MASSIVE STARS: PRESUPERNOVA EVOLUTION, EXPLOSION AND NUCLEOSYNTHESIS Marco Limongi INAF – Osservatorio Astronomico di Roma, ITALY and Centre for Stellar and Planetary Astrophysics Monash University – AUSTRALIA

What is a Massive star ? It is a star that goes through all the hydrostatic burnings in a quiescent way from H to Si and eventually explodes as a core collapse supernova M up ’ M PISN < Massive stars < >120

Why are Massive stars important in the global evolution of our Universe? 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: GRB progenitors The understanding of these stars, is crucial for the interpretation of many astrophysical objects Production of long-lived radioactive isotopes: ( 26 Al, 56 Co, 57 Co, 44 Ti, 60 Fe)

Le SNII contribuiscono in maniera rilevante all’evoluzione chimica della Galassia. Responsabili per la nucleosintesi degli elementi con 16<A<50 and 60<A<90 BB = Big Bang; CR = Cosmic Rays; neut. = n induced reactions in SNII; IMS = Intermediate Mass Stars; SNII = Core collapse supernovae; SNIa = Termonuclear supernovae; s-r = slow-rapid neutron captures

Computation of the Presupernova Evolution of Massive Stars (p,  ) ( ,n) (,)(,) ( ,p) (p,n) (p,  ) (n,  ) (n,p) (n,  ) (  n) (  p) (  )         1. Extended Network Including a large number of isotopes and reactions (captures of light partcles, e ± captures, β± decays)

Computation of the Presupernova Evolution of Massive Stars H/He burnings: + Decoupled Adv. burnings: Coupled 2. Strong coupling between physical and chemical evolution:

Computation of the Presupernova Evolution of Massive Stars 3. Tratment of convection: - Time dependent convection - Interaction between Mixing and Local Burning D = Diffusion Coefficient

Core H burning         CNO Cycle Convective Core The Convective Core shrinks in mass Massive Stars powered by the CNO Cycle

(T  3×10 7 K) 12 C + 1 H  13 N +  13 N  13 C + e C + 1 H  14 N +  14 N + 1 H  15 O +  15 O  15 N + e N + 1 H  12 C + 4 He (99%) 16 O +  (1%) 16 O + 1 H  17 F +  17 F  17 O + e O + 1 H  14 N + 4 He CN-Cycle C  N  O  CNO Cycle NO-Cycle CNO Processed Material

20 Ne + 1 H  21 Na +  21 Na  21 Ne + e Ne + 1 H  22 Na +  22 Na  22 Ne + e Ne + 1 H  23 Na +  23 Na + 1 H  20 Ne + 4 He Ne-Na Cycle 24 Mg + 1 H  25 Al +  25 Al  25 Mg + e Mg + 1 H  26 Al +  26 Al  26 Mg + e Mg + 1 H  27 Al +  27 Al + 1 H  24 Mg + 4 He Mg-Al Cycle Ne-Na and Mg-Al Cycles During Core H Burning the central temperature is high enough (3-7×10 7 K) that the Ne-Na and Mg-Al cycles become efficient  21 Na e 25 Mg destroyed  22 Ne slightly burnt  23 Na e 26 Mg increases  26 Al (~10 -7 ) produced

Evolutionary Properties of the Interior t= yr

Evolutionary Properties of the Surface M min (O) = 14 M  t(O)/t(H burning): 0.15 (14 M  ) – 0.79 (120 M  ) Core H Burning Models

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

        He Convective Core 3  + 12 C(  ) 16 O H burning shell H exhausted core (He Core) Core He burning 4 He + 4 He  8 Be +  8 Be  4 He + 4 He 8 Be + 4 He  12 C +  3 4 He  12 C +   ad  rad Bordo iniziale CC Core Convettivo He  C,O Mix He The He convective core increases in mass

Nucleosynthesis during Core He burning 3 4 He  12 C +  12 C + 4 He  12 O +  16 O + 4 He  20 Ne +  20 Ne + 4 He  24 Mg +  Chemical composition at core He exhaustion: mainly C/O C/O ratio depends on: 1. Treatment of convection (late stages of core He burning) C(  ) 16 O cross section The C/O ratio is one the quantity that mainl affects the advanced evolution of Massive Stars (it determines the composition of the CO core)

Nucleosynthesis during Core He burning 14 N, produced by H burning activates the sequence of reactions: 14 N + 4 He  18 F +  18 F  18 O + e O + 4 He  22 Ne +  22 Ne + 4 He  25 Mg + n For the CNO cycle: X CNO (iniziale)  X 14N For e Solar composition For a Solat composition at core H exhaustion: X( 14 N) ~ ½ Z  In general: The efficiency of the 14 N reactions scales with the metallicity

14 N  22 Ne during the initial stages of core He burning During core He burning, 22 Ne is reduced by a factor of ~2 by the nuclear reaction: 22 Ne + 4 He  25 Mg + n CNO (~1/2 Z) 14 N (~1/2 Z) 22 Ne (~Z) H burning He burning Neutron Mass Fraction s-process nucleosynthesis Nucleosynthesis during Core He burning

84 Se 85 Br 86 Kr 83 As 84 As 85 As 85 Se 86 Se 86 Br 87 Br 87 Kr 88 Kr 73 Ge 74 Ge 75 Ge 76 Ge 74 As 75 As 76 As 72 Ga 73 Ga 77 As 75 Se 76 Se 77 Se 78 Se 79 Se 80 Se 81 Se 82 Se 76 Br 77 Br 78 Br 79 Br 80 Br 81 Br 82 Br 83 Br 77 Kr 78 Kr 79 Kr 80 Kr 81 Kr 82 Kr 83 Kr 84 Kr 80 As 81 As 78 As 79 As 78 Rb 79 Rb 80 Rb 81 Rb 82 Rb 83 Rb 85 Rb 84 Rb 80 Ge 77 Ge 78 Ge 79 Ge 79 Ga 76 Ga 77 Ga 78 Ga 74 Ga 75 Ga n,n,  b-b- b-b- b-b- p s r s,r s-process during Core He burning Both the neutron mass fraction and the seed nuclei abundances scale with the metallicity The abundance of the s-process nuclei scales with the metallicity

Evolutionary Properties of the Interior t= yr WIND

Evolutionary Properties of the Surface Core He Burnin g Models M ≤ 30 M   RSG M ≥ 35 M   BSG

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( ,  ) 16 O cross section) Mass Loss (determine which stars explode as RSG and which as BSG) All these uncertainties affect the size of the CO core that drives the following evolution 22 Ne( ,n) 25 Mg (main neutron source for s- process nucleosynthesis)

Advanced burning stages Neutrino losses play a dominant role in the evolution of a massive star beyond core He burning At high temperature (T>10 9 K  ~0.08 MeV) neutrino emission from pair production start to become very efficient          H burning shell H exhausted core (He Core) He burning shell He exhausted core (CO Core) Core Burning

Advanced burning stages Evolutionary times of the advanced burning stages reduce dramatically

Evolutionary Properties of the Surface M < 30 M   Explode as RSG M ≥ 30 M   Explode as BSG After core He burning Absolute Magnitude increases by ~25 At PreSN stage

Advanced Nuclear Burning Stages: Core C burning H He CO H burning shell He burning shell T~10 9 K

C-burning Main Products of C burning 20 Ne, 23 Na, 24 Mg, 27 Al Scondary Products of C burning s-process nuclesynthesis Advanced Nuclear Burning Stages: C burning At high tempreatures a larger number of nuclear reactions are activated Heavy nuclei start to be produced

H He CO H burning shell He burning shell T~1.3×10 9 K NeO C burning shell Advanced Nuclear Burning Stages: Core Ne burning

Ne-burning Advanced Nuclear Burning Stages: Ne burning Main Products of Ne burning 16 O, 24 Mg, 28 Si Scondary Products of Ne burning 29 Si, 30 Si, 32 S

H He CO H burning shell He burning shell T~2×10 9 K NeO C burning shell Advanced Nuclear Burning Stages: Core O burning O Ne burning shell

Advanced Nuclear Burning Stages: O burning O-burning 28 Si (~0.55) 32 S (~0.24) 38 Ar (~0.10) 34 S (~0.07) 36 Ar (~0.02) 40 Ca (~0.01) Main Products of O burning Secondary Products of O burning

Advanced Nuclear Burning Stages: O burning During core O burning weak interactions become efficient 42 Ca 43 Ca 44 Ca 40 Ca 41 Ca 38 K 39 K 40 K 41 K 42 K 37 K 38 Ar 39 Ar 40 Ar 41 Ar 35 Ar 36 Ar 37 Ar 38 Cl 35 Cl 36 Cl 37 Cl 33 Cl 34 Cl 35 S 36 S 37 S 33 S 34 S 32 S 31 S 33 P 34 P 32 P 31 P 30 P 27 Si 33 Si 32 Si 31 Si 30 Si 28 Si 29 Si 27 Al 26 Al 28 Al 29 P Proton Number (Z) Neutron Number (N) 31 S(  + ) 31 P 33 S(e -, ) 33 P 30 P(e -, ) 30 Si 37 Ar(e -, ) 37 Cl Most efficient processes: The electron fraction per nucleon

H He CO NeO Advanced Nuclear Burning Stages: Core Si burning O SiS H burning shell He burning shell T~2.5×10 9 K C burning shell Ne burning shell O burning shell

Non equilibrium Full equilibrium Advanced Nuclear Burning Stages: Si burning At Oxygen exhaustion Balance between forward and reverse (strong interaction) reactions for increasing number of processes i + k j + l A measure of the degree of equilibrium reached by a couple of forward and reverse processes

At Oxygen exhaustion Si Sc Equilibrium At Si ignition Out of Equilibrium Equilibrium Partial Eq. Out of Eq. At Si ignition (panel a + panel b) A=44 A=45 Eq. Clusters 28 Si 56 Fe Advanced Nuclear Burning Stages: Si burning

56,57,58 Fe, 52,53,54 Cr, 55 Mn, 59 Co, 62 Ni NSE A=44 A=45 Clusters di equilibrio 28 Si 56 Fe 24 Mg 20 Ne 16 O 12 C 4 He Si is burnt through a sequence of  reactions 2.The two QSE clusters reajdust on the new equilibrium abundances of the light particles 3.The matter flows from the lower to the upper cluster through a sequence of non equilibirum reactions Equilibrium Clusters 4.Y e is continuosuly decreased by the weak interactions (out of equilibrium)

H He CO NeO Pre-SuperNova Stage O SiS H burning shell He burning shell T~4.0×10 9 K C burning shell Ne burning shell O burning shell Si burning shell Fe

Evolutionary Properties of the Interior H burning shell He burning shell C burning shell Ne burning shell O burning shell Si burning shell

Chemical Stratification at PreSN Stage Each zone keeps track of the various central or shell burnings 14 N, 13 C, 17 O 12 C, 16 O 12 C, 16 O s-proc 20 Ne, 23 Na, 24 Mg, 25 Mg, 27 Al, s-proc 16 O, 24 Mg, 28 Si, 29 Si, 30 Si 28 Si, 32 S, 36 Ar, 40 Ca, 34 S, 38 Ar 56,57,58 Fe, 52,53,54 Cr, 55 Mn, 59 Co, 62 Ni NSE

FaseTime (yr) L nuc L M cc TcTc cc M shell FuelMain Prod. Sec. Prod. H5.93(6) (7) H1H 4 He 13 C, 14 N, 17 O He6.8(5) (8)4.7(2)6 4 He 12 C, 16 O 18 O, 22 Ne, s-proc. C9.7(2)1.0(6)- 5.0(7) 4.0(7)- 1.0(9) 7.2(8)1.2(5) C 20 Ne, 23 Na, 24 Mg, 27 Al 25 Mg, s-proc. Ne7.7(-1) (280 d) 7.0(9)2.2(9) (9)2.1(6) Ne 16 Ne, 24 Mg 29 Si, 30 Si O3.3(-1) (120 d) 5.0(10) 5.9(11) 4.0(10) (9)4.0(6) O 28 Si, 32 S, 36 Ar, 40 Ca, Cl, Ar, K, Ca Si2.1(-2) (7 d) 1.1(13)1.0(12) (9)7.5(7) Si 54 Fe, 56 Fe, 55 Fe Ti, V, Cr, Mn, Co, Ni Main Properties of the PreSN Evolution

Evolution of More Massive Stars: Mass Loss O-Type: > T(K) > WNL: < H sup <0.4 (H burning, CNO, products) WNE: H sup <10 -5 (No H) WN/WC: 0.1 < X(C)/X(N) < 10 (both H and He burning products, N and C) WC: X(C)/X(N) > 10 (He burning products) Wolf-Rayet : Log 10 (T eff ) > 4.0

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 Y e  hydrostatic and explosive nucleosynthesis  behavior of core collapse) Nuclear Cross Sections (nucleosynthesis of all the heavy elements) Neutrino Losses Partition Functions (NSE distribution)

THE EVOLUTION UP TO THE IRON CORE COLLAPSE The Iron Core is mainly composed by Iron Peak Isotopes at NSE The following evolution leads to the collapse of the Iron Core: The Fe core contracts to gain the energy necessary against gravity T,  increase  nuc lowers becaus the matter is at NSE The Fe core begins to degenerate The Chandrasekhar Mass M Ch = 5.85 × (Y e ) 2 M  is reached A strong gravitational contraction begins The Fermi energy increases  the electron captures on both the free and bound protons incease as well T c ~ K,  c ~ K P e ~ dyne/cm 2 P i ~ 2×10 26 dyne/cm 2 P rad ~ 3×10 25 dyne/cm 2 The main source of pressure against gravity (electron Pressure) lowers The gravitational collapse begins

Fe Core diffusion Neutrino Trapping Core Bounce and Rebounce Shock wave Fe Core Stalled Shock Eenergy Losses 2 x erg/0.1M  “Prompt”shocks eventually stall! -sphere

Strong Shock vs Weak Shock A strong shock propagates. Matter is ejected. A weak shock stalls. Matter falls back.

 diffusion n,p p,n e +,e - heating cooling Gain Radius R G = Km Stalled Shock R S = Km Neutrinosphere R = Km Neutrino-driven explosions Energy deposition behind the stalled shock wave due to neutrino interactions: Shock Wave reheated Explosion

Propatagiont of the shock wave through the envelope Explosive Nucleosynthesis Compression and Heating Explosive Nucleosynthesis Explosion Mechanism Still Uncertain The explosive nucleosynthesis calculations for core collapse supernovae are still based on explosions induced by injecting an arbitrary amount of energy in a (also arbitrary) mass location of the presupernova model and then following the development of the blast wave by means of an hydro code. Piston Thermal Bomb Kinetic Bomb

Induced Explosion and Fallback Injected Energy Induced Shock Compressio n and Heating Induced Expansio n and Explosio n Initial Remnant Matter Falling Back Mass Cut Initial Remnant Final Remnant Matter Ejected into the ISM E kin  erg

Composition of the ejecta The Iron Peak elements are those mostly affected by the properties of the explosion, in particular the amount of Fallback.

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 SNIISNIb/c Fallback RSG Z=Z  E=10 51 erg Initial Mass (M  ) Mass (M  )

Major Uncertainties in the simulation of the explosion (remnant mass – nucleosynyhesis): Prompt vs Delayed Explosion (this may alter both the M-R relation and Y e 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 ~10 51 erg) Nuclear Cross Sections and Partition Functions

Chemical Enrichment due to Massive Stars Different chemical composition of the ejecta for different masses

Chemical Enrichment due to Massive Stars Yields of Massive Stars used for the interepretation of the chemical composition of the Galaxy We can have information on the contribution of massive stars to the solar composition by looking at the PFs of solar metallicity massive star models. ASSUMPTIONS The average metallicity Z grows slowly and continuously with respect to the evolutionary timescales of the stars that contribute to the environment enrichment Most of the solar system distribution is the result (as a first approximation) of the ejecta of ‘‘quasi ’’–solar-metallicity stars. The PF of the chemical composition provided by a generation of solar metallicity stars should be flat

Chemical Enrichment due to Massive Stars Yields averaged over a Salpeter IMF Oxygen is produced predominantly by the core-collapse supernovae and is also the most abundant element produced by these stars Use PF(O) to represents the overall increase of the average ‘‘metallicity ’’ and to verify if the other nuclei follow or not its behavior

Chemical Enrichment due to Massive Stars Elements above the compatibility range  may constitute a problem Elements below the compatibility range  produced by other sources Secondary Isotopes? No room for other sources (AGB) Type Ia AGB No room for AGB process. Other sources uncertain Explosion?

Chemical Enrichment due to Massive Stars Global Properties: Initial Composition (Mass Fraction) X=0.695 Y=0.285 Z=0.020 Final Composition (Mass Fraction) X=0.444 (f=0.64) Y=0.420 (f=1.47) Z=0.136 (f=6.84) NO Dilution M rem = M1M 1M1M IMF: Salpeter

Averaged Yields: Relative Contributions Stars with M>35 M  (SNIb/c) contribute for ~20% at maximum (large fallback) with few exceptions (H,He burning)

CONCLUSIONS Stars with M<30 M  explode as RSG Stars with M≥30 M  explode as BSG The minimum masses for the formation of the various kind of Wolf-Rayet stars are: WNL: M  WNE: M  WNC: M  The final Fe core Masses range between: M Fe = M  for M ≤ 40 M  M Fe = M  for M > 40 M  The limiting mass between SNII and SNIb/c is : M  SNIISNIb/c Salpeter IMF The limiting mass between NS and BH formation is: M  NSBH (uncertainties on mass loss, simulated explosion, etc.)

CONCLUSIONS Assuming a Salpeted IMF the efficiency of enriching the ISM with heavy elements is: H: decreased by f=0.64 He: increased by f=1.47 Metals: increased by f=6.84 For each solar mass of gas returned to the ISM Massive Stars are responsible for producing elements with 4<Z<38 SNIb/c contribute for ~20% to the majority of the elements (large fallback) SNIb/c contribute for ~40% to the elements produced by H and He burning that survive to fallback Depends on: Simulated expl. Mass Loss Binary Systems Pre/Post SN models and explosive yields available at