1	Astro 1050 Mon. Mar. 3, 2004 Today: Start Ch. 10: The Deaths of Stars Questionnaire changes
2	Astro 1050 Questionnaire results Math: About half think too much/too hard Solution: End of Lab problem sessions Best/Worst: Lab, especially Observing good Extra credit good, too Not so good: low lighting, boring lectures (sometimes not clear), test Solutions: Extra Red Buttes Observing (optional), Keep on lights, Flash cards for lecture feedback (also, please, ask more questions!), Extra Credit for sending in an exam question Review Session: YES! Solution: Thursday, Mar. 11, 6PM in labroom
3	Astro 1050 Extra Credit Chance Email me (mbrother@uwyo.edu) a single multiple choice exam question covering Chapters 6-10. Do this no later than class time Monday. This is worth 5 homework questions. I will try to use as many questions that I think are good on the next exam. I like a mix of easy questions, hard questions, mathematical questions, conceptual questions, and descriptive questions. If you write a good question and I use it, hey, that’s an easy point for you! If this works out well, we’ll do it on other exams.
4	Chapter 10 – The Deaths of Stars What happens when the hydrogen runs out? Star once again tries to contract and heat up inside What might stop this? “Unusual” fusion energy sources Þ giant stars Hydrogen shell fusion Heavy element fusion Degenerate electron pressure Þ white dwarfs Loss of material from the star
5	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 “shell burning” Shell burning changes the rules for structure Still need dense core to allow high T, P for fusion But high T on outside of core puts too much energy into outer parts of star Outer parts responds to input energy by expanding and cooling Star in effect separates into two parts Hot dense core where energy generation takes place Extended envelope which shields and insulates core
6	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 L µ R² T⁴ and L is not much greater (so far) R² is »(75/3)²=25² = 625 times larger T must decrease to compensate
7	HR Motion for other Mass Stars For higher mass (already luminous) stars Motion is more horizontal (to red) Hard to increase luminosity above already very high levels
8	Motion in the HR Diagram 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_Coreincreases) Eventually He ignites in the core
9	Motion in the HR Diagram 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
10	Expected Evolution of the HR Diagram for a Cluster A cluster consists of stars 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 move off the main sequence (at the moving “turn off point”) The original supergiants don’t live very long. The lower mass stars produce giants
11	Tests of Stellar Evolution using the HR Diagram Stellar evolution too slow to see changes in given cluster Can observe clusters and look for predicted patterns
12	Complications in Stellar Evolution Pressure forces other than thermal gas pressure 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
13	Pauli Exclusion Principle From quantum rules, electrons don’t like to be packed into a small space, either in atoms or in ionized gas At normal ionized gas densities, electrons are so spread out quantum rules don’t matter. As high enough ionized 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)
14	Effect of Degenerate Electron Pressure 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.
15	Degenerate Pressure Can End Fusion Degenerate Electron Pressure limits contraction and 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
16	Effects of Convection Energy can be moved by radiation or convection Convection in core brings in new fuel Cooler material more opaque making radiation harder and convection more likely Choice also depends on energy flux
17	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 to fast Energy can be generated in “thermal pulses” Low temperature opaque envelope can also “oscillate” Energy is transmitted in “pulses” as envelope expands and contracts Main cause of “Variable Stars”
18	White Dwarfs
19	Simple Planetary Nebula IC 3568 from the Hubble Space Telescope
20	Complicated P-N in a Binary System M2-9 (from the Hubble Space Telescope)
21	A Gallery of P-N from Hubble
22	Complications in Binary Systems Can move mass between stars 1^st (massive) star becomes red giant Its envelope 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 let material 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 away remaining fuel
23	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 r means not only more electrons, but faster electrons, which raises pressure a lot. Relativistic degenerate gas: increasing r can’t increase electron velocity (they are already going close to speed of light) so pressure doesn’t go up as much
24	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 (r), not temperature, so must have higher density P vs. r rule such that higher mass star must actually have smaller radius to provide enough P As M_star ® 1.4 M_Sun v_electron ® c Requires much higher r to provide high enough P, so star must be much smaller. Strong gravity which goes with higher r makes this a losing game. For M ³ 1.4 M_Sun no increase in r can provide enough increase in P – star collapses
25	Implications for Stars Stars less massive than 1.4 M_Sun can end as white dwarfs Stars 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
26	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 Type I: White dwarf in binary, which receives mass from its companion (collapse ignites carbon burning). 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 + n Neutrinos (n) escape, carrying away much of energy, aiding collapse Collapsing outer part is heated, “bounces” off core, is ejected into space Light 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)
27	Supernova in Another Galaxy Supernova 1994D in NGC 4526
28	Tycho’s Supernova of 1572 Now seen by the Chandra X-ray Observatory as an expanding cloud.
29	The Crab Nebula – Supernova from 1050 AD Can see expansion between 1973 and 2001 Kitt Peak National Observatory Images
30	What happens to the collapsing core? Neutron star (more in next chapter) Quantum rules also resist neutron packing Densities much higher than white dwarfs allowed R ~ 5 km r ~ 10¹⁴ gm/cm³ (similar to nucleus) M limit uncertain, ~2 or ~3 M_Sun before it collapses Spins very fast (by conservation of angular momentum) Trapped spinning magnetic field makes it: Act like a “lighthouse” beaming out E-M radiation (radio, light) pulsars Accelerates nearby charged particles
31	Spinning pulsar powers the Crab nebula Red: Ha Blue: “Synchrotron” emission from high speed electrons trapped in magnetic field