Outline - March 25, 2010 H-R Diagram Recap: Evolution and death of low mass stars (pgs. 566-572) About 90% of stars in the sky are Main Sequence stars Evolution and death of high mass stars (pgs. 572-581) Stellar Remnants (white dwarf, neutron star, pulsar, black hole) Novae and supernovae All main sequence stars are stable (gravity exactly balances pressure) and energy source is fusion of HYDROGEN to form HELIUM All of the non-main sequence objects are no longer burning H in their cores (are evolved stars) 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). 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. As the core contracts, the outer envelope expands and the star leaves the main sequence. 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). 1
Red Giant Phase for Low Mass Stars Final Stage of Evolution of Low-Mass Star It s only a matter of time before the star gets in trouble again 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 in the core 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!! 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 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 and creating a Planetary Nebula 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 Sirius B (white dwarf companion to Sirius A) 2
Evolutionary Track on the HR Diagram (Low-Mass Star) 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. Onion Layers of Fusion in a High-Mass Star Timescales of Fusion (M star = 20 M sun ) Star undergoes cycles of core contraction and envelope expansion, fusing heavier and heavier chemical elements, until an iron core forms. H fusion in core: 10 million years He fusion in core: 1 million years C fusion in core: 1000 years Once silicon starts to fuse, the star has about a week to live. O fusion in core: 1 year Si fusion in core: 1 week 3
What s so special about Iron (Fe)? Death of a High-Mass Star Supernova: Implosion followed by Explosion Once substantial amount of iron has built up, star implodes on itself 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! 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 Supernovae Generate Tremendous Amounts of Energy How long does a supernova last? At their maximum brightness, supernovae are as bright as an entire galaxy. Type II supernovae are exploding highmass stars Peak luminosity is about 10 51 ergs = the sun s total output of energy over 10 billion years! Type Ia supernovae are something else entirely (and involve binary star systems) 4
Why should you care about supernovae? Extraordinarily bright, so can use them to measure distances to galaxies that are very far away: b = L / (4π d2) 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) 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) Supernova Remnants (high-mass star guts) Stellar Remnants What s left behind after a star dies? Main sequence mass < 5 M sun: white dwarf Main sequence mass between 5 M sun and 40 M sun: neutron star Main sequence mass > 40 M sun: black hole All of these are stable (neither expanding nor contracting), so long as they are left alone. Pressure in white dwarf and neutron star is somewhat exotic (not normal gas pressure or radiation pressure) due to their highly-compressed states. 5
White Dwarfs White Dwarfs in Binary Systems Pressure comes from degenerate electron pressure Electrons packed together as tightly as quantum mechanics allows; their speeds support the WD against gravitational collapse WD acts a lot like a metal (same temperature and density throughout) Maximum WD mass = 1.4 M sun ( Chandrasekhar limit ) WD with mass = 1 M sun is about the size of the earth, weight of 1 teaspoon of WD material = about the weight of a small truck If all alone in space, WD simply cool off (no internal source of energy) and eventually become black Most stars are found in binary systems May have situation where WD orbits a giant or supergiant star at a relatively close distance Outer layers of the giant or supergiant are very light and fluffy, and may be pulled over onto the WD by gravity Material from companion star builds up in an accretion disk around the WD, and eventually winds up on the surface of the WD White Dwarfs in Binary Systems, II Novae What happens to the WD when mass is dumped onto it depends on how much mass, and how fast. Slow accretion of not much mass (not enough to make the mass of the WD > 1.4 M sun ): nova Fast accretion of a lot of mass (enough to make the mass of the WD > 1.4 M sun rather suddenly): supernova ( Type 1a ) Thin layer of (mostly) H from the companion star builds up on surface of WD Sudden flare in brightness (increases by about a factor of 10,000 or more), then fades over the course of about a month Flare is due to hydrogen fusion on the surface of the white dwarf Novae happen about 2 or 3 times per year in our Galaxy Can recur (i.e., same WD can go nova, but not very predicable) 6
Novae H fusion on the surface of a WD White Dwarf Supernova Supernova Type Ia If the companion star to a WD dumps a lot of mass onto the WD very quickly, making the mass of the WD exceed the Chandrasekhar mass (1.4 M sun), the WD explodes as a supernova! WD is much like a hot metal ball, same temperature and same density throughout Addition of extra mass causes WD to contract (gravity wins over pressure from the electrons) and instantaneously the carbon starts to fuse throughout all parts of the WD, blowing the WD to bits Naked eye nova; picture taken in the Varzaneh Desert in Isfahan, Iran (February 2007) Two Basic Types of Supernovae What s left behind after a massive star goes supernova? Note: Supernovae NEVER repeat! If the mass of the core is less than about 3 M sun, a neutron star is left behind. If the mass of the core is greater than about 3 M sun, there is no source of sufficient pressure to keep the core from collapsing completely under gravity, and a black hole is formed. Remnants of two different supernovae. Left: a Type Ia supernova (WD). Right: a Type II supernova (high mass star). This is a happy alignment of images - the two stars weren t related to each other! 7
Pulsars Neutron Stars Rapidly-Rotating Neutron Stars Even more compressed than WD Typically the size of a city (about 10 km in radius) with mass between 1.4 Msun and 3 M sun Density is such that the weight of one teaspoon of NS material would weigh 100 million tons (vs. 1 ton for WD material) NS supported by neutron degenerate pressure (again, quantum mechanical phenomenon having to do with how tightly neutrons can be packed together) Now know of 100 s of pulsars, most with periods between 0.03 and 0.3 seconds (meaning they rotate between 3 and 30 times per seconds) Must rotate extremely fast (conservation of angular momentum); a star that was originally rotating once per month would now have to rotate a few times per second! The fastest known pulsars have periods of milliseconds and are rotating at speeds approaching 0.25c!!!! Compression of the material also compresses the magnetic field, and amplifies its strength (making it trillions of times larger than the earth s magnetic field) Pulsar at Center of Crab Nebula First discovered by Jocelyn Bell (1967) as pulses of radio light coming from the Crab Nebula The pulses lasted 0.01 seconds, and repeated every 1.4 seconds In 1974, Jocelyn s PhD advisor (Tony Hewish) got the Nobel Prize for explaining what puslars are The radio light from plusars is called synchrotron radiation, a type of light that is emitted by electrons as they move on spiral paths around magnetic field lines. The synchrotron radiation pulses are proof of the fast rotation rates of neutron stars and the presence of an incredibly strong magnetic field. Pulsar Recordings http://www.astrosurf.com/luxorion/audiofiles-pulsar.htm Supernova observed by Chinese astronomers in year 1054 Crab - 1.4 rotations per second Vela - 11 rotations per second PSR 1937+21 (a millisecond pulsar) 642 rotations per second Crab Nebula - remnant of supernova explosion 8
Midterm Exam #2 Curve boundaries for Midterm #2: A > 90.5% A- 87.5% to 90.5% B+ 81% to 87.5% B 75% to 81% B- 69% to 75% C+ 66% to 69% C 63% to 66% C- 57% to 63% D 50% to 57% F < 50% Class letter grade average based on the curve is between B and B- (2.9 / 4.0) Approximate Mid-semester Grades A > 92% A- 86% to 92% B+ 82% to 86% B 78% to 82% B- 70% to 78% C+ 66% to 70% C 62% to 66% C- 57% to 62% D 50% to 57% F < 50% Your score on this exam: 78.5 / 100 Your ranking in the class on this exam: 21 / 47 Approximate letter grade on this exam: B This info is on the last page of your exam. Approximate mid-semester grades based on average of midterm exams, best 3 of 4 home work assignments, and average of 2 labs. 9