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Lecture 1 Overview Time Scales, Temperature-Density Scalings Approximate Author Title Comments price Good overall introduction to Stellar Structure Free O. R. stellar evolution at the upper and Evolution (pdf available) ASTRONOMY 220C Pols division undergraduate level. Stellar Structure Free ADVANCED STAGES OF STELLAR EVOLUTION Collins Web based graduate level text AND NUCLEOSYNTHESIS and Evolution (on line) Kippenhahn and Stellar Structure $88.67 Great “introductory” textbook Winter, 2019 Weigert and Evolution (hardcover) (more for 220A though) Excellent recent (2017) book on just Branch and Supernova the supernova par of the course. http://www.ucolick.org/~woosley $109.81 Wheeler Explosions Expensive. Can be “rented” for half price. A classic. Great on nuclear physics Principles of Stellar Don $<40.85 and basic stellar physics. Good on Evolution and Clayton (paperback) the s-process, but quite dated Nucleosynthesis otherwise. Used $10.00. Buy it. I. Preliminaries Stars are gravitationally confined thermonuclear reactors. The life of any star is a continual struggle between Lecture 1 the force of gravity, seeking to reduce the star to a point, and pressure, which holds it up. A balance is maintained. So long as they remain non-degenerate and have not Overview encountered the any instabilities, overheating leads to expansion and cooling. Cooling, on the other hand, leads to contraction and heating. Hence stars are stable. Time Scales, Temperature-density The Virial Theorem works. Scalings, Critical Masses But, since ideal gas pressure depends on temperature, stars must remain hot. By being hot, they are compelled to radiate. In order to replenish the energy lost to radiation, stars must either contract or obtain energy from nuclear reactions. Since nuclear reactions change their composition, stars must evolve. 18 Theory of supernovae 3.2. Final fate as a function of stellar mass As seen from eq. (3), smaller-mass stars have higher interior densities than larger-mass stars for the same central temperature and thus tend to form electron degenerate cores at earlier stages of evolution. The Virial Theorem implies that if a star (with constant In fig. 3, the earlier evolution of the central density, Pc, and temperature, T o is shown for helium stars density in this example) is neither too degenerate nor too relativistic Central Conditions at death the (radiation or pair dominated) (or cores) of mass M = 8, 6, 3.3, 3.0, 2.8, and 2.2 M e [27]. The approximate ignition lines for carbon, iron cores of 1 3GM 2 neon, oxygen, and silicon, where the nuclear energy generation rate is equalmassive to neutrino stars energy losses, ∼2MN kT (for an ideal ionized hydrogen, constant T ) 2 5R A are shown. The line for $ = 10 approximately separates the electron degenerateare somewhat and nondegenerate 3GM degenerate T ∼ regions, where ~ is the chemical potential of an electron in units of kT. These evolutionary paths clearly 20N AkR 1/3 indicate the effect of electron degeneracy and its dependence on stellar mass. ⎛ 3M ⎞ R ∼ for constant density ⎝⎜ 4πρ ⎠⎟ Helium stars of M~ = 4-8 M e undergo nuclear burning under nondegenerate conditions. In contrast, 1/3 2/3 2/3 the 2.2 M o helium core enters the strongly degenerate region after carbon burning. Stars of M~ = ⎛ 3 ⎞ ⎛ 4π ⎞ GM 1/3 GM 1/3 So T∼ ⎜ ⎟ ⎜ ⎟ ρ = 0.093 ρ ⎝ 20⎠ ⎝ 3 ⎠ N Ak N Ak 2.8-3.3 M e are intermediate: the O + Ne + Mg core becomes semi-degenerate and whether they will This is about 5 MK for the average temperature of the sun, while the enter the degenerate or nondegenerate region is an interesting question. The final stages of evolution central temperature is about 2 to 3 times greater. are classified by the stellar mass as follows (see refs. [34, 25, 45] for reviews and references). If the central temperature has the same sort of scaling (1) Mm, <0.08 Me: The star will become a planet-like blackρ ∝ Tdwarf3 without igniting hydrogen as the average, 2/3 burning. GM 1/3 Tc ∼ ρc N Ak (2) 0.08 M O < Mms < 0.45 Mo: The star will end up as a helium white dwarf, though such a single star has not evolved off the main sequence in a Hubble time. That is as stars of ideal gas contract, they get hotter and since That is, as a star of given mass evolves, its central temperature a given fuel (H, He, C etc) burns at about the same temperature, (3) 0.45 M O < Mms < 8 M®: The star forms a C-O core of mass smaller than 1.06 M e, which then rises roughly as the cube root of its central density more massive stars will burn their fuels at lower density, i.e., becomes strongly degenerate. Most of these stars will become C-O while dwarfs by losing their −2 3 higher entropy. ρ ∝ M T hydrogen-rich envelope, but some (6-8 Me) could reach a supernova stage by increasing the C-O core mass to the Chandrasekhar mass and igniting carbon deflagration. The explosion may look like a Type II-L supernova [9, 40]. (4) 8 M@ < Mms < 10 Mo: The star undergoes nondegenerate carbon burning and forms a degener- ate O + Ne + Mg core [23, 26]. The mass of the O + Ne + Mg core (<1.37 Me) is too small to ignite More generally for helium cores of constant M ≈ 10 − 25 M Burning Processes neon. Further evolution is due to the growth of the mass of the O +Z AMNeS + Mg core⊙ toward the mass, 2.2, 2.8, 3.0, 3.3, 6 and 8 MO (Nomoto and Hashimoto 1988) / I I I I I I I I _ _ I photo disintegration IU.L} F 7" < 415 / I G R /z../H d/'/,/H~-LU'~ I ,IX'- / ~ =~ 95L I/,Y , Si v • "/Ill/In ............... :', .......... I o .- s.s c,." 2 I e+e-pair ,,,~Ne '8"'~"'<Z-~.,~',~3.0 EN,:% / I~ ..... ~ ........ '.',: ~ ~ ...... ~ ...................../ "o:'v"""" ............./"/ 8. FJ "".... / Ff / I" ~k: IO 24,,L rzo. 8hi i i ' I I I I Mg I Ne I "--3 4 5 6 7 8 9 I0 II log Pc (gcm'3) Fig. 3. The evolution of the central density,It turns Pc, out and that temperature, MHe =35 MTo,O iswill shown just forbrush helium the starse+e -ofpair mass instability M s = 8, 6, 3,3, 3.0, 2.8, and 2.2 M e. The ignition lines for carbon, neon, oxygen, and silicon, where the nuclear energy generation rate is equal to neutrino energy losses, are shown. The line for ¢, = 10 approximately separates the electron degenerate and nondegenerate regions, where ~b is the chemical potential of an electron in units of kT. These instabilities (cross hatched lines) have dramatic consequences for the star: Supernovae are powered by one of two sources: • Pair instability can lead to pulsations (pulsational pair- instability supernovae), explosion (pair-instability supernovae), or collapse to a black hole strong • Thermonuclear - white dwarf explosions and force • Electron capture can rob the core of pressure support pair instability and cause collapse to a neutron star – resulting in a supernova) • Gravitational collapse – aka “core collapse” - gravity the a fraction of the binding energy of the neutron star • Photodisintegration can also cause collapse to a weakest force or black hole transported by neutrinos or rotation neutron star or black hole and make a supernova and magnetic fields There are critical masses for all these occurrences Similarly the light curve and spectrum depend on What kind of supernova you see depends on the the properties of the star that blew up, especially whether properties of the star (and its surroundings) in which it had a hydrogen envelope or not. Obviously white dwarfs these instabilities operate and Wolf-Rayet stars will not make Type II supernovae. • White dwarf – explosion shatters the star, but by the time the debris expand enough to let the light out the initial explosion energy has been degraded to essentially nothing. Entirely a radioactive display • Giant star – enough energy is retained (1%) that when the supernova expands and releases it (100 AU), the supernova stays bright for months • Wolf-Rayet star – like a white dwarf, the display is chiefly radioactive with perhaps some early activity from the explosion, but the explosion mechanism is collapse. • Magnetar, circumstellar interaction, and pair instability for special cases Stars inherently turn light elements into heavier ones Pair instability SNe and white dwarf explosions Massive stars and the supernovae (and the neutron stars leave behind nothing. The rest leave an interesting that they make) are responsible for the synthesis of most distribution of compact remnants of the elements heavier than helium. (Some are made by lighter stars and one or two by cosmic ray spallation) Including envelope Helium core Sukhbold and Woosley (2018) 78 WOOSLEY ET AL. Vol. 151 1 Primordial α−process# 1957 4 H He e-process Part of the KEPLER network 3 4 Hydrogen burning 5 6 7 8 9 10 Li Be B C N O F Ne x-process s-process Each nucleus is subject to 11 12 Helium burning r-process 13 14 15 16 17 36 up to 15 strong and weak Na Mg Al Si P S Cl Ar reactions (plus neutrinos) 19 20 21 22 23 24 25 26 27 28 29 30 31 32 33 34 35 36 K Ca Sc Ti V Cr Mn Fe Co Ni Cu Zn Ga Ge As Se Br Kr 37 38 39 40 41 42 43 44 45 46 47 48 49 50 51 52 53 54 Rb Sr Y Zr Nb Mo Tc Ru Rh Pd Ag Cd In Sn Sb Te I Xe 1 Now 4 H Big Bang Oxygen burning He 3 4 Neutrino irradia=on Silicon burning 5 6 7 8 9 10 during supernova Li Be and the e-process B C N O F Ne Carbon and 11 12 Neon burning 13 14 15 16 17 36 Na Mg Al Si P S Cl Ar 19 20 21 22 23 24 25 26 27 28 29 30 31 32 33 34 35 36 Fig.
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