1. Midterm Exam #2 Tuesday, March 23 Closed book Will cover Lecture 8 (Special Relativity) through Lecture 14 (Star Formation) only If a topic is in the book, but was not covered in class, it will not be on the exam! Some combination of multiple choice, short answer, short calculation Equations, constants will all be given Standard calculators allowed Cell phones, PDAs, computers not allowed

2. Outline - March 18, 2010 Protostar to “Main Sequence” star (pgs. 556-57) HR Diagram revisit Define “low”, “intermediate” and “high” mass stars (pg. 565) Evolution and death of low mass stars (pgs. 566-572) Evolution and death of high mass stars (pgs. 572-581)

3. Protostar to Main Sequence Star Protostar becomes a main sequence star with the onset of hydrogen fusion

4. H-R Diagram Review About 90% of stars in the sky are “Main Sequence” stars All main sequence stars are stable (gravity exactly balances pressure) and energy source is fusion of HYDROGEN to form HELIUM What are all of the “non- main squence” objects on the HR diagram?

5. It’s all about stellar evolution… What determines where a star is on the HR diagram? Its evolutionary state! What determines how a star evolves? Its “main sequence” mass! Luminosity on the MS and lifetime on the MS depend on the star’s “main sequence” mass: L = (M / M sun ) 3 L sun Main sequence lifetime = 10 / (M / M sun ) 2 billion years

6. Main Sequence Lifetime (about 90% of a star’s total lifetime) The less massive is a star, the less fuel it has, but the longer it will last on the main sequence. The more massive is a star, the more fuel it has, but the shorter it will last on the main sequence. “Low-mass” stars: born with M < 2 M sun “Intermediate-mass” stars: born with 2 M sun < M < 8 M sun “High-mass” stars: born with M > 8 M sun Main Sequence lifetime of 0.2 M sun star = about 500 billion years Main Sequence lifetime of a 10 M sun star = about 100 million years

7. Evolution of “Low-Mass” vs. “High-Mass” Stars Recall: stability of star is all about pressure-gravity balance Main Sequence star: pressure comes from conversion of H to He Problem: He nucleus has 2 protons (H nucleus has 1 proton), so it takes a higher density and higher temperature to fuse He than it does to fuse H Cores of MS stars: not hot enough or dense enough to fuse He In a main sequence star, He is a nuclear “ash” - it doesn’t contribute to “holding up” the star against gravitational collapse.

8. Build-up of Inert Helium Core Eventually, the star builds up a substantial He core, with H “burning” in a shell around the core. The H burns into layers of the star that are thinner, and thinner, making it harder to hold the star up against gravitational collapse. The He core can provide a little bit of help by contracting (conversion of gravitational energy, just like a protostar). As the core contracts, the outer envelope expands and the star leaves the main sequence.

9. Evolution of a Low-Mass Star As He core contracts, the star moves up the HR diagram. As outer envelope expands, the star becomes physically larger (increases luminosity) and the surface temperature cools (becomes redder). Star becomes a Red Giant. Onset of He burning in the core happens quite suddenly (helium “flash”) once the temperature and density of the core are high enough to fuse He. Helium flash doesn’t disrupt the star (localized region of 1/1000 of the star), but does cause the core to expand a little bit (and envelope shrinks in response).

10. Red Giant Phase for Low Mass Stars Core is now 100 million Kelvin (about 10x hotter than when the star was a main sequence star) Two sources of energy: 1. H to He in a shell 2. He to C (“triple alpha” process) in the core

11. “Triple Alpha” Process (nuclear fusion of Helium to produce Carbon)

12. Red Giants are Truly Enormous (sun as a red giant results in end of life on Earth) Today 5 billion years in the future When the sun becomes a Red Giant it will engulf Mercury, and perhaps Venus. The surface temperature of the sun will be about 1/2 its current temperature, but the sun will be so large that it will take up half the noon-time sky! End of life on earth - we’ll be toasted.

13. Final Stage of Evolution of Low-Mass Star It’s only a matter of time before the star gets in trouble again… This time it’s CARBON ash that has sunk to the center (non-burning carbon core, surrounded by a shell of He burning, surrounded by a shell of H burning). Most low mass stars can repeat the core contraction process, and ignite Carbon fusion (which produces Oxygen). But, once a significant amount of oxygen has built up in the core, it’s game over for the star!!

14. Death of a Low Mass Star Carbon-Oxygen core contracts in an attempt to help hold the star up against gravitational collapse; but there isn’t enough mass in the star to make the temperature and density high enough to fuse the oxgygen Core shrinks down to about the size of the earth, and can’t go any farther because of a quantum mechanical effect Can only compress electrons so far - this is what stops the core contraction Pressure in the core is provided by “degenerate electron gas” and the core becomes stable (no longer contracting) “Burning fronts” (H, He, C) plow out into the very light, fluffy layers of the (enormous!) star, and the outer layers of the star “lift off” due to radiation pressure

15. Formation of White Dwarf and Planetary Nebula (end of a low-mass star) Outer layers of star lift off, revealing small, hot core = White Dwarf Initially, the white dwarf is very hot, but it cools off because it has no internal source of energy (will eventually be black!) Sirius A Sirius B (white dwarf companion to Sirius A)

16. Planetary Nebulae (have nothing to do with planets!) Gas from original envelope of star is heated by the white dwarf

17. Evolutionary Track on the HR Diagram (Low-Mass Star)

18. Evolution of High-Mass Stars Unlike low mass stars, high mass stars make a steady transition from H fusion in the core to He fusion in the core (no “helium flash”), to O fusion in the core, and they keep on going to heavier chemical elements. High-mass stars evolve off the main sequence to become “supergiant” stars.

19. “Onion Layers” of Fusion in a High-Mass Star Star undergoes cycles of core contraction and envelope expansion, fusing heavier and heavier chemical elements, until an iron core forms. Once silicon starts to fuse, the star has about a week to live.

20. Timescales of Fusion (M star = 20 M sun ) H fusion in core: 10 million years He fusion in core: 1 million years C fusion in core: 1000 years O Fusion in core: 1 year Si fusion in core: 1 week

21. What’s so special about Iron (Fe)? Fusion of nucleii that are lighter than iron result in a net gain of energy (takes less energy to bring the nucleii close together than you get from mass loss) Fusion of nucleii that are as heavy or heavier than iron result in a net loss of energy (takes more energy to bring the nucleii close together than you get from mass loss) Bottom line: star can’t use iron as a nuclear fuel to support itself from gravitational collapse, because fusing iron is a losing proposition in the energy balance!

22. Death of a High-Mass Star Supernova: Implosion followed by Explosion Once substantial amount of iron has built up, star implodes on itself Core reaches temperature of 10 billion Kelvin (= tremendously high energy photons), the nuclei are split apart into protons and neutrons (“photodisintegration”) In less than 1 second, the star undoes most of the effects of nuclear fusion that happened in the previous 11 million years!!!!! High-energy photons are absorbed, giving rise to loss of thermal energy in the core, the core becomes even more unstable, and the collapse accelerates Protons and electrons in the core combine together (“neutronization”), resulting in nothing but neutrons in the core Collapse continues until it’s not possible to squeeze the neutrons together any tighter (size of core = size of Manhattan) Collapse starts to slow, but overshoots and outer layers of star are driven out into space (perhaps by “bounce” off the neutron core) in a massive explosion

23. Supernovae Generate Tremendous Amounts of Energy At their maximum brightness, supernovae are as bright as an entire galaxy. Peak luminosity is about 10 51 ergs = the sun’s total output of energy over 10 billion years!

24. How long does a supernova last? “Type II” supernovae are exploding high- mass stars “Type Ia” supernovae are something else entirely (and involve binary star systems)

25. Why should you care about supernovae? Extraordinarily bright, so can use them to measure distances to galaxies that are very far away: b = L / (4  d 2 ) Supernovae are the source of all heavy chemical elements! The heavy chemical elements are produced during the explosion itself, when there is more than enough energy to fuse nuclei heavier than iron (doesn’t matter that there is a “net loss” of energy - the star is already VERY far out of equilibrium)

26. Supernova Remnants (high-mass star guts)

27. Cycle of Star Formation and Supernovae Stars form out of gas in the ISM, evolve, and blow much of themselves back into the ISM Massive stars create heavy chemical elements during the explosions, which “enriches” the ISM with heavy chemical elements New stars form, and make yet more heavy chemical elements It takes about 500 cycles of massive star formation to account for all the heavy chemical elements in the universe More than enough time for this to happen (universe is 14 billion years old, massive stars take a few million years to evolve and explode)