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Supernovae and Neutron Stars Outline of today’s lecture Lecture 17: •Finish up lecture 16 (nucleosynthesis) •Supernovae •2 main classes: Type II and Type I Supernovae and •Their energetics and observable properties Neutron Stars •Supernova remnants (pretty pictures!) •Neutron Stars •Review of formation http://apod.nasa.gov/apod/ •Pulsars REVIEWS OF MODERN PHYSICS, VOLUME 74, OCTOBER 2002 The evolution and explosion of massive stars S. E. Woosley* and A. Heger† Department of Astronomy and Astrophysics, University of California, Santa Cruz, California 95064 T. A. Weaver Lawrence Livermore National Laboratory, Livermore, California 94551 (Published 7 November 2002) Like all true stars, massive stars are gravitationally confined thermonuclear reactors whose composition evolves as energy is lost to radiation and neutrinos. Unlike lower-mass stars (M Շ8M᭪), however, no point is ever reached at which a massive star can be fully supported by electron degeneracy. Instead, the center evolves to ever higher temperatures, fusing ever heavier elements until a core of iron is produced. The collapse of this iron core to a neutron star releases an enormous amount of energy, a tiny fraction of which is sufficient to explode the star as a supernova. The authors examine our current understanding of the lives and deaths of massive stars, with special attention to the relevant nuclear and stellar physics. Emphasis is placed upon their post-helium-burning evolution. Current views regarding the supernova explosion mechanism are reviewed, and the hydrodynamics of supernova shock propagation and ‘‘fallback’’ is discussed. The calculated neutron star masses, supernova light curves, and spectra from these model stars are shown to be consistent with observations. During all phases, particular attention is paid to the nucleosynthesis of heavy elements. Such stars are capable of producing, with few exceptions, the isotopes between mass 16 and 88 as well as a large fraction of still heavier elements made by the r and p processes. CONTENTS 1. Carbon burning 1032 2. Neon burning 1032 3. Oxygen burning 1033 I. Introduction 1016 4. Silicon burning 1034 II. Presupernova Evolution—General Features 1016 5. Nuclear statistical equilibrium 1035 A. Physical overview 1016 B. Stellar models 1035 B. Equation of state and initial composition 1017 1. 8M᭪ to 11M᭪ 1035 C. Opacities 1017 2. 11M᭪ to 100M᭪ 1037 D. Neutrino losses 1018 C. Role of weak interactions 1038 E. Convection 1018 D. Effects of rotation in the late stages 1040 1. Semiconvection 1020 E. Magnetic fields 1042 2. Overshoot mixing 1021 F. Effect of metallicity on the presupernova model 1044 F. Rotation 1021 V. Core Collapse and Explosion 1045 G. Mass loss 1024 A. The iron core 1045 1. Single stars 1024 B. Collapse and bounce 1046 2. Mass loss in binaries 1025 C. Neutrino energy deposition and convection; the III. Main-Sequence and Helium-Burning Evolution 1026 shock is launched 1047 D. Shock propagation and mixing 1050 A. Nuclear physics 1026 VI. Neutron Stars and Black Holes 1050 1. Hydrogen burning 1026 A. Fallback during the explosion 1051 2. Helium burning 1026 B. Fate of ‘‘failed’’ supernovae 1051 B. Observational diagnostics of hydrogen and VII. Pair-Instability Supernovae 1052 helium burning 1027 VIII. Nucleosynthesis Resulting from Gravitationally 1. Red-to-blue supergiant ratios 1027 Powered Explosions 1053 2. SN 1987A 1027 A. Conditions for explosive nucleosynthesis 1053 C. Nucleosynthesis during hydrogen burning 1028 B. Explosive processes 1054 D. Nucleosynthesis during helium burning 1028 1. Explosive oxygen and silicon burning 1054 1. Carbon and oxygen 1028 2. Explosive neon and carbon burning 1054 2. 18O, 19F, and 21,22Ne 1029 3. The p process 1055 E. The s process 1029 4. The neutrino process 1056 IV. Advanced Nuclear Burning Stages 1031 5. The r process 1056 A. General nuclear characteristics 1032 C. Reaction-rate sensitivity 1058 D. The effects of metallicity 1058 E. Nucleosynthesis summary 1058 *Electronic address: [email protected] 1. Processes and products 1058 †Also at Enrico Fermi Institute, University of Chicago, 5640 2. Gamma-ray lines and meteorite anomalies 1060 S. Ellis, Chicago, IL 60637. Electronic address: IX. Light Curves and Spectra of Type-II and Type-IB [email protected] Supernovae 1061 0034-6861/2002/74(4)/1015(57)/$35.00 1015 ©2002 The American Physical Society All heavy elements are formed through The “s-process” fusion processes in stars. The s-process occurs in AGB stars, during Neutrons are produced in the following reactions Nucleosynthesis helium shell burning Process Location Products Heavy elements are formed by the 22Ne + 4He 25Mg + n Big bang addition of neutrons, one at a time, followed big bang H, He, Li 13C + 4He 16O + n nucleosynthesis by beta decay (neutron decays into proton) Hydrogen All MS stars He burning 56Fe + n 57Fe + γ 57 58 Post-MS stars Fe + n Fe + γ Helium burning C, O M>0.5 M 58Fe + n 59Fe + γ ⊙ 59 59 - Post-MS stars Elements up to Fe Co + e + γ C, O, Si burning gray = big bang, red = fusion, brown M>8 M iron 59Co + n 60Co + γ = s-process, teal = r-process ⊙ 60Co 60Ni + e- + γ s-process AGB stars Heavy elements 60 61 (s for “slow”) M<8 M up to bismuth Ni + n Ni + γ ⊙ r-process Type II Remaining heavy (r for “rapid”) supernovae elements etc….. to lead & bismuth The “s-process” The “r-process” The r-process occurs during a Type II supernova explosion s stands for “slow”, because the time between neutron captures (~100 years) is long compared to the time for beta decay (~1 minute) Heavy elements are formed by the rapid addition of neutrons to nuclei (relative to the beta decay timescale) the s-process takes thousands of years to form the heavy elements A nucleus can gain lots of neutrons quickly before undergoing beta decay The s-process occurs at the base of the convective envelope, deep in the star’s interior Optimal conditions for the r-process r-Process Site #1: The Neutrino-powered Wind To build up heavy elements using the r-process, you need lots of neutrons bombarding the nuclei --> very high neutron densities are required During the supernova explosion, temperatures and densities are sufficiently high for the r-process to proceed n + p; Xn > Xp 4He + n For example, at Iron plus neutrons T = 2.5x109 K, 27 -3 nn ~ 10 cm or about a kilogram of neutrons Temperature per cubic cm. Wind, 1 - 10 seconds Duncan, Shapiro, & Wasserman (1986), ApJ, 309, 141 Neutron density Woosley et al. (1994), ApJ, 433, 229 r-Process Site #2: Merging neutron stars Supernovae: Observations SUPERNOVAE • A supernova is the explosive death of a star. Unlike an ordinary nova, it does not repeat. SN 1998dh Two types are easily distinguishable by their spectrum. SN 1998aq • SN 1998bu Type II has hydrogen (Hα). Type I does not. For several weeks a supernova’s 42 • Very luminous. Luminosities range from a few times 10 luminosity rivals that of a large erg s-1 (relatively faint Type II; about 300 million L ) sun . 43 -1 galaxy to 2 x 10 erg s (Type Ia; 6 billion Lsun) - roughly as bright as a large galaxy. (Recently some rare supernovae have been discovered to Supernovae are named for the year in be even brighter) which they occur + A .. Z, aa – az, ba – bz, ca – cz, etc HST Currently at SN 2014R SN 1994D SPECTROSCOPICALLY the sun Dead white dwarf Type II SN has hydrogen (Hα). Type I does not. Planetary White dwarf Nebula in a binary Type Ia Star Supernova Type II Supernova or Type Ib, Ic Supernova Apparent magnitude evolved massive star neutron star black hole Light Curve of Type IIp Supernovae Theoretical light curve of a Type IIp supernova • The most common kind of supernova. Death of a massive breakout SN IIp star that still has its hydrogen envelope. The star is a red (usually) or blue (rarely) supergiant when it dies. • There are three stages – shock breakout, the “plateau”, and Plateau the decline (or tail). (H-envelope recombines) Tail breakout SN IIp ( 56Co decay) 60 days 180 days Plateau (H-envelope • Breakout is the first time the supernova brightens. The recombines) shock wave erupts from the surface heating it to about 200,000 K for about 2000 s. It declines to 30,000 K after Tail one day. Meanwhile the luminosity declines from about 1011 solar 9 ( 56Co decay) luminosities to about 10 solar luminosities. This emission, in UV, has been seen in at least two supernovae less than one day after their explosion (Gezzari et al 2008) Light Curve of Type IIp Supernovae (cont’d) breakout SN IIp Plateau (H-envelope recombines) Tail uv flash poorly ( 56Co decay) resolved - not as big as galaxy! 60 days 180 days • As the hydrogen envelope expands and cools, it eventually reaches 5500 K where the hydrogen starts to recombine. This recombination moves into the expanding envelope as a A sequence of ultraviolet images released in June 2008 shows shock break out. wave over a period of about 3 months. The recombination Just before the explosion, the host galaxy (top left) appears relatively quiet. reduces the opacity and lets out the energy deposited by Then a bright ultraviolet flash signals the onset of the supernova the shock as it passed through the envelope. This is the plateau. The temperature stays pegged to 5500 K. Light Curve of Type IIp Supernovae (cont’d) Type IIp Supernovae (cont’d) breakout SN IIp • The spectrum is dominated by the Balmer lines of hydrogen. On the plateau the spectrum is dominantly absorption lines, but at late time as the supernova Plateau (H-envelope becomes a nebula, one sees emission lines recombines) Radii inferred on the plateau are about 1015cm (100 AU).
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