A540 – Stellar Atmospheres Organizational Details Meeting times Textbook Syllabus Projects Homework Topical Presentations Exams Grading Notes

Basic Outline Textbook Topics –Chapter 1 – Review of relevant basic physics –Chapter 3 – Spectrographs –Chapter 4 - Detectors –Chapter 5 – Radiation –Chapter 6 – Black bodies –Chapter 7 – Energy transport –Chapter 8 – Continuous Absorption –Chapter 9 – Model Photospheres –Chapter 10 – Stellar Continua –Chapter 11 – Line Absorption –Chapter 12,13 – Spectral Lines –Chapter 14 – Radii and Temperatures –Chapter 15 - Pressure –Chapter 16 - Chemical Analysis –Chapter 17 – Velocity Fields –Chapter 18 - Rotation Integrating Stars – –Stars in the astrophysical zoo – –Stellar activity – –Winds and mass loss – –White dwarf spectra and atmospheres – –M, L and T dwarfs – –Non LTE – –Metal poor stars – –Pulsating stars & Asteroseismology – –Supergiants – –Wolf-Rayet stars – –AGB stars – –Post-AGB stars – –Chemically Peculiar Stars – –Pre-main sequence stars – –Binary star evolution – –Other ideas…

Goals Familiarity with basic terms and definitions Physical insight for conditions, parameters, phenomena in stellar atmospheres Appreciation of historical and current problems and future directions in stellar atmospheres

History of Stellar Atmospheres Cecelia Payne Gaposchkin wrote the first PhD thesis in astronomy at Harvard She performed the first analysis of the composition of the Sun (she was mostly right, except for hydrogen). What method did she use? Note limited availability of atomic data in the 1920’s

Useful References Astrophysical Quantities Holweger & Mueller 1974, Solar Physics, 39, 19 – Standard Model MARCS model grid (Bell et al., A&AS, 1976, 23, 37) Kurucz (1979) models – ApJ Suppl., 40, 1 Solar composition – "THE SOLAR CHEMICAL COMPOSITION " by Asplund, Grevesse & Sauval in "Cosmic abundances as records of stellar evolution and nucleosynthesis", eds. F. N. Bash & T. G. Barnes, ASP conf. series, in press: see also Grevesse & Sauval 1998, Space Science Reviews, 85, 161 or Anders & Grevesse 1989, Geochem. & Cosmochim. Acta, 53, 197 Solar gf values – Thevenin 1989 (A&AS, 77, 137) and 1990 (A&AS, 82, 179)

What Is a Stellar Atmosphere? Basic Definition: The transition between the inside and the outside of a star Characterized by two parameters –Effective temperature – NOT a real temperature, but rather the “temperature” needed in 4  R 2  T 4 to match the observed flux at a given radius –Surface gravity – log g (note that g is not a dimensionless number!) Log g for the Earth is 3.0 (10 3 cm/s 2 ) Log g for the Sun is 4.4 (2.7 x 10 4 cm/s 2 ) Log g for a white dwarf is 8 Log g for a supergiant is ~0 Mostly CGS units…

Make it real… During the course of its evolution, the Sun will pass from the main sequence to become a red giant, and then a white dwarf. Estimate the radius of the Sun (in units of the current solar radius) in both phases, assuming log g = 1.0 when the Sun is a red giant, and log g=8 when the Sun is a white dwarf. What assumptions are useful to simplify the problem?

Basic Assumptions in Stellar Atmospheres Local Thermodynamic Equilibrium –Ionization and excitation correctly described by the Saha and Boltzman equations, and photon distribution is black body Hydrostatic Equilibrium –No dynamically significant mass loss –The photosphere is not undergoing large scale accelerations comparable to surface gravity –No pulsations or large scale flows Plane Parallel Atmosphere –Only one spatial coordinate (depth) –Departure from plane parallel much larger than photon mean free path –Fine structure is negligible (but see the Sun!)

Basic Physics – Ideal Gas Law PV=nRT or P=NkT where N=  /  P= pressure (dynes cm -2 ) V = volume (cm 3 ) N = number of particles per unit volume  = density (gm cm -3 ) n = number of moles of gas (Avogadro’s # = 6.02x10 23 ) R = Rydberg constant (8.314 x 10 7 erg/mole/K) T = temperature in Kelvin k = Boltzman’s constant 1.38 x 10 –16 erg K -1 (8.6x10 -5 eV K -1 )  = mean molecular weight in AMU (1 AMU = 1.66 x gm) Don’t forget the electron pressure: P e = N e kT Densities, pressures in stellar atmospheres are low, so the ideal gas law generally applies.

Make it real… Using the ideal gas law, estimate the number density of atoms in the Sun’s photosphere and in the Earth’s atmosphere at sea level. For the Sun, assume P=10 5 dyne cm -2. For the Earth, assume P=10 6 dyne cm -2. How do the densities compare?

Thermal Velocity Distributions RMS velocity = (3kT/m) 1/2 Most probable velocity = (2kT/m) 1/2 Average velocity = (8kT/  m) 1/2 What are the RMS velocities of 7 Li, 16 O, 56 Fe, and 137 Ba in the solar photosphere (assume T=5000K). How would you expect the width of the Li resonance line to compare to a Ba line?

Excitation – the Boltzman Equation g is the statistical weight and  is the difference in excitation potential. For calculating the population of a level the equation is written as: u(T) is the partition function (see def in text). Partition functions can be found in an appendix in the text. Note here also the definition of  = 5040/T = (log e)/kT ) with k in units of electron volts per degree (k= 8.6x10 -5 eV K -1 ) since  is normally given in electron volts.

Ionization – The Saha Equation The Saha equation describes the ionization of atoms (see the text for the full equation). P e is the electron pressure and I is the ionization potential in ev. Again, u 0 and u 1 are the partition functions for the ground and first excited states. Note that the amount of ionization depends inversely on the electron pressure – the more loose electrons there are, the less ionization. For hand calculation purposes, a shortened form of the equation can be written as follows

Make it real… At (approximately) what Teff is Fe 50% ionized in a main sequence star? In a supergiant? What is the dominant ionization state of Li in a K giant at 4000K? In the Sun? In an A star at 8000K?

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

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…… kpc HD 93129A O3If 3.4 kpc Eta Carina B kpc Cyg OB2# B5 Ia + e1.7 kpc Zeta-1 Sco B1.5 Ia kpc

Wolf-Rayet Stars Luminous, hot supergiants Spectra with emission lines Little or no hydrogen L sun Maybe 1000 in the Milky Way Losing mass at high rates, to 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?

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 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?

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?

Solar Type Stars (F, G, K) Pulsators –The delta Scuti stars, etc. –SX Phe stars Binaries –FK Comae Stars –RS CVn stars –W UMa stars –Blue Stragglers

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

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?)

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)

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 ~ M sun per year, km/sec –Dust, gas cocoons (IRC ) some 10,000 AU in diameter Semi-regular and irregular variables (SRa, SRb, SRc) –Smaller amplitudes –Less regular periods, or no periods

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 –…

Amplitude of Mira Light Curve

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

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

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

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

R Corona Borealis Stars A-G type Supergiants Suddenly become much fainter (8 mag) He, Carbon rich, H poor “Dust puff theory” - Mass loss and dust obscuration? Origin - Double degenerate (He + CO with mass transfer)? about 100 known

White Dwarf Merger Scenario The camera aspect remains the same, but moves back to keep the star in shot as it expands. After the star reaches 0.1 solar radii, an octal is cut away to reveal the surviving disk and white dwarf core. The red caption (x) is a nominal time counter since merger. A rod of length initially 0.1 and later 1 solar radius is shown just in front of the star. (Saio & Jeffrey -

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