Stellar Evolution. Eta Carinae
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1 Stellar Evolution Eta Carinae
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3 Evolution of Main Sequence Stars solar mass star: from: Markus Bottcher lecture notes, Ohio University
4 Evolution off the Main Sequence: Expansion into a Red Giant Inner core is inert isothermal helium gas. When the mass of this gas exceeds the Schonberg- Chandrasekhar limit, it begins to collapse. M ic M 0.5 The contraction brings in more hydrogen from the inner envelope, which becomes hot enough to fuse Hydrogen-burning shell around an inert helium core is much more luminous than MS core. 2 µenv µ ic
5 Red Giant Star Fusion rate in H- burning shell is HIGHER than fusion rate during main sequence (due to higher temperature) Higher luminosity causes envelope to adjust larger radius colder photosphere 5 M Sun
6 Subgiant Branch Radius increases by roughly a factor of 3 Temperature decreases significantly Red-giant Branch Temperature of envelope has fallen to point where opacity is high enough that convection carries majority of energy to surface. Higher fusion rate corresponds to higher luminosity, but nearly constant photosphere temperature Radius increases to roughly 100R Sun Helium core continues to contract and heat up
7 Low-Mass Stars (M<8M Sun )
8 Helium Flash (for M<2M Sun ) Contracting helium core becomes very dense and hot (T=10 9 K) Eventually density of electrons becomes high enough where quantum-mechanical effects become important (fermi-dirac statistics) Degeneracy pressure begins to dominate thermal pressure Due to the Pauli exclusion principle: high densities require some electrons to have very high momenta (to avoid duplicating wave functions) Degeneracy pressure will be discussed in great detail later.
9 Helium Flash Helium begins fusing when T core >10 8 K Since the core is degenerate, the onset of fusion doesn t cause the pressure to increase, so there is no thermal expansion. Thus the temperature and fusion rate increase quickly (thermonuclear runaway) L core approaches L sun for a few seconds (this energy goes into expansion, not into radiation of photosphere) Eventually thermal pressure kicks in and causes expansion of the core Expansion causes decrease in core temperature and decrease in luminosity of hydrogen burning shell
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11 Helium Burning Star (Horizontal Branch) Core reaches thermal equilibrium with a mass density of 2x10 7 kg/m 3 and T = 10 8 K Since the core expanded, fusion rate decreases, photosphere decreases, luminosity drops (by an order of magnitude), and temperature of photosphere increases. Some carbon fuses with helium to form oxygen
12 Core mass ends up being somewhat independent of star s initial mass. Typical value is 0.45M Sun Luminosity of helium-burning stars roughly the same, but stars with more massive envelopes are bigger and cooler. Helium-burning stars last for roughly 120 Myr (somewhat independent of overall mass) from: Onno Pols lecture notes (from Harvard University)
13 Asymptotic Giant Branch As the helium in the core becomes depleted, the core is mostly carbon and oxygen. As fusion turns off in the core, core contracts and heats up. Double shell burning star (Red Supergiant): Helium burning shell forms; it has a very high fusion rate and dominates the overall luminosity. Higher luminosity causes photosphere to expand and cool. Core continues to contract, luminosity continues to increase
14 Post AGB and Planetary Nebulae Carbon-Oxygen core continues to contract Core becomes partially degenerate Envelope continues to expand, and material is weakly bound to star. Steep pressure gradients and stellar pulsations cause a pretty strong mass-loss rate ( ) M sun /yr Helium-shell flashes occur in outer core (triple alpha) Envelope temperature becomes low enough for periodic recombination to occur (which increases opacity) Dust particles also increase opacity The outer layers of a red giant become unstable and enter into a series of growing pulsations.
15 Planetary Nebula For stars with a mass M<8M Sun, mass loss causes essentially the entire envelope to be ejected prior to the core becoming massive/hot enough for carbon to fuse. As the envelope mass decreases, effective temperature of star increases, and eventually the spectrum contains enough UV to photoionize the ejected envelope. This gas then emits emission lines via recombination. Fusion turns off as envelope is is ejected; luminosity drops. Timescale=10 4 y
16 Planetary Nebulae Remnants of stars with initial mass < 8M sun Radii: R ~ light years Expanding at ~10 20 km/s ( Doppler shifts) Less than 10,000 years old Have nothing to do with planets! arkus Bottcher lecture notes, Ohio University The Helix Nebula
17 Planetary Nebulae Often asymmetric, possibly due to Stellar rotation Magnetic fields Dust disks around the stars ffrom: Markus Bottcher lecture notes, Ohio University The Butterfly Nebula
18 Ring Nebula
19 Eskimo Nebula
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23 Evolutionary Stages Stage Duration Temperature Luminosity Radius Main Sequence 10 billion Red Giant 1 billion , Helium Star 100 million Asymptotic Giant 10 million ,
24 White Dwarfs Inert carbon/oxygen core Luminosity due to thermal radiation Density is roughly 10 9 kg/m 3 Typical mass is 0.6 M sun Those with M<0.45 M sun are usually He white dwarfs (which lose their envelope during the RGB phase) Chandrasekhar mass limit: 1.46 M Sun Roughly the same size as Earth
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27 Evidence for Stellar Evolution: HR Diagram of the Star Cluster M 55 High-mass stars evolved onto the giant branch Turn-off point Low-mass stars still on the main sequence
28 Estimating the Age of a Cluster The lower on the MS the turn-off point, the older the cluster.
29 High-Mass Stars (M>8M Sun )
30 Evolution of High-Mass Stars (Ch. 15) Initial PMS stages of high-mass stars are similar to those of low-mass stars: Hydrogen core fusion (main sequence) Hydrogen shell burning (supergiant) Helium core fusion (supergiant) Carbon in core can fuse together and also fuse with leftover helium (this occurs if core mass is > 1.06M Sun ).
31 Meynet and Maeder, Astron. Astrophys., 404, 975, 2003
32 Carbon Burning Carbon burning lasts several hundred years Oxygen and Neon are the most prominent elements Timescale for core evolution much faster than dynamical timescale for envelope.
33 Oxygen and Neon Burning Temperature around 10 9 K. Silicon-28 most prominent element due to photodisintegration Timescale is on the order of a year (for neon) to half a year (for oxygen), depending on mass!
34 Silicon Burning T = 2.9 x 10 9 K: Photodissociation produces lots of alpha particles, which then fuse. Timescale for silicon burning one day! Half life of Nickel-56 is 6 days: beta+ decay to iron-56, but core begins to collapse within minutes.
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36 Iron-56 is the most stable element (highest binding energy per nucleon); Nickel-56 can not fuse with anything for an exothermic reaction. The inner core is doomed!
37 Collapse of Inner Core Without fusion, as always, inner core contracts, and temperature increases. Density is > kg/m 3. Electron degeneracy becomes important, but electrons are relativistic so they can t stop core collapse. Electron Capture: Electrons with high enough energy to overcome the nuclear binding energy difference can be absorbed (inverse beta decay); overcall composition of inner core becomes more neutron rich. p+e n+ν e Degeneracy pressure decreases; core collapses more quickly.
38 Photo-disintegration: When contracting core approaches K, thermal photons are energetic to break apart nuclear bonds. A very large fraction of the iron is broken down to helium-4 and neutrons (based on thermal equilibrium) 56 Fe + γ 13 4 He + 4n It takes about 2 MeV per nucleon to break apart Fe-56. Thus photo-disintegration absorbs a lot of internal energy, causing pressure to decrease drastically. Core is now in an almost freefall collapse. Electron capture continues (on heavy nuclei) to make more neutron-rich nuclei. Nuclei eventually merge, forming essentially a gigantic solar-mass nucleus.composition of core is predominately neutrons (in heavy nuclei). Collapse of core takes roughly a second! R ic 20 km Eventually neutron degeneracy pressure kicks in when density reaches nuclear densities (10 17 kg/m 3 ).
39 Boom The change in gravitational potential energy due to the core collapse is roughly (3/5)GM ic2 /R ic = 3x10 46 J The gravitational binding energy of the envelope is roughly J (see hw). Thus only a small fraction of the energy released by core collapse is required to blow away the envelope. (see hw) SN 1987A L L
40 Supernova: Observational Properties Supernova luminosity reaches a peak of L sun (depending on mass), staying bright for several months. At least eight supernovae have been observed in our galaxy with the naked eye in the last 2000 years There are basically two types of supernovae: Type Ia and Type II. Type II have strong hydrogen lines in the spectra, and are due to the collapsing cores of massive stars. Type Ia do not have hydrogen lines, and are due to fusion of carbon/oxygen in white dwarfs.
41 In summary:
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