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 (PopIII): 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: (26Al, 56Co, 57Co, 44Ti, 60Fe) GRB progenitors The understanding of these stars, is crucial for the interpretation of many astrophysical objects

Outline Basic PreSN 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 CNO Cycle Mmin(O) = 14 M t(O)/t(H burning): 0.15 (14 M ) – 0.79 (120 M) MASS LOSS

WNL t=6.8 106 yr t=2 107 yr t=3.6 106 yr t=2.7 106 yr Hs=0.695 Cs=3.18 10-3 Hes=0.285 Ns=1.16 10-3 Os=1.00 10-2 t=6.8 106 yr t=2 107 yr 1H  4He 1H  4He CNO  13C,14N, 17O NeNa,MgAl  23Na, 26Al CNO  13C,14N, 17O NeNa,MgAl  23Na, 26Al Hs=0.566 Hes=0.414 Cs=8.42 10-5 Ns=1.30 10-2 Os=7.18 10-4 26Als=2 10-6 Hs=0.194 Hes=0.786 Cs=1.18 10-4 Ns=1.34 10-2 Os=1.59 10-4 26Als=7 10-6 WIND t=3.6 106 yr WIND WNL t=2.7 106 yr 1H  4He 1H  4He CNO  13C,14N, 17O NeNa,MgAl  23Na, 26Al CNO  13C,14N, 17O NeNa,MgAl  23Na, 26Al

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 g M > 35 M  BSG g g 3a + 12C(a,g)16O g g g g g

11 25 Hs=0.649 Hes=0.331 Cs=2.00 10-3 Ns=4.37 10-3 Os=7.86 10-3 t=2.0 107 yr t=1.5 106 yr t=6.8 106 yr t=5.3 105 yr 4He, 14N 4He, 14N 4He  12C, 16O 22Ne, s-proc 4He  12C, 16O 22Ne, s-proc 120 Hs=0.000 Hes=0.516 Cs=0.397 Ns=0.000 Os=0.06 Hs=0.000 Hes=0.422 Cs=0.432 Ns=0.000 Os=0.119 60 WNL t=3.6 106 yr t=3.6 105 yr t=2.7 106 yr t=3.0 105 yr WNL WNE WNE WC WC 4He, 12C 4He, 12C 4He  12C, 16O 22Ne, s-proc 4He  12C, 16O 22Ne, s-proc

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

g n 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 n g Evolutionary times reduce dramatically

After core He burning At Pre-SN stage 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 Weak Interactions become efficient O-burning Efficiency scales inversely with the mass

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

11 M 25 M 60 M 120 M H H He He 103 yr 3yr 0.3yr 5 days CO CO O Ne/O Si “Fe” Ne/O O Si “Fe” H H 60 M 120 M He He CO CO Ne/O Ne/O O O Si “Fe” Si “Fe”

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

Final Masses at the PreSN stage No Mass Loss Final Mass He-Core Mass He-CC Mass CO-Core Mass Fe-Core Mass WNL WNE WC/WO RSG Radius WIND HEAVY ELEMENTS

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 The explosion can be simulated by means of a piston of initial velocity v0, located near the edge of the iron core 16O 28Si 20Ne 12C 4He 1H “Fe” Piston Si burning O conv. Shell C conv. Shell He Central He Shell H Shell H Central Piston Explosion: 1D PPM Lagrangian Hydrocode (Collella & Woodward 1984) Explosive Nucleosynthesis: same nuclear network adopted in the hydrostatic evolutions v0 is tuned in order to have a given amount of 56Ni ejected and/or a corresponding final kinetic energy Ekin

The Final Fate of a Massive Star No Mass Loss Final Mass He-Core Mass He-CC Mass CO-Core Mass Fe-Core Mass WNL WNE WC/WO Remnant Mass Neutron Star Black Hole SNII SNIb/c Fallback RSG Z=Z E=1051 erg Initial Mass (M) Mass (M)

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

Individual Yields Different chemical composition of the ejecta for different masses

Averaged Yields Yields averaged over a Salpeter IMF Global Properties: Initial Composition (Mass Fraction) NO Dilution Final Composition (Mass Fraction) Mrem=0.186 X=0.695 Y=0.285 Z=0.020 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. 12C(a,g)16O Convection Explosion CL (2004) 13-35 0.00-0.02 300 itosopes Fully Coup. (H-Mo) NO Kunz 2001 Schwarz. Semi NO Not Coupled Hydro/Piston Prompt LC (2006) 11-120 0.02 " YES Fully Coupled Hydro(PPM) Kinetic Bomb WW (1995) 11-40 19 (enuc) + 240 post (H-Ge) CF88x1.7 Ledoux Semiconv. Delayed RHHW (2002) 15-25 700-2000 (adaptive) (H-Pb) Buchmann x 1.2 UN (2002) 13-30 150-270 240 coupled ? CF85 Hydro/Thermal Bomb NH (1988)+ TNH(1996) 13-25 ? HMM (2004-2006) 9-120 0.00-0.04 a network for advanced phases NACRE Overshooting

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): 22Ne(a,n)25Mg Kunz et al. (2001): 12C(a,g)16O Formicola et al. (2004) LUNA collaboration: 14N(p,g)15O LENA collaboration: 14N(p,g)15O 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 Z=Z

Final Composition (for each solar mass returned to the ISM) Global Properties Z=Z Final Composition (for each solar mass returned to the ISM) LC06 WW95 RHHW02 X=0.444 (f=0.64) Y=0.420 (f=1.47) Z=0.136 (f=6.84) X=0.463 (f=0.65) Y=0.391 (f=1.42) Z=0.146 (f=7.30) 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