1. The Stellar Zoo Across the HR diagram: What causes an ordinary star to become weird? basic stellar evolutionbasic stellar evolution mass loss & windsmass loss & winds diffusion & radiative levitationdiffusion & radiative levitation pulsation (radial and non-radial)pulsation (radial and non-radial) rotationrotation mixingmixing magnetic fieldsmagnetic fields binary evolution & mass transferbinary evolution & mass transfer coalescencecoalescence

2. The Upper Upper Main Sequence 100 (or so) solar masses, T~20,000 – 50,000 K Luminosities of 10 6 L Sun Generally cluster in groups (Trapezium, Galactic Center, Eta Carinae, LMC’s R136 cluster) Always variable – unstable. (Some of) The Brightest Stars in the Galaxy StarmVmV MVMV M bol Sp. T.Dist. Pistol Star……-11.8 7 kpc HD 93129A7.0-7.0-12O3If 3.4 kpc Eta Carina6.2-10-11.9B0 02.5 kpc Cyg OB2#1211.5-10-10.9B5 Ia + e1.7 kpc Zeta-1 Sco4.7-8.7-10.8B1.5 Ia + 1.9 kpc

3. Wolf-Rayet Stars Luminous, hot supergiants Spectra with emission lines Little or no hydrogen 10 5 -10 6 L sun Maybe 1000 in the Milky Way Losing mass at high rates, 10 -4 to 10 -5 M sun per year T from 50,000 to 100,000 K WN stars (nitrogen rich) Some hydrogen (1/3 to 1/10 He) No carbon or oxygen WC stars (carbon rich) NO hydrogen C/He = 100 x solar or more Also high oxygen Outer hydrogen envelopes stripped by mass loss WN stars show results of the CNO cycle WC stars show results of helium burning Do WN stars turn into WC stars?

4. WR First identified by French Astronomers Charles Wolf (1827 - 1918) and Georges Rayet (1839 - 1906) in 1867,  candidate object to produce a sufficiently energetic supernova called a hypernova mass at least 20 times that of the Sun expel large amounts of heated gas such as helium in occasional bursts. In this way, WR stars are very "windy" with stellar wind velocities much greater than that of the Sun. The expulsion of so much gas means a WR star is surrounded by an "atmosphere" of gas that is comparable in size to the star itself.  In this way, really just seeing the ionized gas surrounding the star. This gas is so hot it emits visible light, along with a variety of other kinds of photons (radiation) ranging from less energetic radio waves and microwaves to more energetic ultraviolet light

5. WR WR stars emit so much material because of the occasional upwelling of heavier elements fused inside the core, including carbon and oxygen. These heavier elements interrupt the energy flowing outward from the star by absorbing it. Eventually, this causes a tremendous outward pressure that takes the form of a powerful stellar wind blowing off the star.  A typical WR star can lose a mass equivalent to that of the Earth in a year. Because WR stars burn so brightly, fiercely, and erratically, they have very short life spans. A typical WR life span is just a few million years. When a WR star dies, it often undergoes a supernova detonation. Some of these supernovae are sufficiently energetic to classify as a hypernova. And some of these hypernovae are accompanied by a gamma-ray burst as the core undergoes an uneven gravitational collapse into the most awesome object in the Universe -- a black hole, an object whose gravitational pull within a certain distance is so powerful nothing, not even light, can escape

6. More Massive Stars Luminous Blue Variables (LBVs) –Large variations in brightness (9-10 magnitudes) –Mass loss rates ~10 -3 M sun per year, transient rates of 10 -1 M sun per year –Episodes of extreme mass loss with century-length periods of “quiescence” –Stars’ brightness relatively constant but circumstellar material absorbs and blocks starlight –UV absorbed and reradiated in the optical may make the star look brighter –Or dimmer if light reradiated in the IR –Hubble-Sandage variables are also LBVs, more frequent events –Possibly double stars? –Radiation pressure driven mass loss? –Near Eddington Limit?

7. LBV LBVs are massive, intrinsically bright stars which display different scales of light and colour variability, ranging from rapid microvariations to rare outbreaks of catastrophic mass loss. They represent a very short-lived (perhaps as little as 40,000 years) strongly mass-losing phase in the evolution of massive stars, during which they undergo deep erosion of the outer layers before they enter the Wolf-Rayet phase.Wolf-Rayet They include P Cygni stars, S Doradus stars, and Hubble-Sandage variables.P Cygni The LBV class was first defined by Conti (1984). The spectra, luminosity and temperature are highly variable. During quiet periods of minimum brightness (so-called ‘quiescence’) they appear as blue B- type supergiants with temperatures greater than around 15 - 20,000 K showing H and HeI emission. During periods of maximum brightness they appear as later A- or F-type supergiants with temperatures near 8000 K and FeII and [FeII] emission reaches maximum. It is widely believed that the bolometric magnitudes, i.e. total radiation, remain roughly constant throughout these cycles and that the apparent cooling is due to absorption by surrounding shells of ejected matter which re-radiate the stars’ visible and blueward radiation at longer wavelengths (Parker et al. 1993).

8. Chemically Peculiar Stars of the Upper Main Sequence Ap stars (magnetic, slow rotators, not binaries, spots) –SrCrEu stars –Silicon Stars –Magnetic fields –Oblique rotators Am-Fm stars (metallic-lined, binaries, slow rotators) –Ca, Sc deficient –Fe group, heavies enhanced –diffusion? HgMn stars The  Boo stars Binaries?

9. CP At least 25% (Schneider, 1993) of the upper main sequence stars is known as spectroscopically peculiar stars chemically peculiar (CP) stars, implying that unusual chemical composition (element abundances) rather than other causes is responsible for the spectrum anomalies. On and near the main sequence, for spectral types from B to early F, one finds a remarkable diversity of the stellar surface properties and variability. In cooler and hotter parts of the H-R diagram a single, powerful process, such as convection in solar-type stars or mass loss in hot massive stars, dominates the physics of stellar atmospheres. In contrast, several processes of comparable magnitude compete in the A-star atmospheres and envelopes, creating interesting and heterogeneous stellar population.

10. CP The radiative diffusion (Michaud 1970) is the most important process responsible for non-solar surface chemical composition. The diffusion theory suggests that ions heavier than hydrogen are able to levitate or sink under competing influence of the radiation pressure and gravity. Element segregation by the radiative diffusion is easily wiped out by various hydrodynamical mixing effects and, thus, requires a star which is stable over significant part of its outer envelope. Slowly rotating B-F stars with shallow convection zones provide the required stability. The presence of strong, global magnetic field contributes further to the suppression of turbulence and leads to different diffusion velocities depending on the field inclination and strength. Chemically peculiar stars are separated into the two distinct sequences according to their magnetic properties. Am, λBoo, HgMn stars lack strong magnetic fields and show mild chemical anomalies in chemically homogeneous outer stellar layers. Ap and Bp stars have magnetic fields exceeding few hundred gauss, exhibit extreme chemical anomalies and have substantial vertical and horizontal chemical gradients in the photosphere.

11. Pulsating Bootis stars λ Boo stars are Population I early-A to early-F type stars which exhibit significant underabundance of most iron-peak and heavy elements but show solar abundances of CNO and some other light elements (Paunzen et al. 2002a; Heiter 2002). These chemical properties are believed to arise from contamination of the shallow stellar surface convection zones by the accretion of metal-depleted gas from a circumstellar shell (Venn & Lambert 1990) or a diffuse interstellar cloud (Kamp & Paunzen 2002). High-resolution time-series spectroscopy by Bohlender et al. (1999) revealed the presence of high- degree non-radial pulsations in the majority of investigated λ Boo stars λBoo stars tend to pulsate in high-overtone modes. one object with λ Boo chemical characteristics – the planetary host star HR8799 – is known to exhibit the γ Dor type pulsational variability (Zerbi et al. 1999; Gray & Kaye 1999). However, HR 8799 appears to be an exception as other members of the γ Dor group show normal abundance pattern (Bruntt et al. 2008).

12. Pulsating Bootis stars λ Boo stars are Population I early-A to early-F type stars which exhibit significant underabundance of most iron-peak and heavy elements but show solar abundances of CNO and some other light elements (Paunzen et al. 2002a; Heiter 2002). These chemical properties are believed to arise from contamination of the shallow stellar surface convection zones by the accretion of metal-depleted gas from a circumstellar shell (Venn & Lambert 1990) or a diffuse interstellar cloud (Kamp & Paunzen 2002). High-resolution time-series spectroscopy by Bohlender et al. (1999) revealed the presence of high- degree non-radial pulsations in the majority of investigated λ Boo stars λBoo stars tend to pulsate in high-overtone modes.

13. non-magnetic HgMn chemically peculiar stars The non-magnetic HgMn chemically peculiar stars present another challenge for our understanding of the excitation of pulsations in hot stars. Many HgMn stars are situated within the SPB instability strip. Furthermore, an increased opacity due to accumulation of metals by radiative diffusion in HgMn stars is expected to enhance the driving of the SPB pulsations (Turcotte & Richard 2002). Contrary to this theoretical prediction photometric observations show no evidence of pulsational variability in HgMn stars (Adelman 1998). Spectroscopic line profile variations detected for a handful of HgMn stars is limited to lines of 2–3 heavy elements and is, consequently, attributed to chemical inhomogeneities rather than pulsation (Adelman et al. 2002; Kochukhov et al. 2005; Hubrig et al. 2006b). Incompleteness of the theoretical diffusion models in the outer part of the stellar envelope is the most likely explanation for the contradiction between predicted and observed pulsation properties of HgMn stars

14. Rapidly oscillating magnetic Ap stars Rapidly oscillating Ap (roAp) stars represent the most prominent subgroup of pulsating chemically peculiar stars These objects belong to the SrCrEu type of magnetic A stars, and pulsate in high-overtone, low degree p-modes. roAp stars are found at or near the main sequence, at the cool border of the region occupied by the magnetic Ap/Bp stars (Kochukhov & Bagnulo 2006). According to the series of recent spectroscopic studies (e.g., Ryabchikova et al. 2002, 2004; Kochukhov et al. 2002a), effective temperatures of roAp stars range from about 8100 down to 6400 K. Their atmospheres are characterized by diverse chemical abundance patterns, but typically have normal or below solar concentration of light and iron-peak elements and a very large overabundance of rare-earth elements (REEs). Similar to other cool magnetic A stars, roAp stars possess global fields with a typical strength from few to ten kG (Mathys et al. 1997), although in some stars the field intensity can exceed 20 kG (Kurtz et al. 2006b)

15. Spectroscopy of roAp pulsations High-quality time-resolved spectra of roAp stars have proven to be the source of new, incredibly rich information, which not only opened new possibilities for the research on magneto-acoustic pulsations but yielded results of wide astrophysical significance. Numerous spectroscopic studies of individual roAp stars (e.g., Kochukhov & Ryabchikova 2001a; Mkrtichian et al. 2003; Ryabchikova et al. 2007a), as well as comprehensive analysis of pulsational variability in 10 roAp stars published by Ryabchikova et al. (2007b),.  The most prominent characteristic of the RV oscillation in roAp stars is the extreme diversity of pulsation signatures of different elements. Only a few stars show evidence of <50 ms −1 variation in the lines of iron-peak elements, whereas REE lines, especially those of Nd ii, Nd iii, Pr iii, Dy iii, and Tb iii exhibit amplitudes ranging from a few hundred ms −1 to several kms −1.

16. Asteroseismology of roAp stars the κ mechanism acting in the hydrogen ionization zone, with the additional influence from the magnetic quenching of convection and composition gradients built up by the atomic diffusion (Balmforth et al. 2001; Cunha 2002; Vauclair & Th´eado 2004).  theories cannot reproduce the observed temperature and luminosity distribution of roAp stars and have not been able to identify parameters distinguishing pulsating Ap stars from their apparently constant, but otherwise very similar, counterparts (Th´eado et al. 2009).  but some success has been achieved in calculating magnetic perturbation of oscillation frequencies (Cunha & Gough 2000; Saio & Gautschy 2004) and inferring fundamental parameters and interior properties for multiperiodic roAp stars (Matthews et al. 1999; Cunha et al. 2003)

17. Solar Type Stars (F, G, K) Pulsators –The delta Scuti stars, etc.(Population I) –SX Phe stars(Population II) –  exist in the lower part of the classical (Cepheid) instability strip Binaries –FK Comae Stars (caused by large, cool spots) –RS CVn stars –W UMa stars –Blue Stragglers

18. RS CVn & FK com convection coupled with high stellar rotation results in a dynamo mechanism which converts the mechanical energy of rotation and convection into magnetic energy. This dynamo is what powers solar and stellar activity. RS CVn are late-type, evolved (spectral type ~ K) stars found in binary systems. they have deep convection zone. Also, because they are in binary systems, tidal coupling with the companion forces the star to rotate at high velocity.  Consequently, these stars show magnetic activity in the form of starspots ( image via Doppler imaging), flares, coronae, X-ray and chromospheric emission, etc. Because of the deeper convection zone and high rotation rates for RS CVn stars, the level of activity is many orders of magnitudes greater than solar activity. For instance, sunspots only cover a few percent of the solar surface even during sunspot maximum. On RS CVn stars starspots can cover 10-20% of the surface. FK Comae stars are single stars, late-type stars which are believed to be coaelesced binaries. the high rotation stems from the one-time binary nature of the star. Starspots, like sunspots are believed to trace the surface magnetic flux. Doppler imaging of these features on stars gives us valuable information about spot morphology, differential rotation, and activity cycles.

19. W UMa - contact eclipsing binary star W UMa is an eclipsing binary star : the prototype star of an entire class of variable star.  an eclipsing binary, consists of two stars that are so close together that their period is on the order of 8 hours

20. Blue Stragglers In Core of globualr clusters

21. Blue Stragglers Blue straggler stars in the core of Globular clusters Stellar cannabalism in binary systems (whereby gas is drawn from a companion star) has gathered support from the "binary evolution" or "slow coalescence" model (more below) as the primary formation mechanism for blue stragger stars

22. Evidence of the slow coalescence model evidence that six of 43 blue stragglers in globular cluster 47 Tucanae formed by the slow coalescence of contact binary stars. The six blue stragglers contain less carbon and oxygen than the remaining 37, which suggests that the material at the surface of six came from deep inside their parent stars. In theory, however, this deep material can move to the surface of a blue straggler only during the mass transfer process occurring between two stars in a binary system, which is supported by numerical simulations

23. Boesgaard & Tripicco 1986: Fig 2 The famous lithium dip!

24. The Lower Main Sequence – UV Ceti Stars M dwarf flare stars About half of M dwarfs are flare stars (and a few K dwarfs, too) A flare star brightens by a few tenths up to a magnitude in V (more in the UV) in a few seconds, returning to its normal luminosity within a few hours Flare temperatures may be a million degrees or more Some are spotted (BY Dra variables) Emission line spectra, chromospheres and coronae; x-ray sources Younger=more active Activity related to magnetic fields (dynamos) But, even stars later than M3 (fully convective) are active – where does the magnetic field come from in a fully convective star? These fully convective stars have higher rotation rates (no magnetic braking?)

25. On to the Giant Branch… Convection 1 st dredge-up LF Bump Proton-capture reactions CNO, Carbon Isotopes Lithium Gilliland et al 1998 (47 Tuc)

26. Real Red Giants Miras (long period variables) –Periods of a few x 100 to 1000 days –Amplitudes of several magnitudes in V (less in K near flux maximum) –Periods variable –“diameter” depends greatly on wavelength –Optical max precedes IR max by up to 2 months –Fundamental or first overtone oscillators –Stars not round – image of Mira –Pulsations produce shock waves, heating photosphere, emission lines –Mass loss rates ~ 10 -7 M sun per year, 10-20 km/sec –Dust, gas cocoons (IRC +10 216) some 10,000 AU in diameter Semi-regular and irregular variables (SRa, SRb, SRc) –Smaller amplitudes –Less regular periods, or no periods

27. Pulsators Found in many regions of the HR diagram Classical “Cepheid Instability Strip” –Cepheids –RR Lyrae Stars –ZZ Ceti Stars “Other” pulsators –Beta Cephei Stars –RV Tauri –LPVs –Semi-Regulars –PG 1159 Stars –Ordinary red giants –…

28. Amplitude of Mira Light Curve

29. More Red Giants Normal red giants are oxygen rich – TiO dominates the spectrum When carbon dominates, we get carbon stars (old R and N spectral types) Instead of TiO: CN, CH, C 2, CO, CO2 Also s-process elements enhanced (technicium) Double-shell AGB stars Peery 1971

30. Weirder Red Giants Weirder Red Giants S, SC, CS stars –C/O near unity – drives molecular equilibrium to weird oxides Ba II stars –G, K giants –Carbon rich –S-process elements enhanced –No technicium –All binaries! R stars are warm carbon stars – origin still a mystery –Carbon rich K giants –No s-process enhancements –NOT binaries –Not luminous for AGB double-shell burning RV Tauri Stars

31. Mass Transfer Binaries The more massive star in a binary evolves to the AGB, becomes a peculiar red giant, and dumps its envelope onto the lower mass companion Ba II stars (strong, mild, dwarf) CH stars (Pop II giant and subgiant) Dwarf carbon stars Nitrogen-rich halo dwarfs Li-depleted Pop II turn-off stars

32. After the AGB Superwind at the end of the AGB phase strips most of the remaining hydrogen envelope Degenerate carbon-oxygen core, He- and H-burning shells, thin H layer, shrouded in dust from superwind (proto-planetary nebula) Mass loss rate decreases but wind speed increases Hydrogen layer thins further from mass loss and He burning shell Star evolves at constant luminosity (~10 4 L Sun ), shrinking and heating up, until nuclear burning ceases Masses between 0.55 and 1+ solar masses (more massive are brighter) Outflowing winds seen in “P Cygni” profiles Hydrogen abundance low, carbon abundance high (WC stars) If the stars reach T>25,000 before the gas/dust shell from the superwind dissipates, it will light up a planetary nebulae Temperatures from 25,000 K on up (to 300,000 K or even higher) Zanstra temperature - Measure brightness of star compared to brightness of nebula in optical hydrogen emission lines to estimate the uv/optical flux ratio to get temperature

33. R Corona Borealis star Edward Pigott(1797) noticed that it had disapeared but over succeeding months it graduallly reappeared. In the 1930s, it was established by Loreta and O'Keefe that the occasional dimming is due to clouds being ejected from the surface of the star. Observed minima and maxima therefore do not occur on a periodic basis, and are dependent upon line of sight. R Cr B stars are carbon stars, and are generally characterized by extreme hydrogen deciency and pulsate with periods between 40 and 100 days Theories : 1. a merger of two white dwarfs, 2. a late Helium shell ash in a post-ABG star

34. R Corona Borealis Stars A-G type Supergiants Suddenly become much fainter (8 mag) He, Carbon rich, H poor Many lines of neutral atomic carbon and strong bands of molecular carbon(C2, CN) “Dust puff theory” - Mass loss and dust obscuration? Origin - Double degenerate (He + CO with mass transfer)? about 100 known

35. RCB Relativiely low masses (<1 M) Mv ~ -4~-5 LMC RCBs 7000K ~ 5500 K H ~ 1/1000 ~ 1/10 6 H of sun 1. Final flash or last thermal pulse model or 2. He WD + CO WD : canibalised by the more massic CO WD =Double Degenerate(DD) or Mergerd Binary White Dwarf(MBWD) model degenerate C, O core containng ~90% of mass Energy from a nuclear-burning shell at the bottom of a tenuous He-C enveleope core 0.01Ro, overall R ~ 50Ro

36. Other H-defficient stars 1.WR and h-deficient binaries 2. Hydrogen-deficient Carbon stars 3. Extreme Helium stars and Hot RCB stars 4. Hydrogen-deficient central stars of planetary nebulae 5. Carbon stars : cool supergiants with C rich spectra 6. Born-again stars 1) V605 Aql :now hydrogen-deficient central star of a planetary nebular. In 1919 ti brightened as a slow nova & a report that its spectrum at one time resembled an RCB star 2) FG Sge : faint blue star in 1900, progressively redder and brighter, now spectrum of an F or early-G type giant. 1992 ;show RCB type fading events  new-born RCB star 3) Sakurai’s object of V4334 Sgr ;faint blue star suddenly brightened in 1996 ;new-born RCB star www.telf-ast.demon.co.uk www.telf-ast.demon.co.uk www.aavso.orgwww.aavso.org

37. White dwarf merger origin for extreme helium star V652 Her V652 Her has a surface made almost entirely of helium, Every two and a half hours it expands and contracts by over 15,000 km (about 1% of its size).  These pulsations enable us to measure V652's mass Simon Jeffery 

38. White Dwarf Merger Scenario (Saio & Jeffrey - http://star.arm.ac.uk/~csj/movies/merger.html) The merger theory postulates that the binary orbit will shrink due to the loss of orbital angular momentum, either through magnetic interaction between the stars and their atmospheres or by the generation of gravitational waves. There is currently no satisfactory theory which will predict reliable decay times, but they are clearly long. For this simulation, they define an orbital decay time (t d ) of a few orbital periods, 10 in the current case.http://star.arm.ac.uk/~csj/movies/merger.html

39. White Dwarf Soup Single Stars –DO (continuous) –DB (helium) –DA (hydrogen) –DZ (metals) –DC (carbon) Evolutionary sequence still unclear Cataclysmic Variables –WD + low mass companion –Neutron star + companion –Accretion disk