Jul 072017
 

By Brian Williams (STSCI).

In 1572 C.E., a “new star” appeared in the sky, gradually brightening to be as bright as Venus at peak brightness, staying visible for nearly two years. This event, now known as Tycho’s supernova, helped to usher in a new age of science where the heavens were no longer fixed. Nearly 450 years later, we know this object as Tycho’s supernova remnant (Figure 1). The shock wave from this supernova has expanded to a radius of nearly 4 parsecs and is still racing into the interstellar medium at over 6000 km/s in places. Because the remnant is so young (and our telescopes are so great!) a particularly exciting aspect of Tycho is that we can watch it expand over baselines of a few years. This video  shows the expansion of the remnant over five epochs of  observations between 2000 and 2015 taken by the Chandra X-ray observatory.

fig1-5

Figure 1: Snapshot of the video showing Tycho supernova remnant expansion (based on observations with the Chandra X-ray observatory).

Tycho’s supernova remnant is known to be the result of a Type Ia supernova, and while we have a lot of observational data, we do not yet have a great understanding of what the progenitor systems of these supernovae are. It seems that they result from the thermonuclear explosion of a white dwarf star in a binary system, but whether the companion star is another white dwarf or an AGB or red giant star is unknown (both models have pros and cons). We also do not understand how the explosion begins in the degenerate star and propagates through it. However, models of the explosion offer different predictions for the motion of the ejecta after the star has exploded. As an example, Figure 2 shows the velocity distribution of Silicon, an element produced in the supernova, for two different explosion models from Seitenzahl et al. (2013), where the difference between the two is the number of ignition points for the detonation inside the white dwarf. Plotted is the velocity of the ejecta along various cuts in the XYZ coordinate plane. As can be seen, the amount of symmetry in the velocities of the expanding ejecta is not the same between the two models.

apjaa7384f9_hr

Figure 2: Slices through all three coordinate axis planes for the velocity distribution of Si shortly after the explosion. A model with many ignition points for the detonation is shown on the left; a model with only a few is shown on the right.

The question then becomes whether we can constrain these models with observations. My collaborators and I studied new and archival observations of Tycho’s supernova remnant, identifying nearly sixty dense knots of ejecta from the supernova that had a measurable proper motion between the epochs. Of course, this only gives the two-dimensional motion in the plane of the sky, which is only part of the story. However, these ejecta knots, dominated by emission from silicon and sulfur, are also quite bright and Chandra’s CCD cameras produce a spectrum for every pixel. We were able to use the redshift and blueshift of the Si and S lines in the X-ray spectrum between 1.8 and 2.6 keV to measure the velocity along the line of sight. Measuring an accurate shift in the lines required a significant effort to ensure that the atomic physics was properly accounted for. (The rest energy of X-ray lines can be a function of the temperature and ionization state of the gas; see Williams et al. 2017 for more details on this).

Basic RGB

Figure 3: Histograms of the velocity distribution in the X (green), Y (blue), and Z (red) directions.

Our results showed that there is no measurable asymmetry in the expanding ejecta, meaning that, as far as we can infer from this data, the explosion of Tycho was symmetric. Models with more ignition points inside the white dwarf better reproduce the observations in this case. The knots have velocities ranging from just over 2000 km/s to well over 6000 km/s and the 3-D velocity vectors all point away from the center of the remnant (as one would expect). Our work represents the first 3D map of the expanding ejecta from a Type Ia supernova and is an important step in understanding these stellar explosions. We cannot stretch these results too far: Tycho is just one supernova remnant and there may well be more than one way to explode a Type Ia supernova. Still, results like this are encouraging and with Chandra and the next generation of X-ray telescopes studies of more Ia remnants are possible.

References: 

  • Seitenzahl, Ciaraldi-Schoolmann, Röpke et al. 2013, MNRAS, 429, 1156.
  • Williams, Coyle, Yamaguchi et al. 2017, ApJ, 842, 28.
May 272017
 

By Jeffrey Cummings (STSCI and Johns Hopkins University).

Stars evolve through many phases in their lifetime. Towards the end, when they have fused all of the Hydrogen (H) in their cores to Helium (He), they expand, cool down, and become giants. During these final stages, stars also begin to lose a significant amount of mass from their surfaces. A star like our Sun will lose approximately 50% of their mass, while more massive stars (> 8 Msun) will lose as much as 80% to 85%. Nearly all stars will eventually shed their outer layers and expose their hot and ultrahigh density cores. These remnants are called white dwarfs and, unless they experience mass transfer from a companion, their fate is to slowly cool with time.

However, a small percentage (approximately less than 2%) of high-mass stars can reach the necessary high densities and pressures in their cores to undergo a different fate: iron (Fe) will be created in their central core, but because it cannot undergo fusion to continue to generate energy, the star eventually becomes unstable and gravitationally collapses on itself creating a core-collapse supernova (CCSN). While it is commonly adopted in astronomy that stars with initial masses greater than 8 Msun will end their lives as a CCSN, we still do not fully understand at which stellar mass this transition occurs.

Knowledge of this transition mass is important because it sheds light on stellar evolution processes like mass-loss and convection/dredge up, and on the CCSN rate and how large of an effect these supernovae have on the production of elements (important to understand chemical evolution), on the energetics/feedback in galaxies, and on star formation.

Determination of the CCSN Mass Transition

There are three main methods to determine this mass. They all have their challenges and limitations, but jointly they may be able to begin to precisely constrain this transition mass:

1) Supernova studies. With the ability of the Hubble Space Telescope to resolve stars out to nearby galaxies, when a nearby CCSN occurs there is likely deep and resolved archival images of the progenitor star. Additionally, the stellar population of lower-mass stars that formed with the progenitor can subsequently be observed. Using stellar evolutionary models and the photometry of the progenitor and/or of the nearby lower-mass stars that formed with it, it is possible to infer the progenitor’s mass (e.g., Smartt 2009, Smartt 2015, Williams et al. 2014, Jennings et al. 2014). Within the past several years, the sample of nearby supernova has increased enough to provide meaningful statistics, indicating that the lowest-mass stars that undergo supernova have masses of ~9.5 Msun (Smartt 2015; Figure 1). These data also suggest that stars more massive than 18 Msun likely collapse but do not undergo an observable supernova.

BlogFig1Figure 1:  Nearby supernovae are displayed with their inferred initial masses and errors.  These masses are based on STARS and Geneva models (Eldridge & Tout 2004; Hirschi et al. 2004).  In solid black, a trend is fit to the data suggesting a lower limit of ~9.5 Msun.  In dashed black a second trend illustrates the expected number of stars at each initial mass based on a Salpeter initial mass function.  This shows that such a large number of 10 to 17 Msun supernovae progenitors should be accompanied by a smaller but observable number of higher mass supernovae progenitors.  This suggests that stars with mass greater than ~18 Msun may not undergo standard observable supernovae.

2) Stellar evolution models. A second method to estimate the CCSN transition mass is to use stellar evolution models to determine at which mass the core will reach the necessary conditions to create an Fe core and induce core-collapse (or produce an electron-capture CCSN). While a variety of models currently exist to account for the complex physical processes involved, many of them are beginning to converge around 9.3–9.8 Msun (Eldridge & Tout 2004, Poelarends 2007, Ibeling & Heger 2013, Doherty et al. 2015). Figure 2 illustrates the models from Doherty et al. (2015) and the transition masses at differing composition, with solar being represented by Z = 0.02.

BlogFig2

Figure 2: This diagram illustrates the evolutionary result of a star with an initial mass at a given composition (Z = total mass fraction of a star that is not H or He), where solar composition is Z = 0.02.  This model distinguishes between the different types of white dwarfs (WDs) and illustrates where white dwarf formation ends and supernovae begin.  These models suggest a narrow region where electron-capture supernovae (ECSN) occur followed by the traditional Fe-core CCSN and shows that composition is an important secondary variable in these mass transitions.

3) White dwarf studies. Another way to constrain the transition mass is to study the remnants of the stars that do not undergo CCSN, i.e. white dwarfs. Spectroscopic analysis of high-mass white dwarfs in star clusters, together with evolutionary models, allow us to both directly determine the mass of the white dwarf and to infer the mass of the white dwarf’s progenitor. This is known as the initial-final mass relation (IFMR). Figure 3 illustrates the IFMR derived from white dwarfs across multiple clusters and a broad range of masses (Cummings et al. 2015; 2016a; 2016b; in prep.), and with evolutionary timescales based on PARSEC stellar models (Bressan et al. 2012). The IFMR provides a direct constraint on the mass loss that occurs during stellar evolution and how it varies with initial mass. In terms of constraining the CCSN transition mass, there currently are very few ultramassive white dwarfs that have been discovered in star clusters, but our current project is to survey for more white dwarfs approaching the highest mass a white dwarf can stably have (the Chandrasekhar mass limit at 1.375 Msun). The initial mass of a star that will create a Chandrasekhar mass white dwarf will also define the CCSN transition mass, which will provide a critical check of the other methods. However, based on cautious extrapolation, we note that our current relation suggests a consistency with the CCSN mass transition occurring near 9.5 Msun.

BlogFig3

Figure 3: The current IFMR based on the spectroscopic analysis of white dwarfs that are members of star clusters.  Comparison of the spectroscopically determined mass and age of the white dwarf can in comparison to the age of the cluster be traced back to the mass of the star that would have formed a white dwarf at that time.  The solid line represents a fit to the data while the dashed line represents the theoretical IFMR of Choi et al. (2016).  This shows a relatively clean relation with the clear formation of larger and larger white dwarfs from larger and larger stars.  The highest-masses, however, remain poorly constrained, but precise measurement of where this relation reaches the Chandrasekhar mass (1.375 Msun; the upper limit of the y-axis), will define the progenitor mass at which stars begin to undergo CCSN.

The three methods described above, mostly independent from each other, are beginning to suggest a convergence around a CCSN transition mass of ~9.5 Msun. This implies that the commonly adopted value of 8 Msun, which is based on older and more limited data and theoretical models, has led us in the past to overestimate the number of CCSN by ~30%. Such an overestimation would greatly affect our understanding of the chemical evolution, energetics, and feedback in galaxies. Furthermore, the maximum mass for CCSN inferred by Smartt (2015), 18 Msun, further decreases the number of expected CCSN in a stellar population. This may solve the supernova rate problem that resulted from assuming that all stars with a mass greater than 8 Msun should undergo a CCSN. Under that assumption, the number of predicted CCSN is twice the observed rate (Horiuchi et al. 2011). But if only stars in the 9.5–18 Msun mass range undergo CCSN, based on the standard Salpeter stellar initial mass function, the new estimated CCSN rate is almost exactly half of the original estimate. This would bring observations and models into remarkable agreement.

Future work

Significant work remains to be done. More nearby supernovae need to be observed, with subsequent study of their progenitors, to further constrain their lower and upper mass limits. More ultramassive white dwarfs need to be discovered in star clusters, a focus of our research group, to refine our understanding of higher-mass stars and their mass-loss processes. Increasing the white dwarf sample size will provide an independent measurement of the CCSN transition mass, and it may also begin to provide an important observational test for the effects of differing stellar composition on mass-loss, evolution, and this mass transition. For stellar evolution models, a more coordinated approach between these observations and theory will improve the models, which also affects the inference of progenitor masses in the supernova and white dwarf studies. An iterative process, aiming for self-consistency across all steps, will ideally bring a more precise convergence of the mass transition at which supernova begin to occur. In any case, the assumption that all stars with mass greater than 8 Msun will undergo a CCSN is appearing more and more inaccurate.

References:

  • Choi, J., Dotter, A., Conroy, C., et al. 2016, ApJ, 823, 102
  • Cummings, J. D., Kalirai, J. S., Tremblay, P.-E., & Ramirez-Ruiz, E. 2015, ApJ, 807, 90
  • Cummings, J. D., Kalirai, J. S., Tremblay, P.-E., & Ramirez-Ruiz, E. 2016a, ApJ, 818, 84
  • Cummings, J.D., Kalirai, J.S., Tremblay, P.E., Ramirez-Ruiz, E., & Bergeron, P. 2016b, ApJ, 820, L18
  • Doherty, C.L., Gil-Pons, P., Siess, L., Lattanzio, J.C., & Lau, H.H.B. 2015, MNRAS, 446, 2599
  • Eldridge, J. J., & Tout, C. A. 2004, MNRAS, 353, 87
  • Hirschi, R., Meynet, G., & Maeder, A. 2004, A&A, 425, 649
  • Horiuchi, S., Beacom, J. F., Kochanek, C. S., Prieto, J. L., Stanek, K. Z., & Thompson, T. A. 2011, ApJ, 738, 154
  • Ibeling D., Heger A., 2013, ApJ, 765, L43
  • Jennings, Z. G., Williams, B. F., Murphy, J. W., Dalcanton, J. J., Gilbert, K. M., Dolphin, A. E., Weisz, D. R., & Fouesneau, M. 2014, ApJ, 795, 170
  • Poelarends A. J. T., 2007, PhD thesis, Astronomical Institute Utrecht (P07)
  • Smartt, S. J. 2009, ARA&A, 47, 63
  • Smartt, S.J. 2015, PASA, 32, 16
  • Williams, B. F., Peterson, S., Murphy, J., Gilbert, K., Dalcanton, J. J., Dolphin, A. E., & Jennings, Z. G. 2014, ApJ, 791, 105
Apr 012015
 

By John Debes, ESA/AURA astronomer at STScI

It is quite the April Fool’s prank to hide in plain sight, and that’s just what the rather pedestrian-appearing WD 1242-105 did for several years, masquerading as a single white dwarf.  The secret it held, and the story of how my team of researchers unraveled it, makes for a useful lesson in how science sometimes progresses—not always by careful predictions, but sometimes by serendipity.  Furthermore, our discovery also demonstrates how astronomy often requires close collaboration.

WD 1242-105 (See Figure 1), resides near the constellation Virgo, and was first discovered as part of the large Palomar-Green Survey of UV-excess sources (Green et al., 1986).  In that survey, it was promptly misclassified as a subdwarf star.  Sixteen years later, Salim & Gould (2002) recognized that this might be a white dwarf candidate based on its apparent motion on the sky—it was large compared to its faint apparent magnitude.  This is often how new white dwarfs are discovered, since astrophysical objects closer to the Earth have larger apparent motions, and white dwarfs are intrinsically faint.

index

Figure 1: A false color image of WD 1242-105 (center), which is a rather inconspicuous star.  The hint of its high motion on the sky is given by the slight blue/red color–this is due to the fact that the binary had moved between when the two photographs of this star were taken.

Surprisingly, low-resolution spectra were taken and did not detect anything unusual about the white dwarf (Kawka et al. 2004; Kawka & Vennes 2006).  Based on this spectrum, it appeared to be a relatively close, single white dwarf that was about 75 lightyears away.  By a quirk of fate, there was a large high resolution spectroscopic survey of white dwarfs looking for binary objects and using the Very Large Telescope, but this particular white dwarf was not included.

This is where I came in.  As a postdoctoral researcher within the Department of Terrestrial Magnetism (DTM) at the time, I had access to the premiere high resolution spectrograph at Las Campanas Observatory’s Magellan Telescope, called MIKE.  I was using it to survey as many southern hemisphere white dwarfs as I possibly could for traces of metals in their atmospheres.  Since no one had yet published a high resolution spectrum of WD 1242-105, and it appeared to reside within the solar neighborhood, it was a perfect target.  I consequently observed the star three times, each exposure separated by ten minutes.

When I finally got around to analyzing the spectrum of WD 1242-105 a few months later, I was in for a shock.  Instead of the usual single spectral line of Hydrogen (See Figure 2) one sees in white dwarf spectra, there were two separate lines—indicative of a binary system consisting of two hydrogen white dwarfs.  By a sheer stroke of luck, I had caught the binary when both stars were at their maximal relative velocities to our line of sight—the lines were separated by almost 6.5 Angstroms, or 300 km/s.  They were moving so fast, that I could detect changes in their radial velocities on the timescale of each exposure, or ten minutes!

Fit_specphoto_primary

Figure 2: (left) A comparison between synthetic model white dwarf spectra (red lines) with the MIKE observations of WD 1242-105.  The two Hydrogen line components are most visible around the Hα spectral line. (right) A comparison between our model photometry and the observed photometry of this star from optical to mid-IR wavelengths.

Over the next year, I prepared to take more spectra, but by this time I needed to call in some favors.  This binary could have been a progenitor of the famous Type Ia supernovae, and so far not many convincing progenitors have been found.  Many of my colleagues at DTM also used MIKE, and I asked for help in taking additional spectra.  Others were experienced at taking high precision time series photometry of stars, and so I asked them to monitor WD 1242-105.  Finally, another friend of mine was working on a project to measure the parallaxes of nearby stars for evidence of planetary systems—I enlisted that team’s help to measure the distance to this interesting binary.

In the end we had a complete picture for this system’s binary parameters, which enabled us to estimate the individual white dwarf masses (See Figure 3).  We also were able to estimate the masses of the two components using the hydrogen lines themselves as tracers of the white dwarfs’ temperatures and gravities.  Since this was a binary, the distance obtained from our parallax measurement was actually about 60% larger than originally assumed, or 40 pc  (130 light years).  The binary has an orbital period just shy of 3 hours, and a total mass of 0.9 M.  While its mass makes it too light to be a progenitor of a Type Ia supernova, the two components of WD 1242-105 will merge within the next 800 million years in what will no doubt be a remarkable show.  Binaries like WD 1242-105 are believed to be the progenitors of  objects such as R Corona Borealis stars, which are believed to arise from the merger of two white dwarfs.

velocity_fig

Figure 3: Radial Velocity curve of the WD 1242-105 double degenerate binary system.  Red points are derived from the less massive component, while the blue points represent the more massive white dwarf.  The second panel shows the residuals after subtracting off the orbital fit to the data.

Since it is so close to Earth, it is one of the largest known sources of expected gravitational wave radiation at frequencies within the mHz range.  Unfortunately, this is too low a frequency for detection with some proposed space-based future gravitational wave observatories such as eLisa, but nevertheless, it will “shine” brightly in the gravitational wave sky, despite its rather bland optical appearance.

Sometimes, the most interesting things are hiding in plain sight, and they just require the right approach or instrument to discover them.

 

 

The paper related to WD 1242-105 can be found at this link.

Other references:

Green, R. F., Schmidt, M., & Liebert, J. 1986, ApJS, 61, 305

Kawka, A. & Vennes, S. 2006, ApJ, 643, 402

Kawka, A., Vennes, S., & Thorstensen, J. R. 2004, AJ, 127, 1702

Salim, S. & Gould, A. 2002 ApJ, 575, L83

 

You can follow John Debes on twitter, @JohnDebes

Feb 032015
 

By Annalisa Calamida, Postdoctoral Fellow at STScI

The Milky Way bulge is the closest galaxy bulge that can be observed and studied in detail. The bulge could include as much as a quarter of the stellar mass of the Milky Way (Sofue et al. 2009), and the characterization of its properties can provide crucial information for the understanding of the formation of the Galaxy and similar, more distant galaxies. We observed the Sagittarius low-reddening window, a region of relative low extinction in the bulge, E(B-V)  ≤ 0.6 mag (Oosterhoff & Ponsen 1968), with the Wide-Field Channel of the Advanced Camera for Surveys (ACS) and the Wide Field Camera 3 (WFC3) on board the Hubble Space Telescope (HST). The time-series images were collected in the F606W and F814W filters and cover three seasons of observations from October 2011 to October 2014. The main goal of the project led by Dr. Kailash Sahu is to detect isolated stellar mass black holes and neutron stars in the Galactic bulge and disk through gravitational microlensing.

However, these data, covering 8 WFC3 and 4 ACS fields and including a sample of about 2 million stars, are a gold mine for other Galactic bulge stellar population studies. I started a project aimed at characterizing the bulge stellar populations through the study of different evolutionary phases. The work is based on the same data set used for the microlensing project and on observations taken with ACS in 2004. By combining the observations of one of the ACS fields with those taken in 2004, we measured very precise proper motions, better than ~ 0.5 mas/yr (~ 20 Km/s) at F606W ~ 28 mag, in both axes. Proper motions allowed us to separate disk and bulge stars and to obtain the deepest clean color-magnitude diagram of the bulge. As a consequence we identified for the first time a clearly defined white dwarf cooling sequence in the Galactic bulge (Calamida et al. 2014).

The characterization of the white dwarf population is an effective method to understand the formation history of the bulge and the Milky Way itself, since most stars end their life as white dwarfs. The white dwarf population of the bulge contains important information on the early star formation history of the Milky Way. Our knowledge is most extensive for the white dwarf population of the Galactic disk, in which numerous white dwarfs were discovered through imaging surveys, such as the Sloan Digital Sky Survey (Eisenstein et al. 2006, Kepler et al. 2007, Kleinman et al. 2013). White dwarfs have been identified and studied in a few close Galactic globular clusters too, thanks to deep imaging with HST (Richer et al. 2004; Hansen et al. 2007; Kalirai et al. 2007; Calamida et al. 2008; Bedin et al. 2009). Many of the Milky Way bulge stars are metal rich – several of them even reaching super solar metallicity – which means that they are similar to the Galactic disk population, and the bulge stellar space density is closer to that of the disk than that of globular clusters. However, many of the bulge stars are old, like globular cluster stars, which means that they show common properties with both the disk and the cluster populations. Understanding whether the white dwarfs in the Galactic bulge more closely resemble those found in either the disk or the clusters is therefore an important part of developing our understanding of how the Milky Way was formed, as well as the bulge formation, and, indeed, the nature of the Galactic bulge itself.

The color-magnitude-diagram in the F606W and F814W filters of selected bulge stars is shown in Fig.1, where confirmed white dwarfs are marked with larger filled dots. We used theoretical cooling tracks from the BaSTI database (Pietrinferni et al. 2004; 2006) and models from Althaus et al. (2009) to fit the bulge white dwarf cooling sequence, and a distance modulus DM0 = 14.45 mag and reddening E(B-V) = 0.5 mag (Sahu et al. 2006). These models assume different core and atmospheric compositions: CO-core white dwarfs with an hydrogen-rich envelope (DA), CO-core white dwarfs with an helium-rich envelope (DB) and He-core white dwarfs. Assuming an age of about 11 Gyr and an average solar chemical composition for bulge stars, isochrones predict a turn-off mass of ~ 0.95 M, and through the initial-to-final mass relationship, a white dwarf mass of ~ 0.53 – 0.55 M (Weiss & Ferguson 2009, Salaris et al. 2010).  The figure below shows cooling tracks for DA (dashed blue line) and DB (dashed green) CO-core white dwarfs with mass M = 0.54 M and He-core white dwarfs (dashed red) with mass M = 0.23 M plotted on the Galactic bulge cooling sequence. The DA and DB CO-core cooling tracks are unable to reproduce the entire color range of the observed white dwarf cooling sequence. An increase in the mean mass of the white dwarfs would move these models towards bluer colors, further increasing the discrepancy. The lower mass He-core cooling track for M = 0.23 M fits the red side of the bulge white dwarf sequence (note that empirical evidence shows that the lower mass limit for white dwarfs is ~ 0.2 M, Kepler et al. 2007). These results support the presence of a significant fraction (~ 30%) of low-mass (M ≤ 0.45 M) He-core white dwarfs in the Galactic bulge. According to stellar evolution models, in a Hubble time, these low-mass white dwarfs can only be produced by stars experiencing extreme mass loss events, such as in compact binaries. Among the brighter very red white dwarfs we found indeed one ellipsoidal variable (marked with a blue dot in Fig.1), probably composed of a white dwarf accreting from a main-sequence companion, and two dwarf novae (magenta dots). The fainter counterparts of these binaries could populate the region where the reddest white dwarfs are observed in the color-magnitude diagram. These systems could be composed of a white dwarf and a low-mass (M <  0. 3M) main-sequence companion. This hypothesis is further supported by the finding of five cataclysmic variable candidates in the same field (green dots).

cmd_bulge_blog

Figure 1: proper-motion cleaned bulge color-magnitude diagram in the F606W and F814W filters. White dwarfs are marked with larger filled dots. Dashed lines display cooling tracks for CO- and He-core white dwarfs. The ellipsoidal variable and the dwarf novae are marked with blue and magenta dots, respectively, while green dots mark cataclysmic variable candidates. Error bars are also labeled.

Our observational campaign ended in November 2014 and now, by adding the third season of observations and including all the twelve ACS and WFC3 fields, we will be able to increase our sample of stars by one order of magnitude. The increased statistics will allow us to better constrain the nature of the white dwarf population in the bulge, for instance, through comparing white dwarf and main-sequence star counts.

In the future, this study will greatly benefit from the advent of the James Webb Space Telescope (JWST), which will have an improved sensitivity and spatial resolution compared to HST. Moreover, JWST will observe in the near-infrared, where the extinction is a factor of ten lower compared to the optical. The new telescope will allow us to observe the bulge white dwarfs in the near-infrared, something that is now barely feasible with HST. The white dwarf cooling sequence will then be used to estimate the age of stellar populations in the bulge, to be compared to estimates obtained by adopting other diagnostics, such as the main-sequence turn-off. This study will be fundamental for constraining the presence of an age spread in the Galactic bulge, which is now a hotly debated topic.

 

For more details see Calamida et al. 2014, ApJ, 790, 164

 

References:

  • Sofue et al. 2009, PASJ, 61, 227
  • Oosterhoff & Ponsen 1968, BANS, 3, 79
  • Calamida et al. 2014, ApJ, 790, 164
  • Eisenstein et al. 2006, ApJS, 167, 40
  • Kepler et al. 2007, MNRAS, 375, 1315
  • Kleinman et al. 2013, ApJS, 204, 5
  • Richer et al. 2004, AJ, 127, 2904
  • Hansen et al. 2007, ApJ, 671, 380
  • Kalirai et al. 2007, ApJ, 671, 748
  • Calamida et al. 2008, ApJ, 673, L29
  • Bedin et al. 2009, ApJ, 697, 965
  • Pietrinferni et al. 2004, ApJ, 612, 168
  • Pietrinferni et al. 2006, ApJ, 642, 797
  • Althaus et al. 2009, A&A, 502, 207
  • Sahu et al. 2006, Nature, 443, 534
  • Weiss & Ferguson 2009, A&A, 508, 1343
  • Salaris et al. 2010, ApJ, 716, 1241
Apr 142014
 

By Pier-Emmanuel Tremblay, Hubble Fellow at STScI

White dwarfs represent the endpoint of stellar evolution for 95% of all stars. At the present day in our Galaxy, the large majority of stars that were born slightly more massive than the Sun are in their final remnant stage. These degenerate stars are slowly cooling as they lose their internal energy through radiation. We study them both for the purpose of understanding these condensed matter laboratories, and for enhancing their use as probes of fundamental astrophysical relations, such as the expansion of the Universe. The study of white dwarfs in clusters, routinely done by HST, provides very precise ages for the first stellar populations in our Galaxy. By linking the final white dwarf mass to the initial mass of its progenitor, it is also possible to calibrate the core mass growth and stellar lifetime of asymptotic giant branch (AGB) stars [1].

Most of the mass in a C/O white dwarf is a mixture of carbon and oxygen, and there is usually a thin layer of hydrogen (less than 0.01% of the mass) floating at the surface. As a consequence, most degenerate stars have a pure-hydrogen atmosphere. The most accurate method to determine the atmospheric parameters (the effective temperature and surface gravity) of H-rich white dwarfs is to compare the observed line profiles of the hydrogen Balmer lines with the predictions of detailed model atmospheres (Figure 1) [2]. Nevertheless, there was a long-standing problem [3] where cool remnants (0.2 < Cooling Age [Gyr] < 10) with a convective atmosphere have masses up to 20% higher than warmer non-convective objects, which impacts the use of white dwarfs as cosmochronometers.

 

Fig1Figure 1: Observed spectra of the white dwarf WD 1053−290 with a simultaneous fit of the Balmer lines, from Hβ to H8, with a 3D model spectrum. Line profiles are offset vertically from each other for clarity and the best-fit atmospheric parameters are identified at the bottom of the panels. The instrumental resolution is of 6 Å. Source: Tremblay et al. (2013b)

 

We have recently computed the first grid of 3D model atmospheres [4] for hydrogen-atmosphere white dwarfs (Figure 2) in order to improve the convection model. These CO5BOLD [5] radiation-hydrodynamics simulations, unlike the previous 1D calculations, do not rely on the mixing-length theory or any free parameter for the treatment of convective energy transfer.

 

Fig2Figure 2: Snapshot of a 3D white dwarf simulation at effective temperature Teff = 12,000 K and log g = 8. Left: temperature structure for a slice in the horizontal-vertical xz plane through a box with coordinates x,y,z (in km). The temperature is color coded from 60 000 (red) to 7000 K (blue). The arrows represent relative convective velocities, while thick lines correspond to contours of constant Rosseland optical depth, with values given in the figure. Right: emergent bolometric intensity at the top of the horizontal xy plane. The root-mean-square intensity contrast with respect to the mean intensity is 18.8%. Source: Tremblay et al. (2013a)

 

The 3D simulations have been employed to compute 3D spectra for the Balmer lines which were then used in the spectroscopic analysis of the white dwarfs in the Sloan Digital Sky Survey [6]. White dwarfs with radiative and convective atmospheres have derived mean masses that are now the same (Figure 3), in much better agreement with our understanding of stellar evolution. Indeed, both cool and warm degenerates in the Galactic disk are expected to originate from the same populations, but from stars that have formed at slightly different times. We are now in the process of using the 3D simulations as upper boundary conditions for structure models, in order to predict improved ages and more precise ZZ Ceti pulsation properties. We will also improve the metal abundance determinations for white dwarfs that are accreting former disrupted planets in their convective zone.

 

Fig3Figure 3: Mass histograms for DA stars in the Sloan Digital Sky Survey sample with Teff < 40 000 K (black empty histogram) from 1D (top) and 3D spectra (bottom). We also show the sub-distributions for radiative atmospheres (13 000 < Teff (K) < 40 000, blue histogram) and convective atmospheres (Teff < 13 000 K, red histogram). The mean masses and standard deviations are indicated in the panels in units of solar masses. Binaries and magnetic objects were removed from the distributions. Source: Tremblay et al. (2013b)

 

 

 

 

 

 

 

 

 

 

References:

[1] Kalirai, J. S., Marigo, P., & Tremblay, P.-E. 2014 (ApJ, 782, 17)
[2] Bergeron, P., Saffer, R. A., & Liebert, J. 1992 (ApJ, 394, 228)
[3] Bergeron, P., Wesemael, F., Fontaine, G., & Liebert, J. 1990 (ApJL, 351, L21)
[4] Tremblay, P.-E., Ludwig, H.-G., Steffen, M., & Freytag, B. 2013a (A&A, 552, A13)
[5] Freytag, B., Steffen, M., Ludwig, H.-G., et al. 2012 (Journal of Computational Physics, 231, 919)
[6] Tremblay, P.-E., Ludwig, H.-G., Steffen, M., & Freytag, B. 2013b (A&A, 559, A104)