤ 35 M RSG M > 35 M BSG g g g 3 a + 12 C(a, g)16 O g g g

11 Hs=0. 649 Hes=0. 331 25 t=2. 0 107 yr Cs=2. 00 10 -3 Ns=4. 37 10 -3 Os=7. 86 10 -3 t=1. 5 106 yr t=6. 8 106 yr t=5. 3 105 yr 4 He, 14 N 12 C, 16 O 22 Ne, s-proc 4 He Hs=0. 000 Hes=0. 516 60 t=3. 6 106 yr Cs=0. 397 Ns=0. 000 Os=0. 06 120 Hs=0. 000 Hes=0. 422 WNL Cs=0. 432 Ns=0. 000 Os=0. 119 t=3. 6 105 yr t=2. 7 106 yr WNL WNE t=3. 0 105 yr WNE WC WC 4 He, 12 C, 16 O 22 Ne, s-proc 4 He

Major Uncertainties in the computation of core He burning models: Extension of the Convective Core (Overshooting, Semiconvection) Central 12 C mass fraction (Treatment of Convection + 12 C(a, g)16 O cross section) Mass Loss (determine which stars explode as RSG and which as BSG) 22 Ne(a, n)25 Mg (main neutron source for s-process nucleosynthesis) All these uncertainties affect the size of the CO core that drives the following evolution

Advanced burning stages Neutrino losses play a dominant role in the evolution of a massive star beyond core He burning At high temperature (T>109 K) neutrino emission from pair production start to become very efficient g g n n n g n g Evolutionary times reduce dramatically

M < 30 M Explode as RSG M ≥ 30 M Explode as BSG After core He burning At Pre-SN stage

Synthesis of Heavy Elements At high tempreatures a larger number of nuclear reactions are activated Heavy nuclei start to be produced C-burning Ne-burning

Synthesis of Heavy Elements O-burning Weak Interactions become efficient Efficiency scales inversely with the mass

Synthesis of Heavy Elements At Oxygen exhaustion Balance between forward and reverse reactions for increasing number of processes c + d a + b At Oxygen exhaustion At Si ignition (panel a + panel b) A=45 A=44 Sc Si Equilibrium 56 Fe 28 Si Equilibrium Partial Eq. Out of Equilibrium Eq. Clusters Out of Eq. 56, 57, 58 Fe, 52, 53, 54 Cr, 55 Mn, 59 Co, 62 Ni NSE

H 11 M H 25 M He He 103 yr 3 yr 0. 3 yr 5 days CO CO Ne/O O Si Ne/O “Fe” H 60 M O H Si “Fe” 120 M He He CO CO Ne/O O Si Ne/O O “Fe” Si “Fe”

Chemical Composition at the Pre. SN stage O Conv. Shell 28 Si, 32 S, 36 Ar, 40 Ca, 34 S, 38 Ar C Conv. Shell 20 Ne, 23 Na, 24 Mg, 25 Mg, 27 Al s-process He Central 16 O, 12 C + s- + 16 O 28 Si “Fe” 20 Ne process He Shell 16 O, 12 C H Central+Shell 14 N, 13 C, 17 O 12 C H Central H Shell He Central 55 Mn, 59 Co, 62 Ni C conv. Shell 56, 57, 58 Fe, 52, 53, 54 Cr, O conv. Shell Si Burning Main Products Si burning(Cent. +Sehll) Burning Site 4 He 1 H

Final Masses at the Pre. SN stage oss L ass M No us Radi WIND RSG a Fin l M ass WNL WNE WC/WO s Mas e Cor ss s Ma as e M r Co CC e H He CO- HEAVY ELEMENTS Fe-Core Mass

Major Uncertainties in the computation of the advanced burning stages: Treatment of Convection (interaction between mixing and local burning, stability criterion behavior of convective shells final M-R relation explosive nucleosynthesis) Computation of Nuclear Energy Generation (minimum size of nuclear network and coupling to physical equations, NSE/QSE approximations) Weak Interactions (determine Ye hydrostatic and explosive nucleosynthesis behavior of core collapse) Nuclear Cross Sections (nucleosynthesis of all the heavy elements) Partition Functions (NSE distribution) Neutrino Losses

Explosive Nucleosynthesis and Chemical Yields Explosion Mechanism Still Uncertain • Explosion: 1 D PPM Lagrangian Hydrocode (Collella & Woodward 1984) 28 Si “Fe” 16 O 4 He H Central H Shell He Central C conv. Shell Piston O conv. Shell Si burning The explosion can be simulated by means of a piston of initial velocity v 0, located near the edge of the iron core 1 H 20 Ne 12 C • Explosive Nucleosynthesis: same nuclear network adopted in the hydrostatic evolutions v 0 is tuned in order to have a given amount of 56 Ni ejected and/or a corresponding final kinetic energy Ekin

The Final Fate of a Massive Star Z=Z oss L ass M No E=1051 erg SNII SNIb/c al Fin WNL WNE s Mas ss e r Co He Ma s s Ma ss CC a M He e r Co CO WC/WO t n na m Re ss a M Fallback Mass (M ) RSG Black Hole Fe-Core Mass Neutron Star Initial Mass (M )

RADIATION DOMINATED: Sc Ti Fe Co Ni V Cr Mn Ti Fe Si S Ar Ca K NSE/QSE Si-c Si-i Ox f(r, T, Ye) Ne Na Mg Al P Cl Ne/Cx f(r, T, Xi)

Individual Yields Different chemical composition of the ejecta for different masses

Averaged Yields averaged over a Salpeter IMF Global Properties: Initial Composition (Mass Fraction) X=0. 695 Y=0. 285 Z=0. 020 NO Dilution Mrem=0. 186 Final Composition (Mass Fraction) X=0. 444 (f=0. 64) Y=0. 420 (f=1. 47) Z=0. 136 (f=6. 84)

Major Uncertainties in the simulation of the explosion (remnant mass – nucleosynyhesis): Prompt vs Delayed Explosion (this may alter both the M-R relation and Ye of the presupernova model) How to kick the blast wave: Thermal Bomb – Kinetic Bomb – Piston Mass Location where the energy is injected How much energy to inject: Thermal Bomb (Internal Energy) Kinetic Bomb (Initial Velocity) Piston (Initial velocity and trajectory) How much kinetic energy at infinity (typically ~1051 erg) Nuclear Cross Sections and Partition Functions

Present Status of the presupernova and explosion modelling of Massive Stars Authors Mass Range Z Network Mass Loss Rot. 12 C(a, g)16 O Convection Explosion CL (2004) 13 -35 0. 000. 02 300 itosopes Fully Coup. (HMo) NO NO Kunz 2001 Schwarz. Semi NO Not Coupled Hydro/Piston Prompt LC (2006) 11 -120 0. 02 " YES NO Schwarz. Semi NO Fully Coupled Hydro(PPM) Kinetic Bomb Prompt WW (1995) 11 -40 0. 000. 02 19 (enuc) + 240 post (H-Ge) NO NO CF 88 x 1. 7 Ledoux Semiconv. Not Coupled Hydro/Piston Delayed RHHW (2002) 15 -25 0. 02 19 (enuc) + 700 -2000 (adaptive) (H-Pb) YES NO Buchmann x 1. 2 UN (2002) 13 -30 150 -270 0 240 coupled ? NO NO CF 85 NH (1988)+ TNH(1996) 13 -25 0. 02 ? NO NO CF 85 HMM (20042006) 9 -120 0. 000. 04 a network for advanced phases YES NACRE " " Schwarz. Semi NO Not Coupled " Schwarz. Overshooting Not Coupled " Hydro/Thermal Bomb Delayed " NO

Databases of Cross Sections Experimental: Caughlan et al. (1985) Caughlan & Fowler (1988) Angulo et al. (1999) NACRE Bao et al. (2000): (n, g) reactions Iliadis et al. (2001): (p, g) reactions Jaeger et al. (2001): 22 Ne(a, n)25 Mg Kunz et al. (2001): 12 C(a, g)16 O Formicola et al. (2004) LUNA collaboration: 14 N(p, g)15 O LENA collaboration: 14 N(p, g)15 O Theoretical: Woosley et al. 1978 Rauscher & Thielemann (2000) REACLIB Fuller, Fowler & Newmann (1982, 1985) (Weak) Oda et al. (1984) (Weak) Takahshi & Yokoi (1987) (Weak) Langanke & Martinez Pinedo (2000) (Weak)

Z=Z

Global Properties Z=Z Final Composition (for each solar mass returned to the ISM) LC 06 X=0. 444 (f=0. 64) Y=0. 420 (f=1. 47) Z=0. 136 (f=6. 84) WW 95 X=0. 463 (f=0. 65) Y=0. 391 (f=1. 42) Z=0. 146 (f=7. 30) RHHW 02 X=0. 482 (f=0. 65) Y=0. 340 (f=1. 42) Z=0. 178 (f=8. 90)

Strategies for improvements Round Table and Comparison Among: Evolutionary Codes (Assumptions, Numerical Algorithms, etc. ) Input Physics (EOS, Opacities, Cross Sections, Neutrino Losses, Electron Screenings, etc. ) Nuclear Network (extension, how it is included into the code) Computation of Models under the same code setup Input Physics Repository EOS, Opacities, Cross Sections, etc. (Tables and Codes) Additional comments welcome. . . Pre/Post SN models and explosive yields available at http: //www. mporzio. astro. it/~limongi