Aug 012016
 

By Anthony Marston, European Space Agency/STScI

What are Wolf-Rayet stars?

Wolf-Rayet (WR) stars are believed to be evolved massive stars that initially started their lives with masses of  > 20 Msun. With such high masses, they evolve very quickly to the WR state from high-mass hydrogen burning O stars in 1-2Myr. Currently, evidence suggests that the majority of WR stars are either in or affected by having been in relative close binaries, that can affect their evolution.

There are several evolutionary paths and theories as to the evolutionary direction of WR stars. It is postulated that different evolutionary paths exist depending on the how much initial mass exists above 20 Msun, as well as whether they are single or binary stars. For most WR stars, a mass-loss phase of a few tens of thousands of years probably occurs. Evidence for this is seen in the nitrogen-enriched ejecta nebulae that are seen around many WR stars. Ejecta are believed to be associated with a slow wind phase following the fast wind of the main sequence O star phase. Once evolved to a WR star there is again a fast wind phase which can quickly interact with a slow moving ejecta nebula. But not all WR stars are seen to have ejecta.

There are three subtypes of WR stars: WN subtypes show prominent nitrogen emission lines in their spectra, WC subtypes show prominent carbon emission lines, and WO subtypes show strong high excitation oxygen emission lines. These form an apparent evolutionary sequence with the spectra showing the products of Hydrogen burning for WN stars, Helium burning for WC star spectra and higher level burnings for WO stars. WO subtype stars in the Galaxy are very rare (three are known) and probably represent a final WR evolutionary phase before becoming a supernova (probably of type Ib).

How common are they and how are they distributed in the Galaxy?

By the end of 2000 just over 200 WR stars were known in the Galaxy. Most of these were discovered in studies of clusters or serendipitously. They were shown to follow the spiral pattern of the Galaxy and showed a distribution that mimicked other Galactic star formation site indicators. Indeed, WR stars have in various ways been used as markers of very recent high-mass star formation and star formation bursts since they only live a few million years.

In his review of WR stars, Karel van der Hucht (2001) indicated that the likely population of WR stars in the Galaxy could be several thousand rather than the few hundred known. This was in part due to obvious observational restrictions, such as unseen populations on the opposite side of the Galaxy. With the advent of sensitive infrared detectors the possibility of finding distant and/or obscured WR stars became more realistic. Two approaches have been developed for finding Galactic WR stars in recent years.

The “narrow-band” approach (Shara et al. 2009) uses narrow-band images centered on strong emission lines seen in WR stars (e.g. HeII) and subtracts from them narrow-band images covering only the continuum (or broad-band infrared images). The candidates revealed are followed up with infrared and/or optical spectroscopy to confirm their nature.

The “broad-band” approach is based on the near- to mid-infrared colors which are peculiar to stars with strong winds – and in particular WR stars. Figure 1 shows how the free-free emission from the fast WR wind of the nearby WR star WR11 (g Vel) has a distinct spectral index which is substantially different from stellar photospheres leading to WR stars being overabundant in certain areas of broadband infrared (2MASS, Spitzer/IRAC, WISE) point source color-color space (see Figure 2). Even though the this approach is slightly more prone to confusion issues than the “narrow-band” method, it has a couple of advantages over the latter: the potential for picking up weak-lined WR stars or ones where lines are diluted relative to the continuum due to a massive companion or local hot dust emission. It also uses already existing infrared point source catalogs (e.g. the GLIMPSE catalog of source within | b < 1 | in the Galactic plane). As of July 2016, the total number of known Galactic WR stars is 634 (http://www.pacrowther.staff.shef.ac.uk/WRcat/).

figure1Figure 1: Spectral energy distribution of g Velorum (Williams et al. 1990) showing the excess free-free emission from the stellar wind in the infrared wavelengths as compared to photospheric emission (straight black line). The GLIMPSE catalog which used Spitzer/IRAC data will show WR stars with colors distinct from the vast majority of stars.

Our group, with core members Schuyler Van Dyk (Caltech), Pat Morris (Caltech), Jon Mauerhan (UC Berkeley) and Anthony Marston (ESA-STScI), uses the broad-band method. It was first developed by Marston in 2004 to identify candidates in ESO/SOFI infrared spectroscopic observations and it helped identify 60 new WR stars by Mauerhan et al (2011). The color selection uses data from the GLIMPSE catalog, consisting of several 10’s of million sources detected in the Galactic plane using broadband Spitzer/IRAC 3.4 – 8 mm measurements combined with band-merged flux data from 2MASS (broadband near-infrared JHKs). In certain studies, X-ray emission sources and, more lately, WISE point source colors have been used in identifying WR candidates.  Spectroscopic follow-up has concentrated on obtaining K-band spectra, as WR stars are typically identified by strong HeII emission lines such as the 2.189mm line. For the less reddened candidates, optical spectroscopic follow-up has also been possible.

Historically we have found that 10-15% of candidates turn out to be bona fide WR, stars while~ 85% of all candidates are emission-line stars, most often Be stars. Small numbers of O/Of stars B[e] supergiants and stars exhibiting infrared CO bandhead absorption lines have been picked up where combinations of photosphere, dust emission and free-free emission has brought objects into our infrared color space. Improvements to our color-space selection have increased the success rate of WR detections out of the candidates, notably for more reddened/distant objects where the candidate confirmation rate can go as high as 25% (see Figure 2). We are currently looking into a machine-learning capability for assessing the likelihood of an object being a WR star from color-space criteria. The ultimate goal is to be able make accurate predictions of WR numbers of different subtypes in the Galaxy.

figure2

Figure 2: Infrared color-color plots showing the candidate objects observed by Mauerhan et al (2011). The green symbols were newly discovered WN subtypes and red WC subtype stars. Blue points represent candidates that follow up spectroscopy showed were not WR stars. Grey shaded areas indicate the part of the color-color plots where 50% or more candidates were found to be WR stars.

What have WR stars taught us about high-mass star formation?

The number ratios of WR to O stars and Red Supergiant (RSG) or Luminous Blue Variable stars are key values for constraining stellar evolution theories of massive star evolution. In a simple way, ratios provide an indication of relative timescales for lifetimes. Another indication of timescales, and possibly different evolutionary links between subtypes, mass-loss phases and initial stellar masses, is the number distribution of WR subtypes, both WN and WC (plus the rare WO stars).

The distribution of WR stars (studied e.g. using the Spitzer’s GLIMPSE survey across the Galactic plane) marks star formation sites across the Galaxy and indicate likely sites of future supernovae. However, it has become clear over time that many WR stars, that are no more than a few Myr in age, appear to be found well away from the centers of star-forming clusters in the Galaxy. A projection of most of the known WR stars with secure distance shows that some WR stars also appear to be more than 100 pc above/below the plane of the Galaxy (Rosslowe & Crowther, 2015). A possible explanation of why some WR stars appear to be located away from their birth site could be the presence of fast transverse motions caused by expulsion from their cluster formation site. Another possibility could be that these stars were part of small clusters but, being much more luminous than other cluster members, they appear to be isolated. But in recent years, in the study of star-forming regions like the Cygnus OB2 cluster, we have learned of an unexpected third possible explanation.

Various studies suggest that the Cygnus OB2 cluster, being 1 Myr of age, has not evolved significantly from its original distribution. This means that the massive stars, and WR stars in particular, are near the sites where they were born. However, none of the WR stars are in the massive star cluster at the center of the Cyg OB2 association (see Figure 3), and not only that, none show evidence of bow shocks from significant transverse velocities, suggesting these stars were born in situ. We now know, from studies with Herschel, that filaments of high-density gas can extend through star-forming regions with “strings” of star-forming cores being found along them. And in fact, filaments pervade the Cyg OB2 area leading to the possibility of forming high-mass stars outside of major stellar clusters, possibly instigated to form high-mass stars through a triggering event, such as expanding gas shell collisions.

fig3

Figure 3: The Cygnus OB2 association as seen by Herschel PACS/SPIRE (colored background from Schneider et al, 2016). WR stars and Luminous Blue Variable stars (likely precursors of WR stars in stellar evolution) are found well away from the major cluster of O stars shown as white points a bit to the right of center of the field (Comeron et al 2008, Wright et al 2015).

There are therefore two possible scenarios:

  • WR stars are born in situ and away from stellar clusters (but likely within stellar associations) – which means distributed high-mass star formation occurs for some of the most massive stars probably from filaments.
  • WR stars are kicked out of stellar clusters due to dynamics of the early cluster of stars or through binary/supernova interactions, apparently affecting a large fraction of the very massive stars in the stellar cluster.

As we have seen, the study of Wolf Rayet stars has shed new light on unexpected physical processes associated to high-mass star formation. In the future, we will advance their study by: (1) Using machine-learning and improved color-selection techniques to find new WR stars and assess their distributions in the Galaxy, including in high-mass star-forming regions. (2) Pinning down number ratios of WR subtypes and other massive star types. (3) Using the GAIA catalog to get proper motions of WR stars to identify runaway stars. (4) Searching for bow shocks, in particular in the mid-infrared with WISE, as it has been found that they are particularly prominent at IR wavelengths.

 

References:

  • Blaauw,  A.,1993, ASPC, 35, 207
  • Comeron et al, 2008, A&A, 486, 453
  • Crowther et al, 2006, MNRAS, 372, 1407
  • van der Hucht, K., 2001, New AR, 45, 135
  • Mauerhan, J., et al, 2011, AJ, 142, 40
  • Rosslowe, C., & Crowther, P., 2015, MNRAS, 447, 2322
  • Schneider, N., et al, 2016, A&A, 591, A40
  • Shara et al, 2009, AJ, 138, 402
  • Williams, P., et al, 1990, MNRAS, 244, 101
  • Wright, N., et al, 2015, MNRAS, 449, 741

 

Feb 152016
 

By Andrew Fox, ESA-AURA Astronomer at STScI

Like a superpower that imports gas to satisfy its energy needs, the Milky Way depends on fresh gas supplies to fuel its star formation. Our galaxy converts gas to stars at a rate of about one or two solar masses per year, and has enough gas reserves to continue forming stars at this rate for about two billion years. But two billion years is not that long compared to the life of a Galaxy, so without replenishment of its gas supplies, the star formation would eventually cease.

Fortunately for the Milky Way, there is a reservoir of gas to import, seen in the form of the so-called high-velocity clouds (HVCs), a population of gaseous objects that orbit in the halo of the Galaxy. Although not all HVCs are inflowing, they provide an average inflow rate of about one solar mass of gas per year, close to the star-formation rate. HVCs allow us to study the process of Galactic accretion, the delivery of gas to the Galactic disk where star formation takes place.

One of the best-characterized HVCs is the Smith Cloud, named after Gail Bieger (née Smith), who discovered it via its radio emission in her PhD thesis in 1963. We know how far away the Smith Cloud is (13 kpc), how fast it is moving (300 km/s), how much gas it contains (two million solar masses), and even what its orbit looks like: it came out of the Galactic disk about 70 million years ago, and is due to impact the disk about 27 million years from now. It has a comet-like appearance (see Figure 1) and is fragmenting as it interacts with the surrounding gas in the Galactic halo.

sc-zolt2

Figure 1: The Smith Cloud, as observed in neutral hydrogen 21 cm emission from the Green Bank Telescope, superimposed on an optical image of the Milky Way. The Cloud is ~3×1 kpc in size and is moving toward the Galactic disk. This image is color-coded by the intensity of 21 cm emission. Credit: Saxton/Lockman/Levay/NRAO/AUI/NSF/STScI.

Despite our solid understanding of the Smith Cloud, one key property was (until recently) unknown: its chemical composition. This is a vital clue to its origin, since a cloud containing low levels of heavy elements (in the jargon, low metallicity) likely originates outside the Galaxy, perhaps from the intergalactic medium or from a dwarf galaxy. On the other hand, a cloud originating in the Galaxy should have much higher metallicity, close to the levels measured in the Sun, because heavy elements are forged in the cores of massive stars, so their presence indicates prior star formation has taken place.

We used the Cosmic Origins Spectrograph on the Hubble Space Telescope to determine the Smith Cloud’s metallicity and therefore constrain its origin. By using the technique of spectroscopy, the ultraviolet light from three background active galactic nuclei (AGN) was split into different wavelengths, and the absorption-line signatures caused by sulfur atoms in the Smith Cloud were measured. We also used radio observations from the Green Bank Telescope in West Virginia to measure how much neutral hydrogen exists in each direction. Combining the sulfur and hydrogen measurements in the three sightlines allowed us to measure the metallicity of the Smith Cloud, and we found it to be one-half of the solar value. This experiment is illustrated in Figure 2.

sc-cartoon

Figure 2: Illustration of our experiment to determine the Smith Cloud’s metallicity (courtesy Ann Feild/STScI). Ultraviolet and radio observations of three lines-of-sight through the Cloud toward background active galactic nuclei (AGN) allow us to measure the abundance of sulfur atoms.

One-half solar is too high a metallicity for the Cloud to represent a dwarf-galaxy neighbor of the Milky Way (and besides, the Cloud contains no stars). Yet it is too low to have originated in the disk of the Galaxy near the Sun, where we expect solar metallicity. However, one-half solar matches the abundances in the outer disk of the Milky Way. Indeed, if we trace back the orbit of the Smith Cloud to where it last crossed the disk, we find this occurred at about 13 kpc out from the Galactic Center (for comparison, the Sun is at about 8.5 kpc). In other words, both the metallicity and orbit of the Smith Cloud are consistent with it being flung out of the outer disk of the Milky Way, about 70 million years ago. This scenario is illustrated in Figure 3.

sc-external

Figure 3: The orbit of the Smith Cloud. The Cloud is being shaped by gravitational and gas-pressure forces. The Cloud’s kinematics and metallicity suggest an origin in the outer disk about 70 million years ago. 30 million years from now, the cloud is expected to return to the disk. Courtesy Ann Feild.

These results confirm that the Smith Cloud is made of Galactic material. But they do not explain how the Cloud was launched into its orbit. Could it have been blown out of the outer disk by a cluster of supernovae? Supernovae are known to drive winds of gas and dust out of galaxies, and this would naturally explain the high velocity of the Cloud. However, the Cloud is much more massive than those observed in the vicinity of known Galactic supernovae explosions, so it would have had to be an unusually energetic supernova event.

A more exotic possibility invokes regions of dark matter, known as mini-halos, which are thought to be constantly bombarding the disk of the Milky Way. If one of these mini-halos were to accumulate Galactic gas during its passage through the disk, it could create an object like the Smith Cloud: a blob of Galactic gas held together by the gravity of unseen dark matter. Far-fetched as this may sound, mini-halos are predicted by theoretical work on galaxy formation, with simulations expecting that tens or hundreds of such mini-halos should exist for every large galaxy like the Milky Way. This creates a fascinating picture of the disk of the Galaxy, where gas can be transported from one location to another by catching a ride on a dark mini-halo. More research will explore this possibility, particularly by predicting what the observational signatures of these mini-halos should be, and seeing whether the observed morphology of high-velocity clouds can be reproduced.

 

For more reading, check out:

Fox et al. 2016, ApJL, 816, L11 (reporting the above results)

Galyardt & Shelton 2016, ApJL, 816, L18 (discussing dark-matter mini-halos)

Lockman et al. 2008, ApJL, 679, L21 (background on the Smith Cloud)

Jan 152016
 

By Miguel Requena-Torres, Postdoc at STScI

Star formation is a beautiful event that goes on in galaxies. This phenomenon can be quiet or violent. In our own Galaxy, for example, there are quiet clouds of material that evolve slowly to form filamentary structures with growing pockets of gas and dust that eventually acquire enough density to start gravitational collapse, leading to the birth of stars. But our Galaxy also harbors violent regions where cataclysmic events inject enough energy into the surrounding medium that can trigger star formation.  In the last few years, new theories have been developed to explain the formation of stars in our own Galaxy by studying the density of the clumps that could potentially become gravitationally bound. The parameter of interest is the Jean’s Mass, approximately setting the limit between a clump that can support itself by its internal gas pressure and one that cannot, the latter being subject to a runaway gravitational collapse.  This limiting mass seems not to change much throughout our Galaxy, except its inner region: within the inner few hundred parsecs of our Galaxy, the so-called Central Molecular Zone, the situation is quite different!

This Central Molecular Zone is a mix of hot and cold dust, molecules and atoms that are not at all in a quiescent state. The massive black hole at the center of our Galaxy seems to have produced a barred galactic potential, structuring the surrounding material in two different sets of orbits, X1 and X2, with matter moving from the outer orbit into the inner one through some interaction areas. The X1 orbits could have a twisted structure at pericenter, compressing the gas and moving it into the inner orbits. Due to the mass densities observed and the frequency of compressional events expected in the Central Molecular Zone, one would expect this region to be a perfect nursery of stars in the Galactic center. Indeed, evidence that this has been the case in the past is the presence of three of the most massive stellar cluster in the Galaxy. Due to their mass, the stars in these clusters should be young (with a lifetime of few million years).

One of these massive stellar clusters, the so-called Central cluster, is located very close to the supermassive black hole. This cluster, together with other massive stars that orbit the central engine of our Galaxy, dissociate and ionize the molecular and atomic material in the surrounding region, creating at the outer edge of it a very dense molecular structure called the Circumnuclear Disk. The Circumnuclear Disk is not a real disk, but consists of streamers of material that are rotating around the center of the Galaxy and that are probably compressed and ionized as they enter the inner 1 parsec, forming a structure called the mini-Spiral.

 Slide2

Figure 1:  Circumnuclear disk of the Galaxy and mini-Spiral observed by the Submillimeter Array interferometer. CN emission is in green, showing the densest region in the gas streamers. The ionized material of the mini-spiral is in orange. Red and blue correspond to the shock-tracers SiO and H2CO, respectively. (Image credit: Martin et al. 2012).

The other two massive clusters, called the Arches and the Quintuplet, are located further away from the center of the Galaxy, at positive Galactic longitudes. They lie in a very interesting area, surrounded by a lot of dense material and by what looks like magnetic tubes of plasma with bright centimeter continuum emission.  These clusters themselves produce very intense ionization fields that carve-in the molecular material that surrounds them.

In the last few years, the Herschel Space Observatory has been able to map the relatively warm gas in the plane of the Galaxy.  These observations, using different molecular tracers, have clearly shown that in the Central Molecular Zone there is a ring of dense material around the central black hole with a radius of about 150 pc.  There is still a debate regarding how big this ring-like structure really is and whether it closes-in. Its interpretation in terms of individual clouds is problematic because of the range of velocities involved. The Galactic center is a very crowded area, with material spreading in velocity from -200 km/s to 200 km/s. Most of the material at negative longitude shows negative velocity whereas material at positive longitude shows positive velocity. However, in any given region, it is possible to identify more than one component, with velocities differing by more than 50 km/s. This, together with the spread in velocity due to the presence of turbulence, complicates the identification of individual clouds.

check1

Figure 2: Twisted disk of the Galaxy observed by the Hi-Gal Hershel survey at 250 microns. This ring-like features covers most of the Central Molecular Zone (Image credit: Molinari et al. 2011).

The only region in this ring-like structure that currently seems to be forming stars efficiently is the Sgr B2 region. With three very dense cores (S, M and N), this region shines strongly at radio and sub-millimeter wavelengths, lying at one of the edges of the large twisted ring. Each of these cores, currently in a molecular hot-core phase, will eventually form a small cluster of stars. Their current evolutionary phase is fascinating due to their chemical richness and in this regard Sgr B2 is a case of study, with new molecules being discovered there every year (the latest being isopropyl cyanide, C3H7CN, and methyl isocyanate, CH3NCO, discovered in 2014 and 2015, respectively – for more info on this fascinating new discoveries check Astrochymist at http://www.astrochymist.org/astrochymist_ism.html).

The rest of the ring-like structure seems more quiet, although there is a very dense molecular cloud, called The Brick, that has recently raised a lot of interest, having been the target of most of the radio and sub-mm facilities in the world (ALMA, SMA, VLA, APEX, Effelsberg, ATCA). These observations have shown that this region could be dense enough to form the next generation of stars. Indeed, the presence of Maser emission could be interpreted as a sign of on-going star formation, but it could also be produced by strong shocks commonly found in the Galactic center.  Regardless whether or not star formation has already started, this is one of the more prominent regions that will likely form stars in the near future.

Finding star formation in the Circumnuclear Disk of our Galaxy is not an easy task. The streamers that fall into the inner ionized region were thought to harbor some pockets of star formation, however, when we observed them a few years ago using CO and the dense gas tracers HCN and HCO+, we concluded that these objects were not gravitationally bound. The material there seems to be heavily disrupted by the sheer forces that arise due to their close proximity to the central black hole.  The mass of one of the regions was close to be gravitationally bound, however, this region also showed unexpected vibrational excited emission of HCN and, when accounting for its presence in the analysis, its estimated mass density decreased deeming on-going star formation in this region of the Circumnuclear Disk unlikely. More recent observations of this unique region have unveiled the presence of the shock-tracer molecule SiO. A possible explanation is that this is the location where two of the clouds in the vicinity of the high velocity streamers collide. Our new ALMA observations of CO, HCN, HCO+, CN, and tens of other molecules in the Circumnuclear Disk will soon shed some light on this fascinating structure near the heart of our Galaxy.

Another fascinating region at the Galactic center is a bubble-like structure seen in continuum and molecular emission, likely produced by a cataclysmic event.  This elongated bubble is orthogonal to the twisted ring observed by Herschel. In the edges of this bubble, the material has been compressed and is possible to observe clumps in many different molecular tracers. The expected high densities of these clumps deem them as promising sites for on-going star formation. But this still needs to be confirmed as it is not clear that their densities are high enough to keep them gravitationally bound in the extreme physical environment that surrounds the monster black hole at the center of our Galaxy.

Slide3

Figure 3: Composite image of the Central Molecular Zone showing the elongated bubble, outlined by the blue ellipses. Chandra (x-ray) observations are in blue, Hubble (near-IR) in yellow and Spitzer (mid-IR) in red. (Image credit: NASA).

 

References:

Belloche et al. 2014, Sci, 345, 1584

Binney et al. 1991, MNRAS, 252, 210

Halfen et al. 2015, ApJ, 812, 5

Requena-Torres et al. 2012, A&A, 542, 21

Martin et al. 2012, A&A, 539, 29

Molinari et al. 2011, ApJ, 209, 337

Mills et al. 2013, ApJ 779, 47

Mills et al. 2015, ApJ, 805, 72

 

Reviews:

Morris & Serabyn, 1996, ARA&A, 34,645

Genzel et al. 2010, RevModPhys, 82, 2121

 

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)

Mar 172014
 

Using Hubble to probe the dynamic interaction between the Magellanic Clouds

By Andrew Fox, Astronomer at STScI

Interactions between spiral galaxies and their dwarf satellites are often spectacular, producing extended streams of stripped gas and triggering new generations of star formation. The most striking local example of such an interaction lies in the outer halo of the Milky Way in the form of the Magellanic Stream. Extending for over 140 degrees across the Southern Sky, the Stream is a giant ribbon of gas trailing the orbit of the Large and Small Magellanic Clouds as they journey around the Galaxy. Since its discovery over 40 years ago, the Stream has puzzled observers and theorists alike and raised many questions. How was it physically removed from the Magellanic Clouds? Did it originate in the LMC or SMC? And what will its ultimate fate be? New spectroscopic observations with the Hubble Space Telescope and the Very Large Telescope are addressing these questions and finding the origin of the Stream to be surprisingly complex.

Magellanic_stream

Figure 1 :  Top: In this combined all-sky radio and visible-light image, the Magellanic Stream is shown in pink. The radio observations, taken from the Leiden-Argentine-Bonn (LAB) Survey, have been combined with a visible-light panorama. The Milky Way is the light blue band in the center of the image. The brown clumps are interstellar dust clouds in our galaxy. The Magellanic Clouds are seen in white at bottom right. Bottom: close-up of the Stream with our HST/COS sightlines marked with crosses. Credit: NASA, ESA, D. Nidever et al., NRAO/AUI/NSF, A. Mellinger, LAB Survey.

Measuring the chemical abundance (metallicity) of interstellar gas clouds requires finding UV-bright background sources, such as quasars. By splitting the quasar light into its constituent colors, the absorption lines imprinted by foreground gas clouds can be measured. These lines encode detailed information on the chemical composition and motion of the foreground clouds. Using observations from the Cosmic Origins Spectrograph (COS) installed on Hubble in 2009, we observed eight active galactic nuclei (AGN) lying behind or near the Stream. By comparing the strength of the neutral oxygen (O I) and ionized sulfur (S II) UV absorption lines to the strength of the atomic hydrogen (H I) 21 cm emission measured by radio telescopes, we derived the Stream’s metallicity in each direction. O I and S II were chosen for these measurements since they are largely unaffected by ionization and dust-depletion effects, so their ratios with H I provide robust metallicity indicators.

We found the Stream’s metallicity to be only ≈10% of the solar value in three separate directions sampling most of its length, considerably lower than the current-day average metallicity of the SMC (≈20% solar) and the LMC (≈50% solar). However, the age of the Stream is estimated from tidal models to be around 2 billion years, and information on the metallicity evolution of the Magellanic Clouds indicates that 2 billion years ago, the SMC abundance was ~10% solar, matching the value we measure in the Stream, whereas the LMC abundance was much higher, at ~30-40% solar. Our results thus support a scenario in which most of the Stream was stripped from the SMC (not the LMC). It has not self-enriched since its formation, because there is no evidence for ongoing star formation in the gas. In a sense, we have measured a fossil record of the Stream at the time of its birth in the SMC about 2 billion years ago.

However, a fourth sightline we studied (toward the AGN Fairall 9) tells a very different story. In this direction, which lies close to the Magellanic Clouds on the sky, the sulfur abundance in the Stream is found to be 50% solar, five times higher than the value measured in the other directions, and much higher than expected for gas that has been stripped from the SMC. Furthermore, the Fairall 9 direction intercepts a filament of the Stream that appears to connect kinematically to the LMC. Our measurement of a higher metal abundance supports this claim, and points toward a dual origin for the Stream, with two interwoven strands of material, one pulled out of the SMC about 2 billion years ago, and another pulled out of the LMC more recently.

Ongoing work by our team is investigating the total mass and inflow rate of the Magellanic gas onto the Milky Way, where it will potentially be able to fuel future generations of star formation. However, the gas must first survive the perilous journey through the hot Galactic corona, which can evaporate passing gas clouds. This survivability is being tested with computer simulations.

For more details, see Fox et al. 2013 (ApJ, 772, 110) and Richter et al. 2013 (ApJ, 772, 111), and the NASA press release and Huffington post article