1. Core Collapse SNe Inma Domínguez Marco Limongi

2.  Evolution of Massive Stars  Hydrostatic Nucleosynthesis  Explosion Mechanism  Explosive Nucleosynthesis  Contribution to the Chemical Evolution

3. Type II SNe  Chemical Evolution of the Galaxy Type II SNe  16 < A < 50 and 60 < A < 90 16 O 49 Ti 60 Ni 90 Zr BB = Big Bang; CR = Cosmic Rays; neut. = induced reactions in SNII; IMS = Intermediate Mass Stars; SNII = Core collapse supernovae; SNIa = Termonuclear supernovae; s-r = slow-rapid neutron captures INTERPRETATION OF THE SOLAR SYSTEM ABUNDANCES

4. Evolutionary Properties of Massive Stars: Progenitors of CCSNe  M > 12 M  CCSNe  Central Conditions (T,  ) Ignition of ALL Exothermic Nuclear Reactions  The stars is never in degenerate conditions along its evolution

5. STELLAR EVOLUTION EQUATIONS 1 Dimension Lagrangian Hydrostatic Mixing-length theory

6. STELLAR EVOLUTION EQUATIONS + Chemical Evolution Production + Destruction For each time step 1000 (zones) systems of 4+N(isotopes) equations High Computational Time

7. HYDROGEN BURNING - PP 4H  He 1 H + 1 H  2 H + e + + 2 H + 1 H  3 He +  3 He + 3 He  4 He + 2 1 H PP I 3 He + 4 He  7 Be +  7 Be + e -  7 Li + 7 Li + 1 H  2 4 He 7 Be + 1 H  8 B +  8 B  8 Be + e + + 8 Be  2 4 He PP II PP III Proton-Proton Chain Depending on T the different branchings become active. In all cases the result is 4 1 H  1 4 He

8. HYDROGEN BURNING CNO Cycle 12 C + 1 H  13 N +  13 N  13 C + e + + 13 C + 1 H  14 N +  14 N + 1 H  15 O +  15 O  15 N + e + + 15 N + 1 H  12 C + 4 He (99%) 16 O +  (1%) T  3 10 7 K When C and/or N and/or O are present  CNO 16 O + 1 H  17 F +  17 F  17 O + e + + 17 O + 1 H  14 N + 4 He CN NO During the conversion of H into He through the CNO cycle C and O are burnt and N is produced Products of CNO C  N  O 

9. HYDROGEN BURNING – ENERGY GENERATION The CNO cycle is more efficient than he PP chain over a certain T critica CNO PP From Hydrostatic Equilibrium Eq: Central Temperatura scales with Total Mass Massive stars H-burning CNO cycle

10. HYDROGEN BURNING - CONVECTIVE CORE The Energy generated by the CNO-cycle depends strongly on T High Energy Flux  Increases Radiative Gradient  A Convective core Develops Masssive stars burn H within a Convective core At high T the main contribution to the Opacity comes from the Thomson Scattering When the H decreases, the Opacity decreases and the Convective Core receeds and finally, at H-exhaustion, disappears

11. HYDROGEN BURNING – Ne-Na, Mg-Al Cycles If during the central convective H-burning T are high enough log T=7.5-7.8 Active Ne-Na e Mg-Al cycles 20 Ne + 1 H  21 Na +  21 Na  21 Ne + e + + 21 Ne + 1 H  22 Na +  22 Na  22 Ne + e + + 22 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 + + 25 Mg + 1 H  26 Al +  26 Al  26 Mg + e + + 26 Mg + 1 H  27 Al +  27 Al + 1 H  24 Mg + 4 He Mg-Al Cycle Final results of the operation of these cycles Na-Na e Mg-Al  21 Na & 25 Mg practically burnt  22 Ne is reduced by a factor 2  23 Na & 26 Mg increase by a factor 6 & 2, respectively  26 Al produced (~10 -7 )  20 Ne, 24 Mg & 27 Al do not change

12. STRUCTURE AT CENTRAL H-EXHAUSTION The He-core is much more dense than the H-envelope because the mean molecular weight for 4 He is greater than for 1 H  Matter within the He-core is more compact He coreH envelope The synthesis of heavier isotopes increases the mean molecular weight and the structure becomes more compact

13. Convective envelope H conv. core He conv. core H burn.shell He burn. shell He He core CO core dup HYDROGEN SHELL BURNING  At central H-exhaustion H-burning sets in a Shell outside the He-core.  HR diagram: the star moves to the red  A convective envelope forms, the inner border of this envelope reachs zones chemically modified by he central H-burning.  The 1st dredge-up occurs: material processed by nuclear reactions is transported to the surface H exhaustion Start Conv. Env.

14. HELIUM BURNING – 3  At central H-exhaustion, the He core is mainly composed by 4 He (98%) & 14 N (1%) Withouth Nuclear Energy generation within the core, it contracts and T c increases When T c ~ 1.5 10 8 K  Efficient He-burning 4 He + 4 He  8 Be +  8 Be  4 He + 4 He At the beginning 4 He  8 Be and 8 Be rapidly decays to 4 He 4 He + 4 He  8 Be +  8 Be  4 He + 4 He 8 Be + 4 He  12 C +  Later, at higher T and   the equilibrium abundance of 8 Be increases and so increases the probability of the reaction 8 Be + 4 He producing 12 C 3 4 He  12 C + 

15. HELIUM BURNING – REACTIONS Initially: 4 He in 12 C But when 12 C abundance is significant and 4 He abundance is reduced, it is more likely that 4 He is captured by 12 C than by 4 He: 3 4 He  12 C +  12 C + 4 He  12 O +  16 O + 4 He  20 Ne +  20 Ne + 4 He  24 Mg +  3 4 He Nuclear Cross Section depends markedly on T Like H-burning (CNO cycle) He-burning occurs within a convective core The first 2 reactions are more efficient

16. HELIUM BURNING: s-process 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- In Massive  during central He-burning, elements heavier than Fe are synthesized by the s-process. s-process depends on free neutrons and the neutron abundance depends on Z  The final s-element abundances scale with initial metallicity 14 N + 4 He  18 F +  18 F  18 O + e + + 18 O + 4 He  22 Ne +  22 Ne + 4 He  25 Mg + n 14 N produced by the CNO cycle

17. HELIUM EXHAUSTION The most abundant isotopes at central He-exhaustion: 12 C 16 O 20 Ne 25 Mg 26 Mg The first three are produced by: 3 4 He  12 C +  12 C + 4 He  12 O +  16 O + 4 He  20 Ne +  25 Mg & 26 Mg come from the 14 N-chain 14 N + 4 He  18 F +  18 F  18 O + e + + 18 O + 4 He  22 Ne +  22 Ne + 4 He  25 Mg + n 22 Ne + 4 He  26 Mg +  12 C, 16 O, 20 Ne, 25 Mg & 26 Mg are the most abundant isotopes and are produced by He-burning with the surface abundance 12 C/ 16 O ratio depends on the 12 C + 4 He  12 O +   nuclear cross section that it is still NOT well known at the energies of the He burning. This ratio has a strong influence on the subsequent evolution 12 C 16 O 20 Ne 22 Ne 25 Mg 26 Mg ex He c.c. H sh. Conv. Envelope. Core di CO

18. HELIUM EXHAUSTION: s-process elements The most abundant elements are: 70 Ge 74 Se and 80 Kr Heavier nuclei, like 87 Rb, 88 Sr, 89 Y, 90 Zr are not expected to be produced 70 Ge 80 Kr 74 Se ex He c.c. Conv. Envelope.H sh. Core di CO

19. At central He exhaustion, He burning moves to a shell just outside the CO core The following evolution is characterized by the development of a convective He-burning shell limited by the CO core and by the H-burning shell. The chemical composition of this shell, that will be active till the collapse, tends to get frozen because the evolution of the star is more and more rapid at the advanced phases. Convective envelope H conv. core He conv. core H burn.shell He burn. shell He He core CO core dup He conv.shell HELIUM SHELL BURNING – CONVECTIVE SHELL

20. STRUCTURE at He-exhaustion At central H-exhaustion, the  is composed by a CO core, a He-shell and a rich H envelope He core H envelope CO core The two density gradients correspond to the border of the He core (~ 9 M  ) and to the border of the CO core (~ 6 M  ) This density profile is important for the explosion properties

21. ADVANCED EVOLUTIONARY PHASES: NEUTRINO DOMINATED Now the CO core, produced by the central He-burning, contracts During the contraction the  and T within the core favours the production of thermal neutrinos produced by pair anhilation. At T>10 9 K high energy photons produce e + e - pairs That suddenly recombine to produce a photon. BUT once over 10 19 times, e + e - produces a neutrino-antineutrino pair This energy sink increases along the subsequent phases up to the pre-collapse phase Advanced evolutionary phases of massive stars are called “neutrino dominated”

22. ADVANCED EVOLUTIONARY PHASES: NEUTRINO LUMINOSITY From now on the energy losses: Photons from the surface Neutrinos from the center Photon Nuclear Neutrino Up to C central ignition the main energy losses are due to photons and after are due to neutrinos. As the nuclear energy gives the star what is lossing, it follows first the luminosity of photons, and after, the neutrino luminosity 10 8

23. EVOLUTIONARY TIMES E nuc is the energy per gram coming from nuclear reactions, If this is the only energy source in a star of mass M: Nuclear time scale: H burning: 4 1 H  4 He  M = 4 x 1.0078 – 4.0026 = 0.0287 AMU = 0.0287/4 AMU/nucleon = 0.007 AMU/nucleon E nuc = 0.007 x 931.1 x 1.602 10 -6 x 6.022 10 23 = 6.44 10 18 erg/g 1 AMU = 931.1 MeV : 1 MeV= 1.602 10 -6 erg : N A = 6.022 10 23 nucleon/g He burning: 4 4 He  16 O  M = 4 x 4.0026 – 15.9949= 0.0115 AMU = 0.0115/16 AMU/nucleon = 0.0009 AMU/nucleon E nuc = 0.0009 x 931.1 x 1.602 10 -6 x 6.022 10 23 = 8.70 10 17 erg/g O burning: 2 16 O  32 S  M = 2 x 15.9949 – 31.9720= 0.0177 AMU = 0.0177/32 AMU/nucleon = 0.0005 AMU/nucleon E nuc = 0.0005 x 931.1 x 1.602 10 -6 x 6.022 10 23 = 4.98 10 17 erg/g For fix mass, Luminosity and amount of fuel From models: ! The luminosity increases drastically due to neutrino losses  The evolutionary times are drastically reduced

24. 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) neutrino emission from pair production start to become very efficient         Evolutionary times reduce dramatically

25. CARBON BURNING Central C combustion stars ~10 4 years after central He-exhaustion T c ~ 7 10 8 K e  c ~ 1 10 5 g/cm 3 C-burning depends on the 12 C/ 16 O ratio left after central He burning, 12 C(  ) 16 O on the amount of fuel The formation of a Convective Core depends on the existence of a positive energy flux 12 C abundances determines the nuclear energy generation rate A Convective Core develops  nuc >  NO Convective Core  nuc <  In general, for a fix 12 C( ,  ) 16 O reaction rate and mixing technics 12 C abundance decreases for higher initial masses In the 25M   central carbon combustion occurs in radiative conditions

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

27. Synthesis of Heavy Elements O-burning

28. Synthesis of Heavy Elements At Oxygen exhaustion Balance between forward and reverse reactions for increasing number of processes a + b c + d 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 56,57,58 Fe, 52,53,54 Cr, 55 Mn, 59 Co, 62 Ni

29.  =0.000,Ye=0.5000, 56 Ni  =0.038,Ye=0.481, 54 Fe  =0.072,Ye=0.464, 56 Fe  =0.104,Ye=0.448, 58 Fe MATTER PROPERTIES AT HIGH TEMPERATURE :NSE The chemical composition of matter in NSE is a function of T  Y e When the neutronization changes The nuclei with that neutron excess are favoured (with higher binding energies)

30. PRE-SUPERNOVA MODEL: CHEMICAL COMPOSITION Burning SiteMain Products Si Burning 54 Fe, 56 Fe, 55 Fe, 58 Ni, 53 Mn 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 Centrale 16 O, 12 C + s- process He Shell 16 O, 12 C H Centrale+Shell 14 N, 13 C, 17 O Si burning(Cent.+Sehll) O conv. Shell C conv. Shell He Centrale He Shell H Shell H Centrale 16 O 28 Si 20 Ne 12 C 4 He 1H1H “Fe” Studying the different isotope abundances in detail is possible to know from which burning phase they come from or the interior region of the star where they were produced

31. PRE-SUPERNOVA MODEL: Fe-CORE STRUCTURE Fe/Si Si/O CO/He He/H 16 O 20 Ne 12 C 28 Si “Fe”

32. EXPLOSION The gravitational collapse of a stars with M  12 M  could liberate an energy of Most of this energy increases the electron energy and, after electron captures, is converted in neutrino energy Just a small fraction is used to eject (kinetic energy) the envelope So, the key question is to find a mechanism able to transform a small fraction of the binding energy left during the collapse in kinetic energy of the envelope with the observed velocities ( 10 4 km/s)

33. Explosive Nucleosynthesis and Chemical Yields Explosion Mechanism Still Uncertain The explosion can be simulated by means of a piston of initial velocity v 0, located near the edge of the iron core v 0 is tuned in order to have a given amount of 56 Ni ejected and/or a corresponding final kinetic energy E kin Explosion: 1D PPM Lagrangian Hydrocode (Collella & Woodward 1984) Explosive Nucleosynthesis: same nuclear network adopted in the hydrostatic evolutions 16 O 28 Si 20 Ne 12 C 4 He 1H1H “Fe” Piston Si burning O conv. Shell C conv. Shell He Central He Shell H Shell H Central

34. EXPLOSIVE NUCLEOSYNTHESIS Passing through the envelope the Shock Wave increases the density and temperature and nuclear reactions occur We may define the burning time-scales for the available fuels : These time scales are determined by the corresponding destructive reactions Si, O, Ne, C, He and H Assuming the explosion time ~1s Burning products are similar to those obtained in hydrostatic burning He-explosive burning is not efficient in SNII

35. EXPLOSIVE NUCLEOSYNTHESIS Analyzing the most eficient processes: EXPLOSIVE CARBON BURNING: Products: 20 Ne, 23 Na, 24 Mg, 25 Mg, 26 Mg EXPLOSIVE NEON BURNING: Products: 16 O, 24 Mg + 27 Al, 29 Si, 30 Si, 31 P, 35 Cl, 37 Cl EXPLOSIVE OXYGEN BURNING: Products: 28 Si, 32 S, 36 Ar, 40 Ca + 34 S, 38 Ar Still out of NSE: Products are similar to those from hydrostatic burning Starting NSE (direct and inverse process) 2 clusters at quasi-NSE separated by A  44. No connection between the 2 clusters A=44 A=45 Clusters di equilibrio 28 Si 56 Fe

36. EXPLOSIVE NUCLEOSYNTHESIS EXPLOSIVE INCOMPLETE SILICON BURNING: Products: 36 Ar, 40 Ca + 56 Ni( 56 Fe), 54 Fe, 52 Fe( 52 Cr), 51 Cr( 51 V), 55 Co( 55 Mn), 57 Ni( 57 Fe), 58 Ni At this T the 2 clusters connect at A  44. Most of the matter A<44  just part of 28 Si reachs the upper cluster A=44 A=45 Clusters di equilibrio 28 Si 56 Fe EXPLOSIVE COMPLETE SILICON BURNING: At this high temperature: NSE !!!!!! All 28 Si is burnt to Fe-peak elements. Abundances depend on neutronization !! For N  Z 56 Ni is the most abundant nuclei Full NSE Products: Iron Peak Nuclei

37. EXPLOSIVE NUCLEOSYNTHESIS During the explosion Temperatures are very high It could be assumed that matter behind the shock is radiation dominated The shock propagates in all directions (sphere) Each radial coordinate in the presupernova model will reach a maximum temperature = Location and T of the shock Changes in T and  following expansion are crucial for the nucleosynthesis

38. EXPLOSIVE NUCLEOSYNTHESIS For E expl =10 51 erg we could infer in the presupernova model which regions (volumes) experience each burning 37005000 640011750 13400 Complete Si burning Incomplete Si burning Explosive Oxygen Explosive Neon Explosive Carbon Untouched Zone NSE QSE 1cluster QSE 2cluster Ne,Na,MgMg,Al, P, ClSi,S,Ar, K,Ca Cr,V,Mn, Fe Sc,Ti,Fe, Co,Ni

39. EXPLOSIVE NUCLEOSYNTHESIS: PROGENITOR Influence of the Progenitor : 1) M-R RELATION (= density profile): Fix the mass inside a certain volume 2) Ye (neutronization): In those zones that reach NSE or QSE determines the rate between protons and neutrons 3) Chemical Composition : For those zones that experience normal burnings (ie. Explosive Carbon e Neon burnings) fix the amount of fuel available. T=5 10 9 K,  = 10 8 g/cm 3, Y e =0.50  56 Ni=0.63 – 55 Co=0.11 – 52 Fe=0.07 – 57 Ni=0.06 – 54 Fe=0.05 T=5 10 9 K,  = 10 8 g/cm 3, Y e =0.49  54 Fe=0.28 – 56 Ni=0.24 – 55 Co=0.16 – 58 Ni=0.11 – 57 Ni=0.08

40. MASS CUT The Mass Cut depends on the piston initial velocity Mass Cut During the explosion internal zones fall back. At some point part of the matter is Expanding and some Collapsing Depending on v compare to v esc  The mass coordinate at the bifurcation is defined as the Mass Cut 1.110 1.144 1.170 1.220 1.250 1.263 The lack of a explosion model makes the MASS CUT and the KINETIC ENERGY quantities that depend on parameters (initial energy or piston initial velocity and place at which the explosion is started) In general, for greater initial velocities Smaller Mass Cut Greater kinetic Energies

41. 16 O 20 Ne 12 C 28 Si 4 He 1H1H Pre = Dotted Post = Solid OxOx Ne x CxCx Untouched Si-c Si-i FallBack EXPLOSION PROPERTIES: CHANGES IN CHEMESTRY  The changes in composition due to the explosion occur only at the most internal ~3.1 M   Outside the chemical composition remains untouched. It is that from the hydrostatic burning  The complete explosive Si burning and part ot the incomplete explosive Si burning fall back to the compact remant Mass Cut v 0 =1.5550 10 9 cm/sM cut =1.89 M  E kin =1.14 foe Taken:

42. MASS CUT CALIBRATION: LIGHT CURVES From the LC we obtain information for the M cut 56 Ni 56 Co Total 56 Ni=0.15 M  56 Ni=0.07 M  56 Ni=0.01 M  8.8111 Based on the Bolometric LCs and on the distance, we can deduce the amount of 56 Ni produced during the explosion After an initial phase, different for the different types of SNe, the LC is powered by the photons produced by the radioactive decay 56 Ni is produced in the most internal zone depends critically on the Mass Cut  The Mass Cut may be choose to reproduce a certain amount of 56 Ni in agreement with the observations. The theoretical kinetic energy must be compatible with the observed

43. MASS CUT CALIBRATION vs INITIAL MASS From the observed initial mass of the progenitor we may obtain an empirical relaction between this mass and the 56 Ni produced (or M cut)  Few estimations of the progenitor initial mass from the observations  Similar masses give very different 56 Ni masses PROBLEMS !!!! Hamuy et al. 2003

44. CHOOSING A MASS CUT 1) FLAT Case: All masses produce the same 56 Ni mass = 0.05 M   For each model a different mass cut is chosen in order to reproduce this amount of Ni 2) TREND Case: We adopt a relation between Initial Mass and 56 Ni Mass: M i (M  ) M( 56 Ni) (M  ) 130.15 150.10 200.08 250.07 300.05 350.05

45. PRODUCTION FACTORS To compare with Solar Abundances we introduce the Production Factor Two isotopes with the same Production Factor Same Rate as in the Sun Oxygen is produced only by Type II SNe and is the most abundant element produced by SNII  Oxygen Production Factor is a Good Metallicity indicator It is useful to normalize all PF to that of Oxygen to show wich isotopes follow Oxygen (Z)

46. Dots: 13 – 15 – 20 – 25 – 30 – 35 M  Solid line: Salpeter Mass Function Flat 56 Ni => 0.05 M  INTEGRATED YIELDS (Elements) Yields from 13-35 M  + Salpeter Mass Function It is assumed that all masses produce the same amount of 56 Ni (FLAT) We consider “Solar Scaled” with respect to O all elements with a PF within a factor 2 of the O PF The yields produced by a generation of massive stars integrated by a Salpeter IMF depend mainly on the yields coming from a 20-25 M  star

47. Production of Fe  the percentage of SNIa, relative to SNII, has been fixed by requiring that PF Fe =PF O Open circles = No SNIa Filled circles = 12% SNIa 1)SNIa contribute only to the Solar System abundances of nuclei in the range Ti-Ni 2)The inclusion of SNIa brings 50 Ti and 54 Cr into the band of compatibility  50 Ti and 54 Cr become scaled solar compared to O Contribution of Type Ia SNe 3) 14 N and lot of heavy elements come from AGB stars

48. 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 from 12 C (Z=6) up to 90 Zr (Z=40) + r-elements Pre/Post SN models and explosive yields available at http://www.mporzio.astro.it/~limongi http://www.mporzio.astro.it/~limongi Alessandro Chieffi & Marco Limongi (ApJ 1998-2007) with mass loss: 11 -120 M 

49. Uncertainties in the computation PreSN Models  Extension of the Convective Core (Overshooting, Semiconvection)  Mass Loss Uncertainties in the computation of the Explosion Models  Explosion itself Piston:  Mass-cut - M ini  56 Ni (LC)  Energy (v exp)

51. Navegamos sin rumbo a través del obscuro Océano Cósmico ¿ Podemos ganar la liga de campeones ? Estrellas y planetas en un espacio infinito… ¿ Tiene sentido nuestra presencia en el Universo ? IDEAL ORCEMAN by C. Hernández

52. CHEMICAL ENRICHMENT BY A GENERATION OF MASSIVE STARS The 25 M  solar model could be considered as the “typical” case, representative of stars from 13 to 35 M  If we compute the YIELDS (ejected abundances in solar masses) of the different isotopes produced by a grid of models (~13 to 35 M  ), we could compute the chemical contribution of a generation of Massive Stars to the ISM These YIEDS are ingredients in a Chemical Evolution Model for the Galaxy, includes SFR, IMF & Infall In principle, the chemical solar distribution is a consequence of different generations of stars with different initial compositions The metallicity of the ISM is expected to increse continously and with longer time- scales than the evolutionary time of the stars that contributes to the chemical enrichment We expect that the YIELDS of a generation of masive solar metallicity stars explain the solar distribution

53. The only elements that vary between case “Flat” and case “Trend” are Fe and Ni and, at a smaller extent also Ti, Co and Zn (i.e. elements produced in the deep layers of the exploding mantle) The majority of the elements have PFs compatible with that of O  show a scaled solar distribution Flat 56 Ni => 0.05 M  13 15 20 25 30 35 M   processes Int. Mass Stars Trend 56 Ni => 0.15–0.10–0.075-0.07–0.05–0.05 M  Integrated Yields adopting a different M i -M( 56 Ni) relation

54. The Final Fate of a Massive Star with mass loss: 11 -120 M  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  ) Limongi & Chieffi, 2007

55. Individual Yields Different chemical composition of the ejecta for different masses

56. Averaged Yields Yields averaged over a Salpeter IMF 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) M rem =0.186

57. Observed M Pro smaller than LC models predict Li et al. Smartt et al. van Dyk et al.

58. Initial Mass Function m u ~ 100 M  ; m l ~ 0.1 M  m rem  Stellar evolution IMF ≈ Present Day MF for massive stars IMF...universal?

59. Definitions AMU (atomic mass unit, m u )  1/12 mass of 12 C m u c 2 = 931.478 MeV Cross section: Probability per pair of particles of occurrences of a reaction   cm 2   cm 3 /s