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PRESUPERNOVA EVOLUTION AND EXPLOSION OF MASSIVE STARS: CHALLENGES OF THE NEXT CENTURY

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

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PRESUPERNOVA EVOLUTION AND EXPLOSION OF MASSIVE STARS: CHALLENGES OF THE NEXT CENTURY

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  1. 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. WHY ARE MASSIVE STARS IMPORTANT IN THE GLOBAL EVOLUTION OF OUR UNIVERSE? Chemical 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 events

  3. 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 (Limongi & Chieffi 2006, ApJ, 647, 483) • Evolution followed from the Pre Main Sequence up to the beginning of the core collapse • 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) • Convective Core Overshooting: d=0.2 Hp • 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)

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

  5. CORE He BURNING 3a+ 12C(a,g)16O He Convective Core g g g Core He Burning Models g g M ≤ 30 M RSG g g M ≥ 35 M BSG g H burning shell H exhausted core (He Core)

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

  7. MASSIVE STARS: MASS LOSS DURING H-He BURNING WR : Log10(Teff) > 4.0 • WNL: 10-5< Hsup <0.4 (H burning, CNO, products) • WNE: Hsup<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) O-Type: 60000 > T(K) > 33000 • M < 30 Mסּexplode as Red SuperGiant (RSG) • M ≥ 30 Mסּ explode as Blue SuperGiant (BSG)

  8. ADVANCED BURNING STAGES Neutrino losses play a dominant role in the evolution of a massive star beyond core He burning (T>109 K) g n n H-burn. shell g g CO core n n Core burning He core g g n g n He-burn. shell g g n n g The Nuclear Luminosity (Lnuc) closely follows the energy losses Evolutionary times of the advanced burning stages reduce dramatically

  9. MASSIVE STARS: LIFETIMES

  10. 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. 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 (12C, 16O) CO core mass Thermodynamic history Basic fuel for all the nuclear burning stages after core He burning 12C, 16O At core He exhaustion both the mass and the composition of the CO core scale with the initial mass

  12. ADVANCED BURNING STAGES ...hence, the evolutionary behavior scales as well 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!).

  13. 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... 16O,24Mg, 28Si,29S, 30Si 14N, 13C, 17O 14N, 13C, 17O 12C, 16O 28Si,32S, 36Ar,40Ca, 34S, 38Ar 12C, 16O s-proc 20Ne,23Na, 24Mg,25Mg, 27Al, s-proc 56,57,58Fe, 52,53,54Cr, 55Mn, 59Co, 62Ni NSE

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

  15. THE SUPERNOVA PROBLEM The most recent and detailed simulations of core collapse SN explosions show that: • 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. Shock Wave Compression and Heating Matter Falling Back Matter Ejected into the ISM Ekin1051 erg Induced Expansion and Explosion Mass Cut Initial Remnant Final Remnant Initial Remnant Fe core EXPLOSION AND FALLBACK • 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. Z=Z E=1051 erg NL00 No Mass Loss WIND SNII SNIb/c RSG WNL Final Mass WC/WO WNE He-Core Mass Fallback Remnant Mass He-CC Mass CO-Core Mass Black Hole Fe-Core Mass Neutron Star THE FINAL FATE OF A MASSIVE STAR

  18. THE YIELDS OF MASSIVE STARS

  19. CONCLUSIONS. I • Stars with M<30 M explode as RSG Stars with M≥30 M explode as BSG WNL: 25-30 M • The minimum masses for the formation of the various kind of Wolf-Rayet stars are: WNE: 30-35 M WNC: 35-40 M MFe=1.20-1.45 M for M ≤ 40 M • The final Fe core Masses range between: MFe=1.45-1.80 M for M > 40 M 30-35 M • The limiting mass between SNII and SNIb/c is: SNII SNIb/c Salpeter IMF 25-30 M • The limiting mass between NS and BH formation is: NS BH • M>35 M (SNIb/c) do not contribute to the intermediate and heavy elements (large fallback)

  20. MAIN UNCERTAINTIES 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 • 12C(a,g)16O 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, Ye variation, aspherical explosions]

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

  22. THE ROLE OF THE MASS LOSS FOR WNE/WCO IN THE ADVANCED BURNING PHASES LA89 NL00 LA89 NL00

  23. CONSEQUENCES ON THE FALLBACK Final kinetic energy = 1 foe (1051 erg)

  24. Z=Z E=1051 erg LA89 No Mass Loss SNII SNIb/c RSG WIND WNL WNE Final Mass WC/WO He-Core Mass He-CC Mass CO-Core Mass BH Remnant Mass Fe-Core Mass NS NS THE FINAL FATE OF “LA89” MASSIVE STARS

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

  26. 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. He X T>4 108 produced 60Fe 22Ne, a M X M > 35 MO M TREATMENT OF CONVECTION PRODUCTION OF 60Fe IN MASSIVE STARS: 60Fe synthesize within the He convective shell Convection • Preserves 60Fe from destruction • Brings new fuel (a, 22Ne) He convective shell forms in a zone with variable composition

  28. Si conv. shell 28Si Si exhausted Core Mass Fraction Final Fe Core Mass M Si conv. shell 28Si Si exhausted Core Mass Fraction Final Fe Core Mass M TREATMENT OF CONVECTION THE MASS OF THE Fe CORE: Core Collapse and Bounce Fe core Shock Wave n n n n Energy Losses 1 x 1051 erg/0.1M

  29. UNCERTAINTY ABOUT 12C(a,g)16O C0.2 X(12C)=0.2 C0.4 X(12C)=0.4 C0.4 C0.2 (Imbriani et al. ApJ 2001)

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

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

  32. 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 • 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(~1051 erg)/HN(~1052 erg)] • Extension and timing (before/after fallback) of mixing • Efficiency of n-process and changing of Ye

  33. INDUCED EXPLOSION Hypernova model without mixing-fallback (25M, E51=10) Normal SN model (25M, E51=1) Hypernova model (25M, E51=10, mixing-fallback, Ye) Hypernova model with mixing-fallback (25M, E51=10)

  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 • 12C(a,g)16O: ask to nuclear physicists • Explosive Nucleosynthesis and Stellar Remnants: solve the “Supernova Problem”  improve the treatment of neutrino transport

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