1. 1/34 PRESUPERNOVA EVOLUTION AND EXPLOSION OF MASSIVE STARS: CHALLENGES OF THE NEXT CENTURY Marco Limongi INAF – Osservatorio Astronomico di Roma, ITALY email: marco@oa-roma.inaf.it Alessandro Chieffi INAF – Istituto di Astrofisica Spaziale e Fisica Cosmica, ITALY email: alessandro.chieffi@iasf-roma.inaf.it

2. 2/34 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 events Production of long-lived radioactive isotopes: ( 26 Al, 56 Co, 57 Co, 44 Ti, 60 Fe) Chemical Evolution of Galaxies

3. 3/34 OVERVIEW OF MASSIVE STARS EVOLUTION Grid of 15 stellar models: 11, 12, 13, 14, 15, 16, 17, 20, 25, 30, 35, 40, 60, 80 and 120 M  Initial Solar Composition (A&G89) All models computed with the FRANEC (Frascati RAphson Newton Evolutionary Code) release 5.050419  4 physical + N chemical equations (mixing+nuclear burning) fully coupled and solved simultaneously (Henyey)  Nuclear network very extended 282 nuclear species (H to Mo) and ~ 3000 processes (Fully Automated) (NO Quasi (QSE) or Full Nuclear Statistical Equilibrium (NSE) approximation)  Evolution followed from the Pre Main Sequence up to the beginning of the core collapse  Convective Core Overshooting: d=0.2 H p  Mass Loss: Vink et al. (2000,2001) (Teff>12000 K), De Jager (1988) (Teff<12000 K), Nugis & Lamers (2000) (Wolf-Rayet)/Langer 1898 (WNE/WCO) (Limongi & Chieffi 2006, ApJ, 647, 483)

4. 4/34         CNO Cycle Convective Core M min (O) = 14 M  t(O)/t(H burning): 0.15 (14 M  ) – 0.79 (120 M  ) MASS LOSS CORE H BURNING Core H Burning Models

5. 5/34         He Convective Core 3  + 12 C(  ) 16 O H burning shell H exhausted core (He Core) CORE He BURNING Core He Burning Models M ≤ 30 M   RSG M ≥ 35 M   BSG

6. 6/34 H He He  C,O (N) H<0.4 WNE WCO Log (T eff )>4.0 H H/He WNL He/CO M ≥ 30 M  M ≥ 35 M  M ≥ 40 M 

7. 7/34 MASSIVE STARS: MASS LOSS DURING H-He BURNING O-Type: 60000 > T(K) > 33000 WNL: 10 -5 < 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) WR : Log 10 (T eff ) > 4.0  M < 30 M סּ explode as Red SuperGiant (RSG)  M ≥ 30 M סּ explode as Blue SuperGiant (BSG)

8. 8/34 ADVANCED BURNING STAGES Neutrino losses play a dominant role in the evolution of a massive star beyond core He burning (T>10 9 K)          H-burn. shell He core He-burn. shell CO core Core burning The Nuclear Luminosity (L nuc ) closely follows the energy losses Evolutionary times of the advanced burning stages reduce dramatically

9. 9/34 MASSIVE STARS: LIFETIMES FuelT c (K)  c (g/cm 3 ) E nuc (erg/g)Estimated Lifetime (no ) Real Lifetime H4.1(7)4.76.44(18)2.2(7) yr (L  =5·10 38 ) 6.87(6) yr (L tot =5·10 38 ) He2.1 (8)7.2(2)8.70(17)1.6(6) yr (L  =9·10 38 ) 5.27(5) yr (L tot =9·10 38 ) C8.3(8)1.7(5)4.00(17)6.3(5) yr (L  =1·10 39 ) 6.27(3) yr (L tot =8·10 40 ) Ne1.6(9)2.5(6)1.10(17)1.7(5) yr (L  =1·10 39 ) 190 days (L tot =2·10 43 ) O2.1(9)5.8(6)4.98(17)7.9(5) yr (L  =1·10 39 ) 243 days (L tot =9·10 43 ) Si3.5(9)3.7(7)1.90(17)3.0(5) yr (L  =1·10 39 ) 19 days (L tot =2·10 45 )

10. 10/34 ADVANCED BURNING STAGES Four major burnings, i.e., carbon, neon, oxygen and silicon. Central burning  formation of a convective core Central exhaustion  shell burning  convective shell Local exhaustion  shell burning shifts outward in mass  convective shell

11. 11/34 ADVANCED BURNING STAGES The details of this behavior (number, timing, overlap of convective shells) is mainly driven by the CO core mass and by its chemical composition ( 12 C, 16 O) CO core mass Thermodynamic history 12 C, 16 O Basic fuel for all the nuclear burning stages after core He burning At core He exhaustion both the mass and the composition of the CO core scale with the initial mass

12. 12/34 ADVANCED BURNING STAGES In general, one to four carbon convective shells and one to three convective shell episodes for each of the neon, oxygen and silicon burnings occur. The number of C convective shells decreases as the mass of the CO core increases (not the total mass!)....hence, the evolutionary behavior scales as well

13. 13/34 PRESUPERNOVA STAR The complex interplay among the shell nuclear burnings and the timing of the convective zones determines in a direct way the final distribution of the chemical composition... 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 S, 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

14. 14/34 In general the higher is the mass of the CO core, the more compact is the structure at the presupernova stage PRESUPERNOVA STAR The final Fe core Masses range between: M Fe =1.20-1.45 M  for M ≤ 40 M  M Fe =1.45-1.80 M  for M > 40 M ...and the density structure of the star at the presupernova stage

15. 15/34 The most recent and detailed simulations of core collapse SN explosions show that: THE SUPERNOVA PROBLEM the shock still stalls  No explosion is obtained the energy of the explosion is a factor of 3 to 10 lower than usually observed Work is underway by all the theoretical groups to better understand the problem and we may expect progresses in the next future The simulation of the explosion of the envelope is mandatory to have information on: the chemical yields (propagation of the shock wave  compression and heating  explosive nucleosynthesis) the initial mass-remnant mass relation

16. 16/34 EXPLOSION AND FALLBACK Fe core Shock Wave Compression and Heating Induced Expansion and Explosion Initial Remna nt Matter Falling Back Mass Cut Initial Remna nt Final Remnant Matter Ejected into the ISM E kin  10 51 erg Piston (Woosley & Weaver) Thermal Bomb (Nomoto & Umeda) Kinetic Bomb (Chieffi & Limongi) Different ways of inducing the explosion FB depends on the binding energy: the higher is the initial mass the higher is the binding energy

17. 17/34 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 NL00 WIND THE FINAL FATE OF A MASSIVE STAR

18. 18/34 THE YIELDS OF MASSIVE STARS

19. 19/34 CONCLUSIONS. I 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: 25-30 M  WNE: 30-35 M  WNC: 35-40 M  The final Fe core Masses range between: M Fe =1.20-1.45 M  for M ≤ 40 M  M Fe =1.45-1.80 M  for M > 40 M  The limiting mass between SNII and SNIb/c is: 30-35 M  SNIISNIb/c Salpeter IMF The limiting mass between NS and BH formation is: 25-30 M  NSBH M>35 M  (SNIb/c) do not contribute to the intermediate and heavy elements (large fallback)

20. 20/34 PRESUPERNOVA EVOLUTION: Mass Loss during Blue and Red supergiant phases, and Wolf-Rayet stages Treatment of Convection: extension of the convective zones (overshooting, semiconvection), interaction mixing-nuclear burning 12 C(  ) 16 O cross section Rotation EXPLOSION: Lack of an autoconsistent hydrodynamical model (neutrino transport) Induced explosion [Explosion energy (where and how), time delay, fallback and mass cut (boundary conditions), mixing (inner and outer borders), extra-fallback, Y e variation, aspherical explosions] MAIN UNCERTAINTIES

21. 21/34 THE ROLE OF THE MASS LOSS FOR WNE/WCO IN THE ADVANCED BURNING PHASES Strong reduction of the He core during early core He burning Nugis & Lamers (2000) (NL00)Langer (1989) (LA89)

22. 22/34 LA89NL00 LA89NL00 THE ROLE OF THE MASS LOSS FOR WNE/WCO IN THE ADVANCED BURNING PHASES

23. 23/34 Final kinetic energy = 1 foe (10 51 erg) CONSEQUENCES ON THE FALLBACK

24. 24/34 No Mass Loss Final Mass He-Core Mass He-CC Mass CO-Core Mass Fe-Core Mass WNL WNE WC/WO Remnant Mass NS BH SNIISNIb/c RSG Z=Z  E=10 51 erg LA89 WIND NS THE FINAL FATE OF “LA89” MASSIVE STARS

25. 25/34 THE YIELDS OF “LA89” MASSIVE STARS NL00 LA89

26. 26/34 TREATMENT OF CONVECTION Convection is, in general, a hydrodynamical multi-D phenomenon  its inclusion in a hydrostatic 1-D stellar evolution code consititutes a great source of uncertainty Mixing-Length theory: Extension of the convective zones (stability criterion, overshooting, semiconvection)? Temperature Gradient? Interaction between nuclear burning and convective mixing? What about Mixing-Length theory for advanced burning stages of massive stars? Does it make sense?

27. 27/34 TREATMENT OF CONVECTION PRODUCTION OF 60 Fe IN MASSIVE STARS: X M T>4 10 8 produced 22 Ne,  60 Fe He 60 Fe synthesize within the He convective shell Preserves 60 Fe from destruction Brings new fuel ( , 22 Ne) Convection X M > 35 M O M He convective shell forms in a zone with variable composition

28. 28/34 TREATMENT OF CONVECTION THE MASS OF THE Fe CORE: Core Collapse and Bounce Energy Losses 1 x 10 51 erg/0.1M  Fe core Shock Wave Final Fe Core Mass Si exhausted Core 28 Si Mass Fraction M Si conv. shell Si exhausted Core 28 Si Mass Fraction M Final Fe Core Mass Si conv. shell

29. 29/34 UNCERTAINTY ABOUT 12 C(  ) 16 O (Imbriani et al. ApJ 2001) C 0.2  X( 12 C)=0.2 C 0.4  X( 12 C)=0.4 C 0.2 C 0.4

30. 30/34 THE ROLE OF ROTATION Increasing rotation GRATTON- ÖPIK CELL OBLATENESS Von Zeipel Theorem Cells of Meridional Circulation Advection of Angular Momentum Shear Instabilities: - Mixing of chemical species - Transport (diffusion) of angular momentum

31. 31/34 THE ROLE OF ROTATION How include this multi-D phenomenon in a 1-D code? Cylidrical Symmetry: Isobars/Equipotentials 1D problem Average values over characteristic surfaces

32. 32/34 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 How much energy to inject and where: Thermal Bomb (Internal Energy) Kinetic Bomb (Initial Velocity) Piston (Initial velocity and trajectory) How much kinetic energy at infinity [SN(~10 51 erg)/HN(~10 52 erg)] Efficiency of -process and changing of Y e Extension and timing (before/after fallback) of mixing

33. 33/34 INDUCED EXPLOSION Hypernova model without mixing-fallback (25M , E 51 =10) Hypernova model with mixing-fallback (25M , E 51 =10) Normal SN model (25M , E 51 =1) Hypernova model (25M , E 51 =10, mixing-fallback, Y e )

34. 34/34 STRATEGIES FOR IMPROVEMENTS Convection : hydrodynamical simulations in 3D  derive simple prescriptions to be used in 1D hydrostatic models (Arnett) Rotation : implementation of 3D stellar models 12 C(  ) 16 O : ask to nuclear physicists Explosive Nucleosynthesis and Stellar Remnants : solve the “Supernova Problem”  improve the treatment of neutrino transport