Astro 1050 Fri. Apr. 10, 2015
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1 Astro 1050 Fri. Apr. 10, 2015 Today: Continue Ch. 13: Star Stuff Reading in Bennett: For Monday: Finish Chapter 13 Star Stuff Reminders: Ch. 12 HW now on Mastering Astronomy, due Monday. Ch. 13 will be assigned on Monday due Apr. 20 Rooftop Solar Observing: Tues., Wed. Thurs. 11:30pm 1:13 p.m. Note that sessions will be canceled for clouds 1
2 Chapter 13 Cont. The Deaths of Stars What happens when the hydrogen runs out? With no internal heat source the center cools Star once again will contract and interior temperature increases. What is the result? Unusual fusion energy sources giant stars Hydrogen shell fusion Heavy element fusion Degenerate electron pressure white dwarfs Loss of outer gas envelope from the star 2
3 Hydrogen Shell Burning Fusion stops in core when hydrogen runs out Star has core of He, but T too low for fusion there Heat loss makes star contract, T goes up in interior Before T in core reaches He ignition point -- Hydrogen above He core begins rapid shell burning From our text: Horizons, by Seeds Shell burning changes the rules for structure Still need dense core to provide high T, P for fusion But high T on outside of core puts too much energy into outer parts of star (more mass in dense core) Outer parts responds to hotter intermediate layer by expanding and cooling Shell burning layer separates star into two parts Hot dense core Extended envelope which shields and insulates H shell-burning layer 3
4 Expansion into a Red Giant For 5 M Sun red giant Outer part swells to 75 R Sun Inner core still contains most of mass Why is it red? Two equivalent ways to answer: Thick envelope insulates outside from hot core From our text: Horizons, by Seeds L R 2 T 4 and L is not much greater (so far) R 2 is (75/3) 2 =25 2 = 625 times larger Outer T must decrease to compensate 4
5 HR Evolution for other Mass Stars For higher mass (already luminous) stars Evolution is more horizontal (to red) Hard to increase luminosity above already very high levels From our text: Horizons, by Seeds 5
6 Evolution in the HR Diagram From our text: Horizons, by Seeds During H core burning (main sequence) L increases slowly as He accumulates in the core (and T Core increases) During H shell burning At first R increases, T decreases L then increases slowly as more He accumulates in core (and T Core increases) Eventually He ignites in the core 6
7 Evolution in the HR Diagram From our text: Horizons, by Seeds During He core, H shell burning Moving some of E generation back into core shrinks size (and so raises T Surface ) Eventually He in core runs out leaving a C, O core inside the He region Core contracts and heats up till He ignites in shell outside C, O core During He shell, H shell burning At first R increases, T Surface decreases L increases as more C,O accumulate and T Core continues to increase 7
8 Expected Evolution of the HR Diagram for a Cluster A cluster consists of stars of different mass which all started to form at the same time (collapsing from the same fragmenting molecular cloud) Higher mass stars contract to main sequence before low mass ones reach it Higher mass stars are also the first to run out of H and leave the main sequence, becoming supergiants. As time continues, lower and lower mass stars leave the main sequence (at the evolving turn off point ) From our text: Horizons, by Seeds The original supergiants don t live very long. The lower mass stars produce giants 8
9 Tests of Stellar Evolution using the HR Diagram From our text: Horizons, by Seeds Stellar evolution too slow to see any changes in given cluster Can observe clusters and look for predicted patterns 9
10 Complications in Stellar Evolution Pressure forces other than thermal gas pressure (radiation contributes) Reminder: We ve been assuming that when star loses energy it contracts and actually heats up. Clearly not all objects do this (eg. Earth) Convection bringing in fuel from outer regions Mass loss from stellar wind, or mass gain from nearby star 10
11 Pauli Exclusion Principle From quantum rules, electrons don t like to be packed into a small space, either in atoms or in ionized gas (causes shell structure) At normal ionized gas densities, electrons are so spread out quantum rules don t matter. At very high gas densities, quantum rules need to be considered, just has they have been in atoms Think of each atom sized region of space having a set of energy levels associated with it (although it is really more complicated) An energy level can only hold one electron (2 if we consider opposite electron spins ) In multi-electron atoms, you can have at most two electrons in each energy level. In gas, to pack more electrons in same volume, new ones must be placed in higher energy levels (i.e. be going faster) than existing ones 11 From our text: Horizons, by Seeds
12 Effect of Degenerate Electron Pressure Electrons begin to strongly interact with each other Loss of energy does not reduce pressure Star does not contract in response to loss of energy Gravity not available as energy source to heat up star Electrons are already in lowest energy states allowed (equivalent to atoms in ground state) so no energy available there If there is no other energy source, as energy is lost nuclei move slower and temperature drops. 12 From our text: Horizons, by Seeds
13 Degenerate Pressure Can End Fusion Degenerate Electron Pressure can prevent contraction and stabilize core temperature Stars with M < 0.08 M Sun never burn H (brown dwarfs) Stars with M < 0.4 M Sun never burn He (red dwarfs) Stars with M < 4 M Sun never burn C (but do make red giants) Stars with M > 4 M Sun do burn elements all the way to Fe What happens to these objects? Brown dwarfs never become bright sort of giant version of Jupiter Red dwarfs have such long lives none have yet exhausted H Red giants are related to white dwarfs Massive stars explode as supernova 13
14 Effects of Convection From our text: Horizons, by Seeds Energy can be transported by radiation or convection Convection in core brings in new fuel Cooler material more opaque making radiation harder and convection more likely Which dominates also depends on amount of energy from core 14
15 Mass Loss from Giant Stars Envelope of red giant very loosely held Star is so big, gravity very weak at the surface Degenerate core makes nuclear thermostat sluggish Core doesn t quickly expand and cool when fusion is too fast and vice versa Energy can be generated in thermal pulses Low temperature and more opaque envelope can also oscillate Energy is transmitted in pulses as envelope expands and contracts Main cause of many Variable Stars Some Red Giants have very strong stellar winds ejecting envelopes, forming Planetary Nebulae PN called this because some looked like planets in early telescopes PN don t have anything to do with formation of planets 15
16 White Dwarfs The now exposed degenerate core is a White Dwarf For 1 M Sun: R~R Earth ρ ~ gm/cm 3 Faint because it is so small, despite high T Can t contract because P doesn t drop with T (degenerate gas) Slowly cools and fades 16
17 Simple Planetary Nebula IC 3568 from the Hubble Space Telescope 17
18 Complicated P-N in a Binary System M2-9 (from the Hubble Space Telescope) 18
19 A Gallery of P-N from Hubble 19
20 Complications in Binary Systems Mass can be transferred between stars 1 st (massive) star becomes red giant Its envelope expands and is transferred to other star Hot (white dwarf) core exposed 2 nd star becomes red giant Its envelope transferred to white dwarf Accretion disk around white dwarf Angular momentum doesn t allow material to fall directly to white dwarf surface Recurrent nova explosions White dwarf hot enough for fusion, but no Hydrogen fuel New fuel comes in from companion Occasionally ignites explosively, blowing off partially burnt fuel 20
21 Is a star stable against catastrophic collapse? Imagine compressing a star slightly (without removing energy) Pressure goes up (trying to make star expand) Gravity also goes up (trying to make star collapse) Does pressure go up faster than gravity? If Yes: star is stable it bounces back to original size If No: star is unstable gravity makes it collapses Ordinary gas: P does go up fast stable Non-relativistic degenerate gas: P does go up fast stable Relativistic degenerate gas: P does not go up fast unstable Relativistic: Mean are the electrons moving at close to the speed of light Non-relativistic degenerate gas: increasing ρ means not only more electrons, but faster electrons, which raises pressure a lot. Relativistic degenerate gas: increasing ρ can t increase electron velocity (they are already going close to speed of light) so pressure doesn t go up as much 21
22 Chandrasekhar Limit for White Dwarfs Add mass to an existing white dwarf Pressure (P) must increase to balance stronger gravity For degenerate matter, P depends only on density (ρ), not temperature, so must have higher density P vs. ρ rule such that higher mass star must actually have smaller radius to provide enough P From our text: Horizons, by Seeds As M star 1.4 M Sun v electron c Requires much higher ρ to provide high enough P, so star must be much smaller. Strong gravity which goes with higher ρ makes this a losing game. For M 1.4 M Sun no increase in ρ can provide enough increase in P star collapses 22
23 Implications for Stars Cores less massive than 1.4 M Sun can end as white dwarfs Cores more massive than 1.4 M Sun can end as white dwarfs, if they lose enough of their mass (during PN stage) that they end up with less than 1.4 M Sun Stars whose degenerate cores grow more massive than 1.4 M Sun will undergo a catastrophic core collapse: Neutron stars Supernova 23
24 Supernova When the degenerate core of a star exceeds 1.4 M Sun it collapses Type II: Massive star runs out of fuel after converting core to Fe (Iron core cannot fuse, it photo-dissociates first actually absorbing energy!) Type I: White dwarf in binary, which receives mass from its companion (collapses when M > 1.4 M Sun ignites carbon burning). Chronology of events: Star s core begins to collapse Huge amounts of gravitational energy liberated Extreme densities allows weak force to convert matter to neutrons p + + e - n + ν Neutrinos (ν) escape, carrying away much of energy, aiding collapse Collapsing outer part is heated, bounces off core, is ejected into space Light emitted from very hot ejected matter makes supernova very bright Ejected matter contains heavy elements from fusion and neutron capture Core collapses into either: Neutron stars or Black Holes (Chapter 11) 24
25 Supernova in Another Galaxy Supernova 1994D in NGC
26 Tycho s Supernova of 1572 Now seen by the Chandra X-ray Observatory as an expanding cloud. 26
27 The Crab Nebula Supernova from 1050 AD Can see expansion between 1973 and 2001 Kitt Peak National Observatory Images 27
28 What happens to the collapsing core? Neutron star (more in next chapter) Quantum rules also resist neutron over-packing Densities much higher than white dwarfs allowed R ~ 5 km ρ ~ gm/cm 3 (similar to nucleus) M limit uncertain, ~2 or ~3 M Sun before it collapses Spins very fast (conservation of angular momentum) Trapped spinning magnetic field makes it: Act like a lighthouse beaming out E-M radiation (radio, light) Accelerates nearby charged particles Pulsar! 28
29 Spinning pulsar powers the Crab nebula Red: Hα Blue: Synchrotron emission from high speed electrons trapped in magnetic field 29
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