Transcript Slide 1
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 Email: [email protected] 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 Mup’ 8 - 10 < Massive stars < MPISN >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: 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 Log Mass Fraction 2 1 0 -1 -2 -3 -4 -5 -6 -7 -8 -9 -10 -11 -12 BB Novae SNIa 0 20 40 60 80 100 120 CR IMS s-r 140 neut. SNII 160 180 200 Atomic Weight 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 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 Computation of the Presupernova Evolution of Massive Stars 60Zn 61Zn 62Zn 63Zn 64Zn 65Zn 66Zn 67Zn 68Zn 57Cu 58Cu 59Cu 60Cu 61Cu 62Cu 63Cu 64Cu 65Cu 66Cu 67Cu 56Ni 1. Extended Network 57Ni 58Ni 59Ni 60Ni 61Ni 62Ni 63Ni 64Ni 65Ni 54Co 55Co 56Co 57Co 58Co 59Co 60Co 61Co 62Co 52Fe Including a large number of isotopes and reactions (captures of light partcles, e± captures, β± decays) 37K 53Fe 54Fe 55Fe 56Fe 57Fe 58Fe 59Fe 60Fe 61Fe 51Mn 52Mn 53Mn 54Mn 55Mn 56Mn 57Mn 48Cr 49Cr 50Cr 51Cr 52Cr 53Cr 54Cr 55Cr 41Sc 42Sc 45V 46V 47V 48V 49V 50V 51V 52V 44Ti 45Ti 46Ti 47Ti 48Ti 49Ti 50Ti 51Ti 43Sc 44Sc 45Sc 46Sc 47Sc 48Sc 49Sc 50Sc 40Ca 41Ca 42Ca 43Ca 44Ca 45Ca 46Ca 47Ca 48Ca 49Ca 38K 39K 40K 41K 42K 35Ar 36Ar 37Ar 38Ar 39Ar 40Ar 41Ar 1H 3He 4He 2H 3H 25Al 33Cl 34Cl 35Cl 36Cl 37Cl 38Cl 31S 32S 33S 34S 35S 36S 37S 29P 30P 31P 32P 33P 34P 27Si 28Si 29Si 30Si 31Si 32Si 33Si 26Al 27Al 28Al (a,n) (a,g) 23Mg 24Mg 25Mg 26Mg 27Mg n 21Na 22Na 23Na 24Na (p,n) b-,e+ 20Ne 21Ne 22Ne 23Ne 7Be 8Be 6Li 7Li 17F 18F 19F 20F 15O 16O 17O 18O 19O 13N 14N 15N 16N 12C 13C 14C 10B 11B 9Be 10Be (p,g) (n,g) (g,n) (g,p) (p,a) (g,a) (n,a) (a,p) b+,e(n,p) Computation of the Presupernova Evolution of Massive Stars 2. Strong coupling between physical and chemical evolution: P Gm m 4r 4 r 1 m 4r 2 ( P, T , Yi ) Yi ci ( j ) jY j + ci ( j, k ) N A v j ,k Y jYk t j j ,k L e nuc ( P, T , Yi ) + e ( P, T , Yi ) + e grav ( P, T , Yi ) m T GmT ( P, T , Yi ) 2 m 4r P H/He burnings: Adv. burnings: + + ci ( j, k , l ) 2 N A v j ,k ,l Y jYk Yl 2 j , k ,l i 1,........, N ( P, T ) ; e nuc e nuc ( P, T ) ; e e ( P, T ) ; e grav e grav ( P, T ) ; ( P, T ) e nuc e nuc ( P, T , Yi ) Decoupled Coupled Computation of the Presupernova Evolution of Massive Stars 3. Tratment of convection: - Time dependent convection mix t - Interaction between Mixing and Local Burning P m r m L m T m Yi t mix nuc Gm 4r 4 1 4r 2 - e n u c + e + e g ra v Gm T 2 4r P Yi Yi + t n u c t co n v Yi Yi 2 4r D m t conv m D = Diffusion Coefficient Convective Core Core H burning g g M2 Pc 4 R g CNO Cycle g g g M P T 3 T R M Tc R g g Massive Stars powered by the CNO Cycle 3 dT F 3 4acT dr rad Th 0.2 (1 + X H ) The Convective Core shrinks in mass CNO Cycle X i0 Xi A A i i i i 12C + 1H 13N + g 13N 13C + e+ + 13C + 1H 14N + g 14N + 1H 15O + g 15O 15N CN-Cycle 15N + e+ + + 1H 12C + 4He (99%) 16O + g (1%) (T 3×107 K) 16O + 1H 17F + g 17F 17O 17O + e+ + NO-Cycle + 1H 14N + 4He CNO Processed Material C N O Ne-Na and Mg-Al Cycles During Core H Burning the central temperature is high enough (3-7×107 K) that the Ne-Na and Mg-Al cycles become efficient Ne-Na Cycle 20Ne + 1H 21Na + g 21Na 21Ne Mg-Al Cycle 24Mg 21Ne + e+ + 25Al + 1H 22Na + g 22Na + 1H 25Al + g 25Mg 22Ne + e+ + + 1H 26Al + g 26Al 22Ne + 1H 23Na + g 26Mg 23Na + 1H 20Ne + 4He 27Al 25Mg 25Mg + e+ + 26Mg + e+ + + 1H 27Al + g + 1H 24Mg + 4He 21Na e 22Ne slightly burnt 23Na e 26Al 26Mg destroyed increases (~10-7) produced Evolutionary Properties of the Interior t=6.8 106 yr Evolutionary Properties of the Surface Core H Burning Models Mmin(O) = 14 M t(O)/t(H burning): 0.15 (14 M ) – 0.79 (120 M) 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 Core He burning 3a+ 12C(a,g)16O 4He g g g He Convective Core g g + 4He 8Be + g 8Be 8Be 4He + 4He + 4He 12C + g 3 4He 12C + g g g g Tc 1.5 108 K rad He C,O Bordo iniziale CC Mix He Core Convettivo H burning shell H exhausted core (He Core) The He convective core increases in mass ad Nucleosynthesis during Core He burning 3 4He 12C + g 12C + 4He 12O + g 16O + 4He 20Ne + g 20Ne + 4He 24Mg + g Chemical composition at core He exhaustion: mainly C/O 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) C/O ratio depends on: 1. Treatment of convection (late stages of core He burning) 2. 12C(a,g)16O cross section Nucleosynthesis during Core He burning 14N, produced by H burning activates the sequence of reactions: 14N + 4He 18F + g 18F 18O + 4He 22Ne + g 22Ne For the CNO cycle: 18O + e+ + + 4He 25Mg + n Xi X i0 A A i i i i X i0 A 9.3610-4 i i For e Solar composition i (12C,13C,14N ,15N ,16O,17O) X14 9.3610-4 14 1.3 10-2 XCNO(iniziale) X14N For a Solat composition at core H exhaustion: X(14N) ~ ½ Z X 14 In general: The efficiency of the 14N 1 Z 2 reactions scales with the metallicity Nucleosynthesis during Core He burning 14N 22Ne during the initial stages of core He burning X 22 X 14 22 2 10-2 Z 14 He burning 14N (~1/2 Z) 22Ne (~Z) H burning CNO (~1/2 Z) During core He burning, by the nuclear reaction: 22Ne Xn 22Ne is reduced by a factor of ~2 + 4He 25Mg + n 1 X 22 1 4.6 10-4 Z 2 22 40 s-process nucleosynthesis Neutron Mass Fraction s-process during Core He burning p 78Rb 86Kr 77Kr s 85Br 76Br 79Rb 80Rb 81Rb 82Rb 83Rb 84Rb 85Rb 80Kr 81Kr 82Kr 83Kr 84Kr 80Br 81Br 82Br 83Br 78Se 79Se 80Se 81Se 82Se 77As 78As 79As 80As 81As 75Ge 76Ge 77Ge 78Ge 79Ge 80Ge 74Ga 75Ga 76Ga 77Ga 78Ga 79Ga 87Kr 88Kr 78Kr 79Kr 86Br b- 87Br 77Br 78Br 79Br b84Se 75Se r 83As 74As 85Se 86Se 76Se 77Se 84As b- 85As 75As 76As b73Ge 74Ge 72Ga 73Ga n,g s,r 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 WIND t=5.3 105 yr Evolutionary Properties of the Surface Core Core He He Burnin Burnin gg Models 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 12C Convection + mass fraction (Treatment of 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 sprocess nucleosynthesis) All these uncertainties affect the size of the CO core that drives the following evolution 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~0.08 MeV) neutrino emission from pair production start to become very efficient g H burning shell H exhausted core (He Core) He burning shell g e+ + e- e + e g He exhausted core (CO Core) g g g g g g g Core Burning Advanced burning stages Enuc L M t nuc t nuc Enuc M L 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 Lg cost At PreSN stage L 10 10 Lg 8 10 Absolute Magnitude increases by ~25 Advanced Nuclear Burning Stages: Core C burning H He CO H burning shell He burning shell T~109 K Advanced Nuclear Burning Stages: C burning At high tempreatures a larger number of nuclear reactions are activated Heavy nuclei start to be produced C-burning T ~ 109 K Main Products of C burning 20Ne, 23Na, 24Mg, 27Al Scondary Products of C burning 22 Ne(a , n) 25Mg s-process nuclesynthesis Advanced Nuclear Burning Stages: Core Ne burning H He H burning shell He burning shell CO NeO C burning shell T~1.3×109 K Advanced Nuclear Burning Stages: Ne burning Ne-burning T ~ 1.3 109 K Main Products of Ne burning 16O, 24Mg, 28Si Scondary Products of Ne burning 29Si, 30Si, 32S Advanced Nuclear Burning Stages: Core O burning H He CO NeO O H burning shell He burning shell C burning shell Ne burning shell T~2×109 K Advanced Nuclear Burning Stages: O burning O-burning T ~ 2 109 K Main Products of O burning 28Si (~0.55) 32S (~0.24) Secondary Products of O burning 34S (~0.07) 36Ar (~0.02) 38Ar (~0.10) 40Ca (~0.01) Advanced Nuclear Burning Stages: O burning Proton Number (Z) During core O burning weak interactions become efficient 40Ca 41Ca 42Ca 43Ca 37K 38K 39K 40K 41K 42K 35Ar 36Ar 37Ar 38Ar 39Ar 40Ar 41Ar 33Cl 34Cl 35Cl 36Cl 37Cl 38Cl 31S 32S 33S 34S 35S 36S 37S 29P 30P 31P 32P 33P 34P 27Si 28Si 29Si 30Si 31Si 32Si 33Si 26Al 27Al 28Al 44Ca Neutron Number (N) Most efficient processes: 31S(b+)31P 33S(e-,)33P 30P(e-,)30Si The electron fraction per nucleon 37Ar(e-,)37Cl Zi Ye X i 0.5 i Ai Advanced Nuclear Burning Stages: Core Si burning H He CO NeO O SiS H burning shell He burning shell C burning shell Ne burning shell O burning shell T~2.5×109 K Advanced Nuclear Burning Stages: Si burning At Oxygen exhaustion Balance between forward and reverse (strong interaction) reactions for increasing number of processes T ~ 2.5 109 K j+l i+k rik rjl A measure of the degree of equilibrium reached by a couple of forward and reverse processes (ij ) rik - rjl (ij) 1 max( rik , rjl ) (ij) 0 Non equilibrium Full equilibrium Advanced Nuclear Burning Stages: Si burning At Oxygen exhaustion T ~ 2.5 109 K (ij ) 0.1 Sc At Si ignition At Si ignition (panel a + panel b) T ~ 3.5 109 K T ~ 3.5 109 K (ij ) 0.01 Si A=44 Equilibrium Equilibrium 56Fe 28Si 0.01 (ij ) 0.1 (ij ) 0.1 Partial Eq. (ij ) 0.1 Out of Equilibrium A=45 Out of Eq. (ij) 0.1 Eq. Clusters Advanced Nuclear Burning Stages: Si burning T ~ 3.5 109 K A=45 (ij) 0.1 1. 56Fe A=44 24Mg 3. The matter flows from the lower to the upper cluster through a sequence of non equilibirum reactions 20Ne 16O Equilibrium Clusters Clusters di equilibrio 4He is burnt through a sequence of (g,a) reactions 2. The two QSE clusters reajdust on the new equilibrium abundances of the light particles 28Si 12C 28Si 43 Ca (n, g ) 44 Ca 42 Ca (a , p ) 45 Sc 44 Sc (n, p ) 44 Ca 42 Ca (a , g ) 46 T i 42 Ca (a , n) 45 T i 44 T i(n, g ) 45 T i 44 Sc ( p, g ) 45 T i 41 43 Ca (a , n) 46 T i 44 41 K (a , p ) 44 Ca 4. Ye is continuosuly decreased by the weak interactions (out of equilibrium) Ca (a , g ) 45 T i Sc (n, g ) 45 Sc 56,57,58Fe, 52,53,54Cr, 55Mn, 59Co, 62Ni NSE Pre-SuperNova Stage H He CO NeO O SiS Fe H burning shell He burning shell C burning shell Ne burning shell O burning shell T~4.0×109 K Si burning shell Evolutionary Properties of the Interior H burning shell He burning shell Ne burning shell O burning shell C burning shell Si burning shell Chemical Stratification at PreSN Stage 16O,24Mg, 14N, 13C, 17O 28Si,29Si, 30Si 12C, 16O 28Si,32S, 36Ar,40Ca, 34S, 38Ar 12C, 16O s-proc 20Ne,23Na, 56,57,58Fe , 52,53,54Cr , 24Mg,25Mg, 27Al, s-proc 55Mn, 59Co, 62Ni NSE Each zone keeps track of the various central or shell burnings Main Properties of the PreSN Evolution Fase Time (yr) H 5.93(6) Lnuc L Mcc Tc c Mshell Fuel Main Prod. Sec. Prod. 12.8 3.7(7) 7.2 8.7 1H 4He 13C, 14N, 17O He 6.8(5) 6.02 1.5(8) 4.7(2) 6 4He 12C, 18O, 16O 22Ne, s-proc. C 9.7(2) 1.0(6)5.0(7) 4.0(7)1.0(9) 7.2(8) 1.2(5) 2.39 12C 20Ne, 25Mg, 23Na, s-proc. 24Mg, 27Al Ne O 7.7(-1) (280 d) 7.0(9) 3.3(-1) (120 d) 5.0(10) 5.9(11) 2.2(9) 4.0(10) 0.62 1.05 1.2(9) 1.8(9) 2.1(6) 4.0(6) 2.39 1.7 20Ne 16O 16Ne, 29Si, 24Mg 30Si 28Si, Cl, Ar, K, Ca 32S, 36Ar, 40Ca, Si 2.1(-2) (7 d) 1.1(13) 1.0(12) 1.08 3.1(9) 7.5(7) 1.5 28Si 54Fe, 56Fe, 55Fe Ti, V, Cr, Mn, Co, Ni Evolution of More Massive Stars: Mass Loss O-Type: 60000 > T(K) > 33000 Wolf-Rayet : 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) Final Masses at the PreSN stage WIND RSG WNL WNE WC/WO 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 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 enuc lowers becaus the matter is at NSE The Fe core begins to degenerate The Chandrasekhar Mass MCh=5.85×(Ye)2 M is reached Tc ~ 1010 K, c ~ 1010 K Pe ~ 1028 dyne/cm2 Pi ~ 2×1026 dyne/cm2 Prad ~ 3×1025 dyne/cm2 A strong gravitational contraction begins The Fermi energy increasesthe electron captures on both the free and bound protons incease as well The main source of pressure against gravity (electron Pressure) lowers The gravitational collapse begins Fe Core 3 1012 g/cm 3 3 1011 g/cm 3 e- + p e + n 10 g/cm 14 Fe Core diffusion -sphere Neutrino Trapping 3 Shock wave Core Bounce and Rebounce Eenergy Losses 2x 1051 erg/0.1M Stalled Shock “Prompt”shocks eventually stall! Strong Shock vs Weak Shock A strong shock propagates. Matter is ejected. A weak shock stalls. Matter falls back. Neutrino-driven explosions Stalled Shock RS=200-300 Km Energy deposition behind the stalled shock wave due to neutrino interactions: heating cooling diffusion p,n e , e e + e- e + e e + e e- + eShock Wave reheated e+,en,p Explosion Gain Radius RG=100-150 Km Neutrinosphere R=50-700 Km Explosive Nucleosynthesis Propatagiont of the shock wave through the envelope 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 Induced Shock Compressio n and Heating Induced Expansio n and Explosio n Initial Remnant Injected Energy Matter Ejected into the ISM Ekin1051 erg Matter Falling Back Mass Cut Initial Remnant Final Remnant 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 Z=Z E=1051 erg SNII SNIb/c WNL WNE WC/WO Fallback Mass (M) RSG Black Hole Neutron Star Initial 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 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 Chemical Enrichment due to Massive Stars Mto t PFi X i dm Mcu t Mto t S un X i dm Mcu t 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 Mup Yi tot Yi (m)dm (m) km-a a 2.35 Mlow 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 No room for other sources (AGB) No room for AGB Secondary Isotopes? process. Other sources uncertain Type Ia Explosion? AGB Chemical Enrichment due to Massive Stars Global Properties: IMF: Salpeter 1 M Initial Composition (Mass Fraction) X=0.695 Y=0.285 Z=0.020 NO Dilution Mrem=0.186 Final Composition (Mass Fraction) X=0.444 (f=0.64) Y=0.420 (f=1.47) Z=0.136 (f=6.84) 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: 25-30 M WNE: 30-35 M WNC: 35-40 M The final Fe core Masses range between: MFe=1.20-1.45 M for M ≤ 40 M MFe=1.45-1.80 M for M > 40 M The limiting mass between SNII and SNIb/c is : Salpeter IMF SNIb / c 0.22 SNII 30-35 M SNII SNIb/c 25-30 M The limiting mass between NS and BH formation is: (uncertainties on mass loss, simulated explosion, etc.) NS BH CONCLUSIONS Massive Stars are responsible for producing elements with 4<Z<38 Assuming a Salpeted IMF the efficiency of enriching the ISM with heavy elements is: For each solar mass of gas returned to the ISM H: decreased by f=0.64 He: increased by f=1.47 Metals: increased by f=6.84 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 http://www.mporzio.astro.it/~limongi