Monday, September 04, 2006

D. Koch: The Kepler Mission and Binary Stars (S 240 invited talk, Wed, Aug. 23)

The talk started with a detailed description of the Kepler mission concept. It is a photometric mission which will look for habitable planets (0.5 to 10 Earth masses). The instrument is a 1m Schmidt telescope with a very large field of view (FOV; a 42 CCD array larger than 100 square degrees - one would need six Palomar sky survey plates to cover the FOV). Kepler will observe in the Cygnus-Lyra region along the Orion arm of our Galaxy. It will achieve a differential photometric precision of 6.6 ppm for a 6.5 hour integration. A single FOV will be continuously observed for 4-6 years (except for less than 1 day every month). Simultaneous observations of more than 103000 main-sequence stars will be obtained (3000 "guest objects" can be proposed). The time resolution is 30 minutes for most stars, and one minute for a subset of 512 objects. Kepler will observe a single bandpass from 430 to 890 nm, where the cutoffs are set to avoid CaII H&K and fringing. The PSF FWHM will be about six arcsec.

The detection capability was also discussed - it depends on many different parameters.

The selected targets are late dwarf stars (F,G,K,M). A "Stellar Classification Program" is gathering new multi-band photometric data (SDSS griz filters + filter for Mg b lines). The "Kepler Project" is producing a catalog of stars in the target field.

Data processing will partly be performed on board (only pixels of interest are extracted). Ground processing will consist of cosmic ray, bias, smear and common-mode noise removal and the analysis of fluxes for threshold events. All light curves will be archived at STScI.

Other results from the data include results relevant to binary stars - follow-up radial velocity observations to get masses and to differentiate between planetary transits and grazing eclipsing binaries. 1000 to 1500 eclipsing binaries with high precision light curves are expected. Non-eclipsing binaries will be identified using astrometry (distance), effective temperature and luminosity. The data will also be useful for other astrophysical purposes (oscillations, etc.).

Community participation and access will be possible through the guest observer program. Scientists may propose to observe targets in the FOV not being observed as part of Kepler planet search. The objects can be brighter or fainter than the nominal dynamic range. The observation duration can be three months or longer.

The launch of Kepler is scheduled for October 2008.

More information is available on the Kepler website.

D. Pourbaix: Binaries in Large-Scale Surveys (S 240 invited talk, Wed, Aug. 23)

This talk gave a thorough overview on ground- and space-based surveys relevant to binary stars.

Starting with unresolved binaries, Wielen (1996) suggested two ways of unveiling unresolved binaries:
  • Variability-induced movers (VIM), for which the position of the photocenter moves back and forth along a line, and
  • Colour-induced displacement binaries (CID), for which one takes images through different filters and sees the displacement. This was illustrated by a plot for the five SDSS filters. The five positions of the photocenter will follow the path of the binaries (a line), ordered by wavelength (as opposed to random displacements for a single star). Thus, one can detect binaries with separations below the resolving power of the instrument.

This was followed by a list of eclipsing binary surveys:
  • Robotic Optical Transient Serach Experiment
  • Optical Gravitaional Lensing Experiment
  • All Sky Automated Survey
  • Discovery Channel Survey
and others. All surveys are essentially automated.

Then followed a discussion of the usefulness of general surveys with respect to binary stars. From 2MASS, one can get positions and IR magnitudes to confirm binarity, and it serves as a source of candidates.

From SDSS, binaries can be detected in several ways:
  • Spectrum analysis (Balmer lines from hot stars and TiO from a cool companion, or narrow H emission lines superposed on the WD spectrum)
  • Outliers in the color-color diagram
  • Spectroscopic binaries (but there are only 10000 objects with two or more spectra, the maximum number of spectra is 13)
  • CID binaries
  • Spectrum analysis: 747 detached close binaries (Silvestri et al. 2005)
  • Color outliers: 863 WD+MD pairs (Smolcic et al. 2004)
  • SBs: 675 candidates (19 orbits, Pourbaix et al. 2005)
  • CID binaries: 542 candidates (Pourbaix et al. 2004)

Other useful surveys are
  • High proper motion stars (bulge): A group at CfA and others have obtained more than 300 spectroscopic orbits.
  • Open clusters: Mermilliod & Paunzen (2003)
  • By-products of Planet Quest (late-type stars): Nidever et al.

Potentially useful surveys are
  • Guide star catalog II (proper motions, BVI magnitudes, positions)
  • Palomar Quest
  • CFHT Legacy Survey
  • UKIRT (looking for WD+BD and subgiant+BD binaries)
  • Radial Velocity Experiment (RAVE)
  • Pan-STARRS

Space missions:
  • IUE has observed about 150 spectroscopic orbits (Stickland et al.).
  • FUSE took over in 1999.
  • Hipparcos (double and multiple star annexes, DMSA), the so-called "Hipparcos Poor Astronomer Sky Survey" (What can be done with existing Hipparcos data?). The percentage of binaries among variable stars with a large brightness variation increases with V-I. This is a CIM effect. There is correlation between the position of the photocenter and the brightness. HIP 88848 served as an example for "changing a mess into science" (Fekel et al. 2005, third body). Torres et al. (2006) discusses the eclipsing binary V1061 Cyg. A poster by Halbwachs & Pourbaix was also mentioned. Makarov et al. discuss the identification of binaries comparing Tycho-2 proper motions with Hipparcos ones.
  • Gaia will enable the detection of binaries and multiple stars by astrometry, spectroscopy and photometry (variability, outliers). There will be enough data to investigate the binary frequency over the HR diagram, across stellar populations, etc.

  • Do not put all your eggs in the same pocket (do not survey only eclipsing binaries)
  • Someone's garbage can be science for others (e.g. stars in the SDSS)
  • Public archives are gold mines for immediate scientific results (enough data exist to give one paper for each of us) and training sets for future science
  • Do not push garbage recycling too much, beyond the specs (cf. the Hipparcos-planets controversy)

End of Convection posts

The previous post was the last one for the "Convection in Astrophysics" symposium. Note that some of the posters are available on the conference website as pdf files. A few more posts for the "Binary stars" symposium will follow.

Convection in Astrophysics, Session E, Invited Talk 1 (Wed, Aug. 23)

P.R. Wood: The Relation of Convection to Pulsation and Mass Loss in Red Giants

Concerning pulsations in red giants, the good news is that there have been one or two attempts of global modelling of convection in red giants. One is Woodward et al., where the star pulsates and shows a dipolar structure. The other is Freytag et al. (1999).

The bad news is that studies of pulsation in red giants use convection theories that belong in the past. They use spherically symmetric models and mixing length theory.

This introduction was followed by a discussion of the relation of convection and pulsational stability. In the red giant envelope, almost 100% of the energy flux is transported by convection. The driving of pulsation depends on the energy transport into and out of mass zones, hence it depends on convection. It follows that one has to know how convection varies throughout the pulsational cycle. The work integral shows that most work is done in the convective region.

The linear pulsation models by Xiong et al. (1998) show that turbulent pressure and turbulent viscosity introduce damping.

Pulsation in the presence of turbulent convection: Turbulent eddies transport momentum in a medium with a pulsation velocity gradient. This gives rise to turbulent viscosity.Wood mentioned that this can be seen as the macroscopic analogy of molecular viscosity, but in the discussion after the talk, Stein noted that turbulent viscosity is very different from molecular viscosity.

Nonlinear pulsation models show that work dissipated by turbulent viscosity is the major damping factor. How good are these models? Not very good. In general, pulsation models (linear and nonlinear) with current convective treatments are too unstable. Nonlinear models give large luminosity spikes that are not observed - the convection is too efficient (Olivier & Wood 2005).

On the observational side, microlensing surveys have given enough data that one can begin to do asteroseismology on red giants (MACHO, OGLE). Sample MACHO light curves were shown. They are often semiregular, multimode, and show a large amplitude range.

Next, the period-luminosity diagram (Wood et al. 1999) was discussed. Can one see different modes of oscillation? The periods and period ratios fit only approximately, which indicates that the envelope structure is inadequate. What is the cause of closely-spaced periods? Are they stochastic excitation by convection, or structural changes resulting from convective irregularities, or nonradial modes?

The closely-spaced periods were also discussed using O-C diagrams for red giants (O ... date of observed maximum, C ... predicted date). A straight line in these diagrams means that the period is constant. It was shown that red giant periods vary abruptly by a few percent.

Another way to look at closely-spaced periods are the so-called Peterson diagrams, which show period ratios vs. log(period). In these diagrams, the models do not fit well the observed ratios.

Period-luminosity relations by Lebzelter and Wood (2005) were also discussed. The importance of mass loss was stressed. But - what causes the mass loss?

For an ideal study of pulsations in red giants, one needs to follow the pulsation for 50 to 100 years.

As a final curiosity, red giants with long secondary periods were mentioned. The light amplitudes are up to 0.8 mag, the velocity amplitudes only a few km/s. Are they associated with the chromosphere?

R.F. Stein: Applications of Convection Simulations to Oscillation Excitation and Local Helioseismology (Review, Wed, Aug. 23)

This was the review talk of session E in the Convection Symposium: "Oscillations, mass loss, and convection".

One application of numerical simulations is to understand the excitation process for p-mode oscillations. The excitation equation contains an integral, in which the terms for pressure fluctuations and mode compression are multiplied (don't take mode compression out of the integral!). Excitation decreases at low and high frequencies, which results in typical oscillation periods, e.g. 5 minutes for the Sun.

In the power spectrum for the turbulence, spatial and temporal factors cannot be separated. Therefore, a generalized Lorentzian is fit (where the width and power depend on wavenumber). This is another improvement needed for analytical work.

In a discussion on turbulent (Pt) and gas pressure (Pg), an analysis of local heating and cooling shows that excitation is due to entropy fluctuations. Pt and non-adiabatic Pg work comparably near the surface, Pt extends deeper. p-mode driving is primarily by turbulent pressure. There is some cancellation between Pt and non-adiabatic Pg.

Next, some work by G√ľnther Houdek was shown, and his analytical model compared to simulations by Stein & Nordlund (2001). Excitation rates for p modes for the Sun, Procyon, and alf Cen A where shown - models and simulations agree.

Another application is helioseismology, i.e. to test and refine local helioseismic models. One needs a large size and a long time sequence (so far, the simulation is 48 Mm wide by 20 Mm deep, the intention is to double the size next month). As a bonus, the simulations lead to an understanding of supergranulation. The simulations run for 48 hours (1 turnover time), include f-plane rotation (surface shear layer), but no magnetic field yet. The resolution is 100 km horizontal, 12-70 km vertical.

After a presentation of the numerical method, some results were shown. The mean atmosphere is highly stratified. A slice through simulations for the velocity in the vertical plane shows that downflows are being swept aside by the diverging upflows. Thus, the downflows become larger and make it to deeper layers. Some of the downflows are not swept aside and disperse. Displaying the vertical velocity on horizontal planes shows the continuous change of scales from granules to supergranules. Upflows at the surface come from a small area at the bottom. Downflows at the surface converge to supergranule boundaries.

The horizontal and vertical velocities from simulations were then compared to MDI observations. At larger scales, the oscillatory component dominates, at smaller scales, the convective component dominates. Simulations were also compared to MDI observations in a k-omega diagram, a time-distance diagram, and in a diagram of f-mode travel times vs. simulated flow fields (divergence and horizontal). North-going and south-going travel time differences from MDI were compared to simulated velocities.

Note that the simulated data is available from the MDI data base: Full data sets (~200GB per hour solar time) and slices of Vxyz and T at selected depths (~2 GB per hour solar time).

After a movie of temperature and vertical velocity by Viggo Hansteen, showing non-linear wave propagation in 2D, some answers to the question "Why are linear calculations useful?" were given. They complement non-linear calculations, they are much faster than non-linear calculations, one can explore the parameter space, and one can isolate physical effects.

At the end of his talk, Stein mentioned that he is looking for a solar physics post-doc.

Convection in Astrophysics, Session D, Oral Contributions (Wed, Aug. 23)

H. Shibahashi: The DB gap of white dwarfs and semiconvection

This talk was given by Mike Thompson on behalf of Shibahashi.

First, a general overview of the classification and properties of white dwarfs (WDs) was given. The classification of WDs reflects their surface composition, but not their temperature. The spectra of DA WDs show only hydrogen lines (their atmospheres consist of pure H), those make up about 80% of all WDs. DB WDs show only He I lines (pure He). There are also DO (showing He II lines), DC, and other classes. DAs are found from hottest to coolest temperatures, DOs for Teff > 45000 K, DBs for Teff < 30000 K. No He-rich WDs are found between 45000 and 30000 K, this constitutes the DB gap.

A plot of the number of DA WDs vs. Teff from McCook & Sion (1987) catalog was shown. At about 11000 K, one can see a step in the ratio of DA/non-DA WDs.

Why does the DB gap exist? The simple picture of parallel sequences of H and He-rich objects doesn't work. There is a theory of spectral evolution (by Fontaine and Wesemael), which proposes a common origin for WDs (PNN). At 30000 K, He ionization creates convection, and He is mixed to the surface. In the spectral evolution model, a wide range of H layers is expected (10^-4 to 10^-13 solar masses).

In observations, there has been much progress recently, mainly by large surveys.

The new working hypothesis is that all WDs have some H. Only about 10^-15 (solar masses?) is needed to produce an optically thick H layer at the surface. For Teff > 45000 K, the He II/III zone creates turbulence which mixes H with He and leads to He stars. For gap stars, the He II/III zone is too deep to mix up H, and gravitational settling leads to the He envelope. At 30000 K, He I/II creates turbulence which mixes H with He and leads to He stars.

One can make a prediction of semiconvection based on this scenario. At 30000 K, the He ionization zone turns into a convectively stable layer, which is nontheless superadiabatic. This is a plane-parallel, gravitationally stratified layer of fluid in hydrostatic and radiative equilibrium, with a steep chemical gradient. The equations for this situation were shown, employing the Boussinesq approximation. The physical cause of the overstability is that radiative heat exchange brings about an assymmetry in the oscillary motion, and this leads to overshooting.

There are some open issues.

In summary, there are two groups of WDs, and a DB gap. Convective mixing and/or chemical separation might be responsible for the gap. A new type of WD variables is predicted near the red edge of the DB gap.

M. Spite: Extra-mixing in Extremely Metal-Poor red giants

This talk presented some results of the "First Stars" project. The aim of this project is an analysis of the chemical composition of the galactic matter in the early times. This works only if there is no mixing. 18 dwarf and 33 giant stars (not C-rich) with Fe/H <= -3.0 were selected, and high resolution spectra (R ~ 45000) obtained.

As an example, spectra of the Mg b lines and the NH band were shown. An LTE abundance analysis with OSMARCS models was performed.

For the discussion of Mg in turnoff stars (dwarfs and giants), a plot of [Mg/Fe] vs. [Fe/H] was shown. One could see that the abundance ratio is constant, with a scatter of about 0.1 dex. This is more or less the same for all elements. Exceptions are C and N in giant stars - [C/Fe] shows an extremely large scatter, and N is even worse.

What is the reason for the large scatter for C and N in giants? In a primordial scenario, it would mean that there was a real scatter in the ISM of the early Galaxy. In the in situ scenario, the original C and N abundances in the atmospheres of the giant stars have been altered.

Indicators of mixing in giants are:
  1. An anticorrelation of C and N: In a plot of [N/Fe] vs. [C/Fe], we see two groups of stars: mixed stars with [N/Fe] of about 1 and [C/Fe] <>
  2. A very low abundance of Li in mixed stars is expected. The Li abundance decreases as carbon 13 increases.
  3. The phenomenon must appear at a specific location in the HR diagram.

Next, comparison to the results of Gratton et al. (2000) was made. The mean metallicity of the Gratton et al. stars is -1.5 dex, whereas the mean metallicity of the "first stars" is -3.1 dex. Mixing appears at a higher luminosity for the "first stars", but in both cases at the location of the bump.

A discussion of the abundances of Na and Al in these stars showed that some mixed stars are Na or Al rich (none of the unmixed stars). This could be due to deep mixing. Maybe some mixed stars are AGB stars (no effect is seen in oxygen, maybe the effect is too small).

As soon as an extremely metal poor star reaches the luminosity of the bump, its atmosphere is mixed with the H burning layer and the abundances of the light elements are altered.

P.P. Eggleton: Two Instances of Convection and Mixing in Red Giant Interiors

The actual title of the talk was "Formation and destruction of 3He in low-mass stars - Big Bang nucleosynthesis rescued".

The discovery of a very important mixing process taking place on the AGB was presented. The mechanism is a Rayleigh-Taylor instability, driven by an unusual nuclear reaction of 3He. He is produced in the interior on the main sequence, then mixed into the convection zone.

The reaction equation is: 3He + 3He -> 4He + p + p

As an example, the evolution of an 0.8 solar masses population II star was shown in the HR diagram. The element distribution in the star at the end of the main sequence is as follows. There is a small exhausted core, and a lot of 3He is produced and mixed into the surface layers. There is a maximum in molecular weight at the bottom of the H burning shell at 3 million years after the turnoff (?). This leads to a little 3He burning shell.

For the simulations, the 3D hydrodynamics code "Djehuti" was used, which is described in Dearborn, Lattanzio and Eggleton (2006, ApJ 639, 405), a paper on the He flash. It implements explicit hydrodynamics, is run on 351 processors, and the timestep is Courant limited.

The Rayleigh-Taylor instability does not remove the molecular weight gradient that drives it. It is constantly replenished. Mixing advects fresh 3He in at the same rate as it advects products out.

The results were presented as a movie of the global stellar surface.