ASTR377: A six week marathon through the firmament

1. ASTR377:A six week marathon through the firmament by Orsola De Marco orsola@science.mq.edu.au Office: E7A 316 Phone: 9850 4241 Week 5, May 17-20, 2009

2. Overview of the course • Where and what are the stars. How we perceive them, how we measure them. • (Almost) 8 things about stars: stellar structure equations. • The stellar furnace and stellar change. • Stars that lose themselves and stars that lose it: stellar mass loss and explosions. • Stellar death: stellar remnants; PN, WDs, NS, BH! • When it takes two to tango: binaries and binary interactions.

3. What death becomes her? • Depending on the main sequence mass, a star will end in a different way. • M<0.08Mo no H burning – BD [L or T dwarfs]. • 0.08Mo < M < 0.5Mo no He burning. Stars become He WDs. • 0.5Mo < M < 5Mono C burning: stars become CO WDs. • 5Mo < M < 7Mo yes C burning to Ne and Mg: stars become ONeMg WDs. • M > 7Mo burn all the way to Fe: these stars go through a type II SN and become NSs or BHs.

4. What happens after the AGB? • Spectra of planetary nebulae indicate that the “shell” is expanding with speeds of 20-30 km/s. The central star is often visible inside the PN. Jacoby, De Marco & Sawyer 1998

5. What happens after the AGB? • Knowing angular size, distance and speed of expansion and assuming that the speed did not change since the time of the material ejection, we can determine the age of the PN. • What is wrong with the assumption above (see in a few slides)?

6. Base Moon diameter = sin  = sin  Hypotenuse D Angular size and distance  ? 3476 km D

7. Bob O’Dell American (alive) The PN clock • The central stars of PN on the HRD: they sit in a locus and have increasing ages along the red arrow. • Conclusion: the star evolves to the blue in a very short time. PN “ignition”: log Teff ~ 4.4 AGB logTeff=3.5 Youngest PNe Oldest PNe O’Dell 1968

8. At the end of the AGB • Core mass growing because of the He shell burning. This generates an increase in L. • Two shell sources, H (out) and He (in). Every time a star has multiple shell sources, all burning outwards through the fresh fuel supply, there is instability. • The relative speed at which they burn out can create an instability where the helium shell L increases with no release of pressure, leading to even more L. • Eventually the L pushes the envelope out, the H shell extinguished and the He burning rate decreases, L gets out, then the entire star returns to equilibrium. • At the end of each thermal shell flash the envelope convection zone extends downwards and dredges up the results of He burning: C and O. This is the third and last dredge up chance a star has and makes of AGB star the C factories of the Universe. http://outreach.atnf.csiro.au/education/senior/astrophysics/stellarevolution_postmain.html

9. Departing the AGB • When the AGB H envelope mass drops below a few times 10-4 Mo the star loses its equilibrium and contracts. • The contraction and the increased transparency of the envelope result in a fast increase of the effective temperature. • All the while the H or the He shell source are still burning. • Eventually the burning stops and the star cools on the WD cooling track. Vassiliadis & Wood 1994

10. The PN phase • As the star heats a fast but this time tenuous wind sweeps AGB wind material up and creates a shells. • When the photosphere of the heating star passes the ~25,000-K mark the swept up shell is ionized and can be seen in forbidden lines.

11. Stellar mass loss: the birth of a nebula Animation from the Space Telescope Science Institute

12. What do Planetary Nebulae look like in the sky

13. What do Planetary Nebulae look like in the sky 3 arcmin

14. Planetary Nebula shapes: round…. Abell 39; WIYN image; G. Jacoby

15. Planetary Nebula shapes: “elliptical”…. The Helix nebula; Spitzer image; K. Su

16. Planetary Nebula shapes: “elliptical++”…. The Cat’s Eye nebula; HST image

17. OH321.8+4.2; Bujarrabal; HST Planetary Nebula Shapes: bipolar

18. PN shapes and shaping • Young PN and pre-PN (shining from shocks not from radiative ionization) are always non-round.

19. Planetary Nebula Halos 2.5 pc NGC6543 HST/NOT [OIII]/[NII]/Ha. (P. Harrington, R. Corradi)

20. Planetary nebulae as we teach themHow do PN form? The super-wind: sudden intensification of mass-loss by a factor of 3-100 at the end of the AGB 3.4 CSPN post-AGB AGB WD Empirically shown to happen, theoretically unexplained

21. Planetary nebulae as we teach themHow do PN form? Envelope mass < 10-3 to -4 Mo: departure from the AGB. Fast thin post-AGB wind compresses the super-wind. 3.4 CSPN post-AGB AGB WD Kwok 1982; Balick 1987

22. Planetary nebulae as we teach themHow do PN form? Post-AGB star heats up T>25,000K: ionized PNA39: a well behaved PN CSPN post-AGB AGB WD

23. … but how do PN acquire their shapes?

24. … but how do PN acquire their shapes? Relatively-fast rotation during AGB super-wind … Garcia-Segura et al. 2003

25. … but how do PN acquire their shapes? … and/or magnetic fields Garcia-Segura et al. 2003

26. … but how do PN acquire their shapes? … result in circumstellar material with an equatorial enhancement

27. … but how do PN acquire their shapes? When the star heats up, on its way to becoming a white dwarf, a fast wind rums into the previously-ejected gas.

28. … but how do PN acquire their shapes? Let’s zoom out …

29. … but how do PN acquire their shapes? Let’s zoom out … … the lobes perpendicular to the plane of the disk continue to expand.

30. … but how do PN acquire their shapes? When the star heats up the gas “shines” and we see the bipolar PN

31. An orbiting companion can spin up the AGB star atmosphere … but how do PN acquire their shapes? The problem: giant stars do not rotate fast enough.

32. How can a companion spin up a giant? Animation from the Space Telescope Science Institute

33. White dwarfs: key properties • Small – about Earth radius. • Not (all) white: some are very cool (~4000K). • Super dense. Srius A and B; HST/FGS image

34. Friedrich Wilhelm Bessel German 1784-1846 History of WDs • 1844 Friedrich Bessel discovers that Sirius changes position and the change is not due to annual parallax…. There must be an unseen companion. • In 1864 the companion was found by Alvan Graham Clark. It was 25,000K bit only 10,000th the luminosity of the Sun. What does this mean? • A better solution of the orbit gave us Sirius B’s mass (0.9Mo) This meant that this star was much denser. Alvan Graham Clark American 1832-1897

35. Ralph H. Fowler, UK 1889-1944 History of WDs • A better solution of the orbit gave us Sirius B’s mass (0.9Mo) This meant that this star was much denser. • In 1917 Adriaan van den Maaren discovered another WD (a single one). • It was not till quantum mechanics that Ralph Fowler determined that degeneracy pressure was supporting the WDs. • Eventually Chandrasekhar determined the mass limit above which WDs cannot exist (among many many other things that got him the Nobel Prize) Subrahmanyan Chandrasekhar Indian 1910-18995

36. WD density and pressure • You can determine the WD density from a value of its mass and radius (which for Sirius B come from observations). • You can then determine a value for the central pressure of the WD using the equation of hydrostatic equilibrium.

37. The mass-radius relation for WDsaka: the more you have the smaller it is! (No demonstration)

38. The Chandrasekhar limit for WDs • As you increase the mass of a WD the radius decreases, the density increases and eventually the electrons speeds approach the speed of light, i.e., they become relativistic. • By equating the core pressure for a star in hydrostatic equilibrium, to the pressure for relativistic electrons we see that there is no radius for which the star will be in equilibrium: the star just collapses. • From the same equation we can derive the mass of such star. For masses larger than the Chandrasekhar mass limit, the star collapses. (Board demonstration)

39. Neutron Stars: formation • Si-burning adds Fe to the core which increases in mass. And contracts under its own weight. • Electron degeneracy provides the pressure but when the electrons become relativistic at the Chandrasekhar limit the star collapses in a free-fall time (<1 second). • Protons and electrons combine to form neutrons and neutrinos (the neutrinos take energy out). • Collapse is halted by neutron degeneracy pressure. • We have a neutron star. • A neutron star is not technically a star because it is not gaseous and it is not powered by thermonuclear fusion….

40. Neutron stars: radius • The mass radius relations for WD is the same as that for neutron star except that the neutron mass is in place of the electron mass.

41. Neutron stars: escape velocity • The escape velocity from a NS is about ½ of the speed of light.

42. Neutron stars: luminosity • When first formed a NS has T~106K. So, despite their small radius the luminosity is not terribly low. • Using Wien’s law we see this spectrum peaks at ~30A or 400 eV in the X-ray range. • But there is another type of radiation from NS.

43. Neutron stars: pulsars • As the star collapses conservation of angular momentum makes it spin at about 0.1c, implying a rotation period of the order of milliseconds. • Magnetic flux is also conserved such that the surface B fields is intensified. • The rotating B field creates an E field that rips charged particles from the surface of the star, which later get beamed by the B field and ejected at the poles. • They were discovered during a radio survey of the Galaxy and the first one was named LGM-1. • There are about 500 known pulsars and considering the selection effects, there must be a lot more.

44. Neutron stars: pulsars • Why are pulsars not WDs? • If the periodicity was due to pulsations, the material pulsating would have to be very dense indeed, much denser than a WD can be. • If it were a rotating WD, the rotation speed at the periphery of the WD would have to be 100c. • Pulsars are often seen in the middle of SN remnants.

45. John Wheeler American 1911-2008 BH History • Postulated by John Mitchell 200 years ago as objects with such high g that not even light could escape (he called them black stars). • Calculated by Schwarzschild in 1916, as soon as GR was invented. • Died shortly after at the front of some disease– very sad. • John Wheeler called them Black Holes Karl Schwarzschild German 1873-1916

46. Black Holes: the Oppenheimer-Volkov limit • As for the Chandrasekhar limit there is a mass above which neutron degeneracy pressure fails to support the star. It is about 3Mo– hard to calculate because of the effects of the strong nuclear force.

47. Black Holes: the Schwarzshild radius • By equating the escape velocity to the speed of light you see the size of a BH’s event horizon. • Nothing going inside can emerge, not particles, not light, not information.

48. BH: tidal ripping • At what distance can a BH rip a person approaching feet first? (assuming that a force that can rip a person is about 16 tonnes) • For a 2000 Mo BH you will be ripped at exactly the Schwarzshild radius.

49. Observations of BHs • V404 Cyg – binary system where a K0 star donates mass to a dark object via an accretion disk. • The BH is revealed by the orbital motion of the companion or the luminosity of the accretion disk.