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The Physics of Supernovae Inma Domínguez Universidad de Granada Santiago de Chile, octubre de 2007.

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Presentation on theme: "The Physics of Supernovae Inma Domínguez Universidad de Granada Santiago de Chile, octubre de 2007."— Presentation transcript:

1 The Physics of Supernovae Inma Domínguez Universidad de Granada Santiago de Chile, octubre de 2007

2 SN 1987A  Chemical Evolution  Cosmology  Trigger Star formation  Neutrinos  BH, NS, GRBs  Reionization of the Universe etc etc

3 Supernovae are one of the most energetic explosive events in Nature BRIGHT A SN in 10 sec releases 100 times the energy that the sun releases in all its life SN1054 was as luminous as the moon for some days RARE: About 1 per century in our Galaxy Last recorded seen by naked-eye :1006 (Lupus), 1054 (Chinese), 1572 (Brahe), 1604(Kepler) BRIEF: Luminosity falls by a factor of 100 in 4 months

4 Standard Candles Fainter  Further  Distance Modulus  Luminosity Distance

5 SNe Classification Core collapse of massive stars Thermonuclear explosion I b (strong He) I c (weak He) SNe II P Type II II L No H H Type I I a (strong Si) Based on spectra and light curve morphology

6 Basic SN type spectra

7 Light Curves Type Ia SN Similar luminosity Similar spectral evolution  Good distance indicators Cosmological parameters Type II SN Dramatic differences II-P (plateau) II-L (rapid declination) Cosmology

8 SNe RATE Galaxy IaIb/cII E-S00.04< 0.01 S0a-Sb S0c-Sd Irr Mannucci et al SN rate per unit Mass ( M  yr (H o /75) 2

9 SN Ia in E-S0 Old populations Long lived progenitors Low mass  in Binary Systems SN Ib/c & SNII Absent in E-S0 Young populations Short lived progenitors Massive  SN Ia rate in Spirals Galaxies-with SFR Part of SN Ia comes from a younger population Cappellaro et al 2003, Mannucci et al. 2005, Sullivan et al. 2006

10 Stellar Evolution M<0.8 M  0.8100 M   Myr  Gyr 0.5

11 Classification of SNe ~ 4000 SNe (nowadays > 300 /yr)

12 The most abundant isotopes:  1 H  4 He  16 O 12 C 20 Ne (  -elements) O 16 O 4 He 1H1H 12 C 20 Ne 56 Fe N=50 N=82 N=126 Solar System Abundances

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14 50 yrs !!

15 Origin of the Elements: Inside the Stars Observational Evidences:  Pop II  Less heavy elements by a factor of 100 Our Galaxy has synthesized 99 % of the heavy elements during ¡ts evolution  Merril (1952) discovered Tc in  All Tc isotopes decay  1/2  10 6 yr Tc has been synthesized inside the star

16  Klein, Beskow & Treffenberg (1947) Studied the abundances at NSE in function of T and  rate nuc. re. = inverse rate This mechanism could not reproduce the observed abundances  But NOT bad for the Fe peak !! Origin of the Elements: Nuclear Statistical Equilibrium (NSE) ?

17 Binding Energy per nucleon BE/c 2 =[Zm p + (A-Z)m n - m(A,Z)] © Rolfs & Rodney 1988 BE/A 56 Fe smallest mass per nucleon  to 56 Fe exothermic reactions

18 The interpretation of the abundances  The Peaks in the abundances of 4 He, 12 C, 16 O, 20 Ne and other  elements  capture nuclear reactions inside the stars  Fe-peak elements 56 Fe is the isotope with higher binding energy 56 Fe is the last product of exothermic nuclear fusion reactions, NSE  Elements heavier than Fe High Coulomb barrier for charge reactions Neutron captures

19 Most abundant nuclei Anders & Grevesse 1989  Nuclear Physics  Physical Conditions  Where & When ??

20 Solar System Abundances Abundances peak at the “magic numbers”,Z: 2, 8, 20, 28, 56, 82 He, O, Ca, Fe, Ba, Pb © Cameron 1982

21 The familiar picture  H burning (the most effective, with an average of 7MeV per nucleon of generated energy): produced 4He, 3He, and gives (generally secondary) contributions to intermediate nuclei up to Si.  He burning (the second-most effective): produces 12 C, 16 O, some 20 Ne, plus secondary chains starting from 14 N or 13 C and leading to neutron generation.  Fusion of intermediate nuclei - 12 C, 16 O, 20 Ne, 28 Si  nuclei below and up to the Fe-peak.  Nuclear statistical equilibrium (NSE) processes, crossing the peak at 56 Fe - 56 Ni.  Explosive nucleosynthesis, starting from NSE and reorganizing abundances up to 65 Cu, occur in CCSNe and in SN Ia.  Neutron captures (slow and rapid – s and r - processes).

22 Solar System Abundances Anders & Grevesse 1989 Cameron 1982 SNII BBN SNII SNIa AGB SNII ? AGB BBN

23 Some definitions… “ Metals”: elements heavier than helium, Z “Metallicity”: [Fe/H] = log (Fe/H)  – log (Fe/H)  “Abundance ratio”: [X/Y]= log (X/Y)  – log (X/Y)  * Abundance scale by number: 12  log N(H) * Mass fractions: X= Hydrogen (X  ~0.71) Y= Helium 4 (Y  ~0.27) X+Y+Z= 1 Z= Metals (Z  ~0.02)  Population I objects (stars): Z ~ Z   Population II : Z << Z   Population III : Z ~ 0 (not detected yet ?)

24 Stellar Evolution & Nucleosynthesis Mass  AGB Planetary Nebulae White Dwarfs (if) Binary Systems  Novae  SNe Ia  AIC: Neutron  (Pulsars)  Neutron  (Pulsars)  Black Holes  CCSNe The activation of a nuclear burning phase The stellar life-time DEPEND on “ Less” in Z…

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26 Low mass stars M < 8 M  AGB/Planetary Nebulae return C, N, s-elements etc to the ISM

27 Exploding CO WDs (accreting mass from a companion) SN Ia produce ~2/3 of the observed Fe in the Universe Type Ia Supernovae (SN Ia or Thermonuclear SNe)

28 25 M  Massive  Chieffi, Limongi, Straniero 1998

29 Massive stars M ≥ 8-10 M  Core Collapse Supernovae eject O, Mg, Ti and likely r-p-elements into the ISM

30 BB = Big Bang; CR = Cosmic Rays; neut. = ν induced reactions in SNII; IMS = Intermediate Mass Stars; SNII = Core collapse supernovae; SNIa = Thermonuclear supernovae; s-r = slow-rapid neutron captures Origin of the elements

31 The Origin of the Elements up to Zn ApJS 1995 L * M < 8M  neut. Irra CR Cosmic Rays s shell x Explosive  rich freeze out

32 Yields  Low and Intermediate Mass Stars 4 He C N s-process (A > 90) elements Lattanzio et al., Meynet & Maeder, Marigo et al., Siess et al. Straniero et al. (TERAMO), Siess et al., Van den Hoeck & Groenewegen Ventura et al.  Type Ia Supernovae Fe and Fe-peak Nomoto et al., Iwamoto et al. Höflich et al., Thielemann et al.  Massive stars   -elements (O, Ne, Mg, Si, S, Ca),  some Fe-peak, s-process elements (A < 90)  and r-process elements. Woosley & Weaver / Limongi & Chieffi (ORFEO)

33 Some definitions  Yields  Production Factor Mass Loss !! in M 

34 Yields + Evolution-Time  Chemical Evolution time SN II SN Ia + SNII Chemical Evolution  -elements Fe  20 Ne 24 Mg 28 Si 32 S 36 Ar 40 Ca

35  -enhancements appear naturally due to the different life-times between SNII and SNIa… but at what level? and when? Modification of the IMF: more massive stars produce more “alphas” Modification of the SFR: more “alphas” produced before SNIa appear © McWilliam (1997)

36 Ingredients of GCE Initial conditions Big Bang abundances Prompt initial enrichment Initial mass function (IMF) Relative birthrates of stars with different masses Star formation rate (SFR) Constant, burst, interruptions etc Stellar yields vs. stellar mass and metallicity SNII, SNIa, AGB, Novae, etc Galactic gas inflow/outflow Late infall of primordial gas etc Supernova-driven galactic winds etc Stellar & gas dynamics

37 STELLAR EVOLUTION EQUATIONS 1 Dimension Lagrangian Hydrostatic + Chemical Evolution

38 STELLAR EVOLUTION EQUATIONS Convection (a problem !!)  Time-dependent convection  Mixing-Nuclear burning coupled Micro-physics  EOS  Opacity  Nuclear Cross Sections (Strong & Weak)  Screening factors  Neutrinos

39 Extensive Nuclear Networks  Automatic Adaptive Network (p,  ) ( ,n) (,)(,) ( ,p) (p,n) (p,  ) (n,  ) (n,p) (n,  ) (  n) (  p) (  )         NUCLEAR NETWORK High number of Isotopes High Number of Nuclear Reactions p, n and  captures e ± captures  ± Decay

40 THE FRANEC CODE MAIN PROGRAM (Finite difference Henyey Method) Strong reactions Weak reactions Neutrinos Opacities Equation of State Initial stellar parameter (mass, chemical composition) First model at the beginning of the Pre-MS Definition of Convective borders Mixing Adaptive re-zoning Mass loss Atmosphere Physical evolution Chemical evolution New temporal step Output

41  AGB  Thermonuclear SNe  Core Collapse SNe

42 Evolution of Low & Intermediate Mass Stars

43 Schematic structure of an AGB star (not to scale)

44 Evolutionary track toward the WD 0.6 CO 0.55 He 0.2 CO 0.1 He 0.5 He 0.6 CO WD MS RGBHB AGB PN M=1 M  t =10 Gyr Remnant: CO WD 0.6 M  Prada Moroni & Straniero 2002

45 A WD in a binary system toward a thermonuclear explosion 2 WDs WD + 

46 Light Curve L time 56 Ni  56 Co  56 Fe Thermonuclear Explosion of a CO WD M~M Chandrasekhar L max  M Ni ~ 1.4 M  “Universally” accepted model for Ia: Supernova Cosmology Project

47 WD is degenerate Pressure for relativistic electrons:  1926 Fowler  Pauli Exclusion Principle P independent of T Thermonuclear Explosion e - Degenerate Pressure (EOS) The Chandrasekhar limit   nuc <  hyd

48 Thermonuclear Explosions C-deflagration C or He detonation C-delayed detonation RG WD SD DD Detonation v burn  v sound Deflagration v burn < v sound Delayed detonation Deflagration  Detonation Propagation of the burning front M Ch Compressional heating WD ignition

49 Still Key Problems to control SNIa !!  Progenitors ? CO WD + companion SD vs DD… both ?? Accretion ?? 1D parametrization 3D still … fighting !! (Barcelona, Chicago, MPI, NRL) begin subsonic  Explosion Mechanism ? CSM : 2002ic Hamuy et al. Nature gj Aldering et al X Patat et al. Science 2007 NORMAL SNIa

50 Massive  Core Collapse At the end... Layered Structure Dense Iron Core   10 7 g·cm -3 T  K M Core  1.4M  R Si-Core  4000 km R Fe-Core  800 km

51 Massive  Core Collapse FusingMain Fusion Products Time H He 6 million years He C, O years C Ne, O 1000 years Ne O 9 Months O S, Si, Ar 4 Months Si Fe, Cr 1 day End result ? A star whose core looks like an onion

52 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” M=25M  Chieffi & Limongi Collapse and Explosion

53 Core-Collapse Mechanism Once the star has finished its fuel the core cools because of two reasons : c) Contraction turns into a free-fall collapse, vast amount of neutrinos are produced In less than 1 second the inner core radius goes from 4000 km to 10 km (matter from the rest of the core is falling inward) a)Iron dissociation  fusion of light nuclei  the star continues emitting energy b)Degenerate e - gas  p + e - (2.25 MeV)  n + e (neutronization)  e escape and remove energy

54 Core-Collapse Mechanism Making Stars Explode PROBLEM: Turning the implosion into an explosion !!! There are several models explaining the explosion, but until now simulations do not succeed in obtaining an explosion Because the neutrinos free path is small the falling matter becames very hot and expands outwards. Finally, the star explodes and ejects the star’s outer layers into space. All that remains of is a very dense object: neutron star or black hole

55 Core Collapse SNe: LCs L  M 56Ni 56 Ni  56 Co  56 Fe II-P 1. Rise: thermal energy (envelope is fully ionized) 2. Plateau: recombination of H Lenght  M H 3. Radioactive Tail: 56 Co decay II-L No Plateau Small H-envelope simulated by a piston of initial velocity v 0, located near the edge of the Fe core Explosion Mechanism Still Uncertain

56 Numerical Methods STELLAR EVOLUTION FRANEC (Chieffi, Domínguez, Imbriani, Limongi, Piersanti, Straniero) 1D Hydrostatic Code  Extended Nuclear Network (700 isotopes)  Physics and Chemestry coupled  Time dependent mixing PMS  AGB  WD  Accretion  Explosive C-ignition  TPs PMS  Fe-core Low-mass  Massive 

57 Numerical Methods EXPLOSION & LIGHT CURVES 1D Radiation-Hydrodynamic Code (PPM) (Höflich, Khokhlov )   Ray transport Monte Carlo  Frequency dependent transport eq. (1000 )  Extended Nuclear Network (postprocess)  Radiation transport via moments eq.  Expansion opacities (scatt., bf, bb)  Explosion mechanism: detonation deflagration piston  CCSNe LCs Eddington fac. Mean opacities +  SNIa

58 1999ee SNIa Hamuy et al em IIP Hamuy et al el SNIa Krisciunas et al Observations L max  LC L max  B-V L max  V Ca L max  V Ni  LCs  Spectra (evolution)  Observed Relations

59 Information from the spectra -4 days + 15 days C-burning Star of Si burning Duration of these phases lower limit to the mass SN1999by MgII 1.05  m CaII 1.15  m Hoflich et al SNIa Sub-L

60 Visible X-ray IR Radio SN Remnants Crab Nebula SN 1054

61 Type Ia SN remnants: shocked ejecta Tycho SN 1572 X-ray emission spectra Interaction with the Ambient Medium  AM ~ g/cm 3   T X i ionization XMM-Newton DDT PDDT Sub-Ch Identify Explosion Mechanism DDT  Fe Ca S Fe O Si Ar Badenes et al. 2003

62 Cas A Asymmetrically expanding  Explosion ?? Age ~ 300 yr SN1680 Good spatial resolution X and Optical data  CCSNe He-rich envelope SiXIII/MgXI Vink et al Hwang et al Chandra Si Fe

63 Bibliography  BÖHM-VITENSE 1993, Introduction to Stellar Astrophysiscs, University of Chicago Press.  CLAYTON 1992, Principles of Stellar Evolution and Nucleosynthesis, University of Chicago Press.  HANSEN & KAWALER 1994, Stellar Interiors: Physical Principles, Structure and Evolution, Springer-Verlag  KIPPENHAHN 1990, Principles of Stellar Structure and Evolution, Springer-Verlag.  OSTLIE & CARROLL 1996, An Introduction to Modern Stellar Astrophysics, Addison Wesley.

64 Bibliography  PAGEL 1997, Nucleosynthesis and Chemical Evolution of Galaxies, Cambridge University Press.  BUSSO, GALLINO, WASSERBURG 1999, Nucleosynthesis in AGB stars, Ann. Rev. A. &A., 36, 369.  WALLERSTEIN et al. 1998, Synthesis of the elements in stars forty years of progress, Reviews of Modern Physics, Volume 69,


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