A Black Hole Lurks in the Trapezium

The Orion Nebula Cluster (ONC) is a young star cluster whose age is estimated to be less than ~3 million years. Due to its proximity, the ONC is one of the best studied star clusters. At the heart of the ONC is the Trapezium, a tight cluster of several massive OB stars. A study by Subr et al. (2012) suggests that a massive black hole with more than 100 times the mass of the Sun might be lurking in the Trapezium. This is due to the large velocity dispersion observed for the 4 brightest Trapezium stars - Θ1A, Θ1B, Θ1C and Θ1D.

Considering the small number of stars in the sample, such a velocity dispersion measurement is not particularly robust. Nevertheless, the large velocity dispersion indicates there is more mass holding the cluster together than can be accounted for by just the stars. As a result, the presence of a black hole with more than 100 times the Sun’s mass is hypothesized. A black hole of this mass would represent the collapsed remnant of what was once a massive runaway-mass star. At 1300 light years away, this black hole would be the closest known to Earth.

This image shows the heart of the ONC. Clearly visible are the 4 Trapezium stars. These 4 hot and massive stars dominate the core of the ONC. Image credit: Hubble Legacy Archive, Robert Gendler

The ONC was once more compact that it currently is. For instance, the ONC does not contain wide stellar binaries with separations larger than ~1000 AU which would have been disrupted when the cluster was more compact. A consequence of being more compact is that gravitational interactions among massive OB stars would have been frequent enough to cause a fraction of them to be ejected from the cluster and a fraction of them to undergo ‘runaway’ stellar collisions to form a massive runaway-mass star. Subsequently, the massive runaway-mass star collapses directly to form a black hole, which can still continue to experience runaway growth by accreting more mass.

Interestingly, the observed number of massive OB stars in the ONC appears to be fewer than predicted. This is consistant with the loss of massive OB stars through ejection from the cluster and ‘runaway’ stellar collisions, leading to a deficit in massive OB stars. If a massive black hole is indeed lurking in the Trapezium, it probably isn’t growing much at present. The intense radiation from the hottest and most massive stars in the ONC would have driven much of the star-forming gas out of the cluster, causing the cluster to expand in size as there is less mass keeping the cluster together. As the cluster swells, stars collide less frequently and the runaway growth of the black hole slows to a halt.

Reference:

Subr et al. (2012), “Catch me if you can: is there a runaway-mass black hole in the Orion Nebula Cluster?”, arXiv:1209.2114 [astro-ph.GA]

Tantalizing Possibility of a Hidden Ocean on Mimas

Mimas is the smallest and innermost member of Saturn’s mid-sized icy moons, a group which includes Enceladus, Tethys, Dione, Rhea and Iapetus. With a diameter of 396 km, Mimas is also one of the smallest objects known in the Solar System whose self-gravitation is sufficiently strong to keep it rounded in shape. Mimas whizzes around Saturn about once every 23 hours. As it goes around Saturn, Mimas wobbles back and forth. This type of motion is called libration.

By carefully analysing images of Mimas taken by NASA’s Cassini spacecraft using a technique known as stereophotogrammetry, one particular component of Mimas’ libration was found to have an amplitude roughly twice as large as predicted. Since this component of libration depends on the interior structure of Mimas, its large value suggests that Mimas has a ‘weird’ interior. A few interior models of Mimas have been proposed to explain the large libration.

An image of Saturn’s moon Mimas taken by the Cassini spacecraft on 13 February 2010. Credit: NASA/JPL/Space Science Institute.

The 140 km wide Herschel impact crater on Mimas makes it resemble the Death Star from the Star Wars franchise. One interior model suggests the presence of a large mass buried beneath the Herschel impact crater, making Mimas more massive on one side. However, this model is inconsistent because the presence of such a large buried mass would have permanently reoriented Mimas such that the Herschel impact crater would face more towards Saturn, which is not the case.

Mimas is basically comprised of a shell of icy material overlying a denser rocky core. A more plausible interior model to explain the large libration involves Mimas having an elongated rocky core. However, such an elongated rocky core is expected to have an effect on the global shape of Mimas. If the icy shell is fully relaxed over the elongated rocky core, then the overall shape of Mimas should appear more elongated, which is not the case. Nevertheless, a low gravity object like Mimas can maintain large internal porosities which can create space for an oddly shaped core without affecting its overall shape.

A more exciting interior model of Mimas suggests that this small icy moon of Saturn might have an internal global ocean of liquid water located 24 to 31 km beneath its battered icy surface. For a small object like Mimas, it is difficult to keep an internal ocean from freezing. Heat generated from the decay of radioactive isotopes in the rocky core of Mimas would easily escape through the icy shell and cause the internal ocean to quickly freeze.

However, such an internal ocean on Mimas can still be kept liquid through heat generated from tidal heating. This is because Mimas’ orbit around Saturn is somewhat eccentric and the eccentricity may even have once been higher. As a result of its eccentric orbit, Mimas is sometimes closer to Saturn and sometimes further away. This causes Mimas to feel a difference in the gravitational pull from Saturn, which has the effect of alternately squeezing and stretching Mimas. Such a flexing motion creates friction in the interior of Mimas and friction generates heat. A tidally heated internal ocean of liquid water on Mimas is not inconceivable. Its neighbour, Enceladus, is known to have an internal ocean of liquid water sustained by tidal heating.

Reference:

Tajeddine et al., “Constraints on Mimas’ interior from Cassini ISS libration measurements”, Science 17 October 2014: Vol. 346 no. 6207 pp. 322-324

Polar Ice Deposits on Mercury

NASA’s MESSENGER spacecraft has obtained the first ever optical images showing the presence of water-ice and other frozen volatiles within the permanently shadowed interiors of craters near Mercury’s north pole (Chabot et al., 2014). It may come as a surprise that water-ice is present on Mercury since Mercury is the closest planet to the Sun and surface temperatures at its equatorial regions can soar above 400°C. However, near Mercury’s poles, there are a number of craters whose interiors are permanently shadowed from the Sun. Since Mercury does not have an atmosphere to transport heat around the planet, the permanently shadowed interiors of these craters serve as cold traps where water-ice and other volatiles can remain frozen there.

Figure 1: Mercury.

 Figure 2: Locations of water-ice deposits in the shadowed interiors of craters on Mercury.

Over two decades ago, Earth-based radar observations provided the first indications that water-ice might be present on Mercury’s poles. MESSENGER entered orbit around Mercury on 18 March 2011 and has been observing the planet from orbit ever since. In late 2012, the presence of polar water-ice deposits on Mercury was confirmed by MESSENGER through a combination of observations involving neutron spectrometry (Lawrence et al., 2013), measurements of surface reflectance at the near-infrared wavelength of 1064 nm (Neumann et al., 2013) and thermal modelling (Paige et al., 2013).

The polar deposits of water-ice and other frozen volatiles on Mercury were imaged using the Wide-Angle Camera (WAC) on MESSENGER’s Mercury Dual Imaging System (MDIS). Although the polar deposits never receive direct sunlight, they could still be imaged by taking advantage of the very low levels of sunlight scattered off illuminated crater walls. Images of the permanently shadowed interior of the 112 km wide Prokofiev crater, the largest crater near Mercury’s north pole, show an area with widespread surface water-ice deposits. The area shows up in the WAC images as a region with higher reflectance compared to its surroundings.

Numerous smaller craters cover the floor of Prokofiev crater. The area with surface water-ice deposits within Prokofiev crater has a similar cratered terrain as the neighbouring sunlit surface. This indicates that the water-ice deposits were placed there after the formation of the underlying craters, suggesting that the water-ice deposits were placed there relatively recently instead of billion of years ago. Furthermore, the water-ice deposits appear to be uniform, again implying a recent emplacement. Because if the water-ice deposits were there before impacts excavated the craters, a patchy appearance would result since the craters and their ejecta would have buried parts of the water-ice deposits.

WAC images of other craters with permanently shadowed interiors show areas of lower reflectance believed to be water-ice deposits covered by a thin, overlying layer of dark, organic-rich volatile material. These lower reflectance deposits extend to the edges of the permanently shadowed regions and terminate sharply. The sharp boundaries indicate that the deposits are relatively young since the long process of lateral mixing by impacts has yet to smudge the boundaries. WAC images of surface volatile deposits in Mercury’s polar craters show that these deposits are relatively young. The deposits were either delivered to the planet recently or continuously restored at the surface through an ongoing process.

References:

- Chabot et al., “Images of surface volatiles in Mercury’s polar craters acquired by the MESSENGER spacecraft”, Geology 15 October 2014, v. 42, no. 10

- Lawrence et al., “Evidence for Water Ice Near Mercury’s North Pole from MESSENGER Neutron Spectrometer Measurements”, Science 18 January 2013: Vol. 339 no. 6117 pp. 292-296 DOI: 10.1126/science.1229953

- Neumann et al., “Bright and Dark Polar Deposits on Mercury: Evidence for Surface Volatiles”, Science 18 January 2013: Vol. 339 no. 6117 pp. 296-300 DOI: 10.1126/science.1229764

- Paige et al., “Thermal Stability of Volatiles in the North Polar Region of Mercury”, Science 18 January 2013: Vol. 339 no. 6117 pp. 300-303 DOI: 10.1126/science.1231106

Brown Dwarf in a Distant Orbit around an A-Type Star

Brown dwarfs are substellar objects that populate the gap between the most massive planets and the least massive stars. ζ Del B is a newly discovered brown dwarf around the A-type star ζ Del A. Before this discovery, ζ Del A was simply known as ζ Del without the suffix “A”. ζ Del A is a main-sequence star with 2.5 ± 0.2 times the mass of the Sun and about 50 times the Sun’s luminosity. A main-sequence star is basically a star that is currently generating energy by fusing hydrogen into helium and it is neither at the start nor near the end of its life. The Sun is one example of a main-sequence star.

Figure 1: Artist’s impression of a brown dwarf.

Only a small number of brown dwarfs are known to orbit stars that are significantly more mass and luminous than the Sun. ζ Del B is the latest addition to this short list of brown dwarfs. ζ Del A and its brown dwarf companion, ζ Del B, are both estimated to lie at a distance of roughly 220 light years away. Estimates show that the ζ Del system is 525 ± 125 million years old. Spectroscopic observations of the spectrum of ζ Del B indicate it is an L-type brown dwarf (L5 ± 2) with an effective temperature of 1650 ± 200 K. Brown dwarfs are classified into 4 spectral types - M, L, T and Y. M-type brown dwarfs are the hottest, while Y-type brown dwarfs are the coolest.

Based on its near-infrared brightness, temperature and age, ζ Del B is estimated to be 50 ± 15 times the mass of Jupiter. This gives the brown dwarf a mass ratio of 0.019 ± 0.006 with respect to its stellar companion. Additionally, ζ Del B has a projected separation of 910 ± 14 AU from ζ Del A. Such a projected separation means the orbital period of ζ Del B around ζ Del A is probably on the order of ~10,000 years. ζ Del B is one of the most widely-separated and lowest mass ratio substellar companions known around a main-sequence star.

Figure 2: Estimated mass of ζ Del B based on its near-infrared brightness, temperature and age. De Rosa et al. (2014).

 Figure 3: Mass ratio as a function of separation for brown dwarf companions (blue open circles), directly imaged brown dwarf and planetary companions (red open squares), and brown dwarf and planetary companions detected by radial velocity and transit techniques (black points). ζ Del B (black filled star) is among the most widely separated, lowest mass ratio companions imaged to date. De Rosa et al. (2014).

A number of different scenarios might explain the formation of ζ Del B and other widely-separated substellar companions around main-sequence stars. The formation of ζ Del B at its current location would require an unusually massive circumstellar disk at a large distance around ζ Del A and such a formation scenario is quite unlikely. Nevertheless, it cannot be ruled out that ζ Del B might have formed much closer to ζ Del A before migrating outward due to interactions with the circumstellar disk or with an unseen companion.

The ζ Del system could also have formed from the fragmentation of a single pre-stellar core of gas and dust into two cores, where one core is much more massive than the other. The more massive core collapsed to form ζ Del A, while the less massive core collapsed to form ζ Del B. Such a scenario can produce companions with separations on the order of ~1,000 AU. Another formation scenario, albeit with a low probability of occurring, involves ζ Del B forming independently before being captured by ζ Del A to form a low mass ratio binary.

Reference:

De Rosa et al. (2014), “The VAST Survey - IV. A wide brown dwarf companion to the A3V star ζ Delphini”, arXiv:1410.0005 [astro-ph.SR]

Adaptation of Antarctic Lichens to Conditions on Mars

Billions of years ago, Mars was a warm and wet planet. Life could have evolved on Mars and then receded to micro-habitats as the planet subsequently became colder and dryer. Present-day micro-habitats for life on Mars can include subterranean aquifers and cracks or fissures in rocks. J.-P. de Vera et al. (2014) conducted a study using the lichen Pleopsidium chlorophanum and found that this Antarctic lichen can adapt, within a span of 34 days, to the conditions expected to be present in the micro-habitats on Mars today.

The sample used in the study was collected from the granites and volcanic rocks of North Victoria Land in Antarctica during the 10th German North Victoria Land Expedition in 2009/2010. Pleopsidium chlorophanum is an extremophile that lives in very cold and dry places. Its native habitat in Antarctica somewhat approximates the conditions on Mars. Pleopsidium chlorophanum is usually found within cracks and fissures in rocks. It can remain metabolically active down to -20°C and can absorb water directly from snow.

To simulate Mars-like environmental conditions, the lichens were placed in the Mars Simulation Chamber (MSC) at the Mars Simulation Facility (MSF) of the DLR Institute of Planetary Research in Berlin. The atmosphere in the MSC was 95 percent carbon dioxide, 4 percent nitrogen and 1 percent oxygen. The pressure was held at 800 Pa, with a diurnal relative humidity cycling of 0.1 to 75 percent and a diurnal temperature cycling of 21°C to -50°C (i.e. similar to the temperatures observed in the equatorial to mid-latitude regions on Mars).

The lichens were embedded within a Mars analogue soil material. In the experiment, 3 samples of lichens were subjected to Mars-like niche conditions (i.e. conditions expected in the micro-habitats on Mars) and another 3 samples of lichens were subjected to the unprotected Mars-like surface conditions for 34 days. For lichens in the unprotected Mars-like surface conditions, they were subjected to much more intense UV irradiation.

Results from the experiment indicated that for Mars-like surface conditions, photosynthetic activity dropped to 18 percent of pre-experiment levels and it was unclear if the lichens remained photosynthetically active at the end of the 34 days. However, lichens that were subjected to Mars-like niche conditions and experienced a much lower radiation dose fared very differently. For these lichens, photosynthetic activity only dropped to 55 percent of pre-experiment levels after the 34 days. In fact, photosynthetic activity at the end of the experiment was 17 percent higher than what was measured for the lichens in their native habitat in Antarctica.

Under simulated Mars-like niche conditions, the lichens appeared to have experienced an initial period of shock lasting ~7 days. Following that, the lichens rapidly adapted. Photosynthetic activity increased over the subsequent days and the increase continued to the end of the experiment. It appears that the lichens were much more sensitive to the intensity of UV irradiation than to other Mars-like parameters such as high carbon dioxide concentration, very low temperatures, extreme humidity fluctuations and very low atmospheric pressure. This study supports the notion that Earthly life can adapt to the present-day conditions on Mars. Furthermore, life which may have originated during the early warm and wet Mars might still survive and thrive in present-day micro-habitats on Mars.

Reference:

J.-P. de Vera et al., “Adaptation of an Antarctic lichen to Martian niche conditions can occur within 34 days”, Planetary and Space Science 98 (2014) 182-190

Scattering of Super-Earths

In the core-accretion model, the formation of gas giant planets begins with the rapid coalescence of solids to form massive cores. When a massive core reaches ~10 times the mass of Earth, it accretes gas in a runaway fashion and eventually becomes a full-fledged gas giant planet. During the formation of gas giant planets around a young star, multiple massive cores can form. These massive cores can be gravitationally scattered out to large distances by other massive cores and by newly-formed gas giant planets. Typically, over 80 percent of massive cores between 1 to 15 times the mass of Earth get scattered out to the peripheries of their natal planetary systems.

The scattering process hurls the massive core, also known as a scattered planet, onto an eccentric orbit which takes the planet out to a large distance from its host star. Subsequently, the planet can interact with the gaseous disk around its nascent host star. Planet-disk interactions can damp the planet’s orbital eccentricity, causing the planet to settle into a more circular orbit in the outer regions of its planetary system. However, the outcome of such a planet-disk interaction depends on the planet’s mass and both the characteristics and evolution of the gaseous disk.

A scattered planet needs to be at least as massive as the Earth in order for its orbital eccentricity to be damped. In fact, orbital circularization is most effective for massive planets as they interact most strongly with the gaseous disk. The orbits of scattered Earth-mass or smaller planets tend to remain eccentric. Gaseous disks that are more massive and longer-lived can damp eccentricities more effectively than those which are less massive and short-lived. Additionally, gaseous disks that decay with an expanding inner cavity can circularise orbits at larger distances than those which dissipate homogeneously.

Typically, a gaseous disk around a young star has a lifespan of a few million years before it decays completely. Gaseous disks that dissipate homogeneously are only effective in circularising the orbits of more massive super-Earths and Neptune-mass planets. The orbits of these planets are circularised at smaller orbital distances (i.e. within 50 AU for Sun-like stars). For smaller planets, planet-disk interactions are less effective and these planets tend to remain in moderately eccentric orbits that continue to take them out to larger orbital distances (i.e. ~200 AU for Sun-like stars). Gaseous disks that decay with an expanding inner cavity can circularise the orbits of super-Earths at orbital distances beyond 100 AU. Smaller planets can remain at larger orbital distance, although their orbital eccentricities cannot be damped completely.

The orbital parameters of a few distant trans-Neptunian objects such as Sedna and 2012 VP113 suggests it is plausible a super-Earth with 2 to 10 times the mass of Earth is lurking in a low-eccentricity orbit between 200 to 300 AU from the Sun. If this planet exists, it is unlikely to have formed where it currently is and also unlikely to have migrated from inside 30 AU to its present orbit. However, the existence of such a planet in the far outer reaches of the Solar System is conceivable if it was scattered from inside 30 AU and subsequently interacted with a massive gaseous disk around the young Sun. Such a gaseous disk would have to dissipate from the inside out in order for planet-disk interactions to “circularise” the planet’s orbit to where the planet is currently predicted to reside.

Reference:

Bromley & Kenyon (2014), “The Fate of Scattered Planets”, arXiv:1410.2816 [astro-ph.EP]

Tidally Distorted Rocky Exoplanets

Exoplanets in close-in orbits around their host stars can become tidally distorted. Being in a close-in orbit also makes it likely that the planet is tidally-locked with the same planetary hemisphere perpetually facing its host star. For a tidally-locked exoplanet, the effect of tidal distortion tends to stretch the planet into a triaxial ellipsoid where the planet’s longest axis is always oriented towards its host star. If the planet transits its host star, the effect of tidal distortion and the resulting asphericity in the planet’s shape can cause the planet’s size to be underestimated, and subsequently, if the mass of the planet is measured, the planet’s density to be overestimated.

Figure 1: Artist’s impression of a planet transiting a star.

The effect of tidal distortion on close-in, tidally-locked exoplanets has been explored only for gas giant planets due to the general assumption that tidal distortion is only relevant for such planets. It is commonly assumed that the effect of tidal distortion is too small for rocky exoplanets. A study done by Saxena et al. (2014) show this is not the case. Rocky-exoplanets in close-in orbits around smaller red dwarf stars can become tidally distorted to an observable extend. In particular, the study focuses on rocky exoplanets with 1, 1.5 and 2 Earth radii orbiting M5V and M1V red dwarf stars. An M5V red dwarf star is smaller and less massive than an M1V red dwarf star.

Results from the study show that for a 1.5 or 2 Earth radii rocky exoplanet circling an M5V red dwarf star near the fluid Roche limit, the effect of tidal distortion can stretch the planet sufficiently such that the ratio of the planet’s longest to smallest axis can approach 3:2. Basically, the fluid Roche limit is the minimum distance a planet can be from its host star before the planet’s own gravity can no longer hold itself together and the planet starts to disintegrate. For rocky exoplanets around M1V red dwarf stars, the effect of tidal distortion is less significant such that the longest axis is at most only ~10 percent larger than the shortest axis.

Figure 2: The ratio of the longest to smallest axis for tidally distorted rock exoplanets of 1, 1.5 and 2 Earth radii orbiting M5V and M1V red dwarf stars. Saxena et al. (2014)

Tidal distortion decreases the projected area of a planet when viewed along the planet’s longest axis. As a result, when a planet transits its host star, the effect of tidal distortion causes the transit depth (i.e. fraction of starlight the planet blocks) to be shallower. This can lead to the planet’s size being underestimated. For rocky exoplanets with 1 to 2.25 Earth radii around M5V red dwarf stars, underestimates can reach ~10 percent at near the fluid Roche limit. For M1V red dwarf stars, underestimates only reach 2.5 to 5.5 percent. The size underestimates can lead to density overestimates for these planets.

Figure 3: Radius underestimates for different sized rocky exoplanets due to tidal and rotational distortions. Saxena et al. (2014)

The planetary models used in this study assume an Earth-like composition. Nonetheless, a more rigid planet (e.g. a solid, pure-iron planet) would be less susceptible to tidal distortion than a less rigid planet. Since tidal distortion is sensitive to a planet’s rigidity, observations of a planet’s asphericity due to tidal distortion might provide insights to the planet’s interior structure and bulk composition.

Red dwarf stars tend to have very compact planetary systems and these stars also tend to have smaller planets. As a result, it is more likely for red dwarf stars to harbour rocky exoplanets on close-in orbits that can cause these planets to be subjected to strong tidal distortion. Furthermore, red dwarf stars are the smallest and least massive stars. This means that a planet around a red dwarf star can induce a proportionally larger observational signature compared to a same planet around a Sun-like star.

Reference:

Saxena et al. (2014), “The Observational Effects and Signatures of Tidally Distorted Solid Exoplanets”, arXiv:1410.2251 [astro-ph.EP]

An Ancient Relic of the Universe’s First Galaxies

Ultra-faint dwarf galaxies (UFDs) are the faintest galaxies known in the universe. These galaxies are smaller than 1,000 light years in radius and have very low metallicities. As a result, UFDs are believed to represent the first generation of galaxies in the universe. UFDs are the most dark matter dominated galaxies since nearly all their mass is admittedly in the form of dark matter. In fact, the total mass of all stars in a UFD is typically less than a million solar masses. For comparison, the Milky Way galaxy contains 200 to 400 billion stars. UFDs are also much smaller and fainter than classical dwarf spheroidal galaxies. So far, UFDs have only been discovered around the Milky Way galaxy and the neighbouring Andromeda galaxy.

Using images acquired by the Hubble Space Telescope, Jang & Lee (2014) reported the discovery of a UFD in a region of space far from any massive galaxy in the Virgo Cluster. This newly discovered UFD is named Virgo UFD1. Observations indicate that Virgo UFD1 contains low metallicity red-giant-branch (RGB) stars but no asymptotic-giant-branch (AGB) stars. It means that Virgo UFD1 is very old, most likely older than 10 billion years, because AGB stars mark the final evolution of stars that are at least as massive as the Sun and such stars do not live more than about 10 billion years.

The stars in Virgo UFD1 probably formed in a single burst of star formation more than 10 billion years ago. As a consequence, the absence of subsequent generations of stars to fuse hydrogen and helium into heavier elements results in the low metallicity of Virgo UFD1. There is also no sign that Virgo UFD1 experienced tidal interaction with any galaxy. Estimates place Virgo UFD1 at a distance of more than 50 million light years away. The old age, low metallicity and large distance from any massive galaxy suggest Virgo UFD1 may represent an ancient relic of the universe’s first galaxies.

Left panel: Effective radius versus absolute total magnitude of Virgo UFD1 (large starlet symbol) in comparison with other stellar systems. Right panel: The central surface brightness versus absolute total magnitude of Virgo UFD1 in comparison with other stellar systems. Note: Circles and lenticular symbols for the giant ellipticals and bulges in spiral galaxies, downward triangles for the UCDs, pentagons for the Milky Way galaxy globular clusters, squares and diamonds for the Local Group satellite galaxies and UFDs, and upward triangles for the dwarf galaxies in M81 and M106 and the low surface brightness galaxies in M101, and small yellow starlet for Virgo dSph-D07. Jang & Lee (2014).

Reference:

Jang & Lee (2014), “Discovery of an Ultra-Faint Dwarf Galaxy in the Intracluster Field of the Virgo Center: A fossil of the First Galaxies”, arXiv:1410.2247 [astro-ph.GA]

A Warm Gas-Giant Planet Orbiting a Giant Star

A study by Ortiz et al. (2014) reported on the spectroscopic confirmation of a Jupiter-like gas-giant planet in a close-in, eccentric orbit around a giant star. This planet is identified as KOI-1299b, with the suffix “b” denoting its planetary nature. The host star of KOI-1299b is entering its later stages of stellar evolution and is about to swell into a red giant star. This star has 1.35 ± 0.10 times the Sun’s mass, 4.15 ± 0.12 times the Sun’s radius and an effective surface temperature of 5020 ± 60 K. KOI-1299b is observed by NASA’s Kepler space telescope to transit its host star once every 52.5 days. The size of KOI-1299b was estimated by measuring the drop in the star’s brightness each time the planet transits in front of it. A separate paper by Ciceri et al. (2014) on the discovery of KOI-1299b also reported similar results.

Figure 1: Artist’s impression of a gas-giant planet with a system of planetary rings circling it.

The planetary nature of KOI-1299b was confirmed using high-resolution spectroscopic follow-up observations conducted between June and October 2014 using the Calar Alto Fiber-fed Échelle spectrograph (CAFE) on the 2.2 m telescope of Calar Alto Observatory in Almería, Spain, and the Fibre-fed Échelle Spectrograph (FIES) on the 2.56 m Nordic Optical Telescope of Roque de los Muchachos Observatory in La Palma, Spain. These follow-up observations measured the amount of gravitational tugging KOI-1299b exerts on its host star, allowing the mass of KOI-1299b to be estimated and its planetary nature to be confirmed. Additionally, these follow-up observations also show that KOI-1299b has a highly eccentric orbit around its host star.

Figure 2: Radial velocity measurements of the host star of KOI-1299b. Upper panel: CAFE (blue circles) and FIES (red squares). Lower panel: Residuals. Ortiz et al. (2014)

 Figure 3: Left panel: Eccentricity and semimajor axis of the extrasolar planets discovered around main sequence stars (black dots) and giant stars (magenta circles). Right panel: Orbital period versus stellar mass. The position of KOI-1299b is marked with a green triangle in both panels. Ortiz et al. (2014)

KOI-1299b has 5.86 ± 0.05 times the mass of Jupiter, 1.08 ± 0.03 times the radius of Jupiter and an orbital eccentricity of 0.479 ± 0.004. With the mass and size known, the estimated density of KOI-1299b is 5.7 ± 0.5 g/cm³. The highly eccentric orbit of KOI-1299b brings the planet as close as ~0.16 AU (periastron) from its host star and out as far as ~0.45 AU (apastron). Between periastron and apastron, KOI-1299b receives ~450 to ~56 times as much insolation as Earth receives from the Sun. On average, the estimated equilibrium temperature of KOI-1299b is 942 ± 20 K. However, the planet’s highly eccentric orbit can cause temperatures to vary by ~500 K. KOI-1299b is classed as a warm-Jupiter since is not as hot as typical hot-Jupiters whose temperatures are well over 1000 K. As the host star of KOI-1299b swells into a red giant, tidal interactions between the star and planet will increase, eventually causing KOI-1299b to be engulfed by its host star.

References:

- Ortiz et al. (2014), “Spectroscopic confirmation of KOI-1299b: a massive warm Jupiter in a 52-day eccentric orbit transiting a giant star”, arXiv:1410.3000 [astro-ph.EP]

- Ciceri et al. (2014), “KOI-1299b: a massive planet in a highly eccentric orbit transiting a red giant”, arXiv:1410.2999 [astro-ph.EP]

Polluting a Red Supergiant Star with Heavy Elements

A recent study by Levesque et al. (2014) found that the red supergiant star HV 2112 in a nearby galaxy known as the Small Magellanic Cloud is enriched with various peculiar heavy elements. As a result, HV 2112 is postulated to be a Throne-Zytkow Object (TZO). Basically, a TZO is a red supergiant star with a neutron star at its center. The neutron star most likely got there by in-spiralling within the envelope of the red supergiant star, all the way down to the core. The neutron star destroys the red supergiant star’s core and part of the core forms an accretion disk around the neutron star. Temperatures and densities in the accretion disk are high enough to synthesize a range of peculiar heavy elements.

HV 2112 is observed to be enriched with heavy elements such as calcium, rubidium, lithium and molybdenum. Although most of these heavy elements can be synthesized by a red supergiant star on its own, the high calcium abundance requires HV 2112 to be a TZO since only in an accretion disk around a neutron star are the temperatures and densities sufficiently high to synthesize calcium. Nevertheless, a study by Sabach & Soker (2014) provides an alternative explanation whereby the high calcium abundance observed for HV 2112 came from a companion star that had exploded as a core collapse supernova (CCSN). During the supernova event, HV 2112 was already a large red supergiant star and could intercept more of the supernova’s ejecta, including the calcium that was synthesized in the supernova.

HV 2112 and its companion star started of as a pair of massive stars circling one another. The companion star is slightly more massive than HV 2112. Since more massive stars evolve quicker, the companion star first evolves to a red giant star, puffs up, and transfers some of its mass to HV 2112. Eventually, HV 2112 becomes more massive than its companion star and also evolves to become a red giant star. The companion star subsequently explodes as a CCSN. When that happens, HV 2112 has already evolved further to form a red supergiant star. Its large size allows HV 2112 to intercept a good fraction of the supernova ejecta that is rich in heavy elements, including calcium. As a result, HV 2112 could be a red supergiant star that was “polluted” by ejecta from a supernova, rather than a TZO.

References:
- Levesque et al. (2014), “Discovery of a Thorne-Zytkow object candidate in the Small Magellanic Cloud”, arXiv:1406.0001 [astro-ph.SR]
- Sabach & Soker (2014), “A super asymptotic giant branch star enriched with calcium by a supernova as the origin of HV2112, rather than a Thorne-Zytkow Object”, arXiv:1410.1713 [astro-ph.SR]

One Planet, Two Stars

Welsh et al. (2014) present the discovery of KIC 9832895b, a circumbinary planet in a 240.5 day orbit around an eclipsing binary. Basically, an eclipsing binary is a pair of stars that appear to eclipse one another as they orbit around each other. In the case of KIC 9832895b, it orbits around an eclipsing binary consisting of a pair of stars with 0.93 and 0.194 times the Sun’s mass. Both stars orbit around one another with a period of 27.3 days. The orbital period of KIC 9832895b is 8.8 times the orbital period of the eclipsing binary. This places the planet safely outside the dynamical instability zone.

Figure 1: Artist’s impression of a gaseous planet which KIC 9832895b might resemble.

Figure 2: Face-on view of KIC 9632895b’s orbit, showing the habitable zone (HZ). The dark green region corresponds to the narrow (conservative) HZ and the light green corresponds to the nominal (extended) HZ. The dashed red circle represents the dynamical instability zone. The orbit of KIC 9632895b is shown in white. Welsh et al. (2014).

KIC 9832895b was detected from its three transits across the primary star (i.e. the more mass star) of the eclipsing binary. The transit depth indicates that KIC 9832895b is 6.2 times the radius of Earth, indicating that KIC 9832895b is somewhat larger than Neptune. The mass of KIC 9832895b is estimated to be most likely less than 16 times the mass of Earth due to the absence of any noticeable perturbations it has on the eclipsing binary. This constrains the mean density of KIC 9832895b to be less than 0.38 g/cm³ and demonstrates that it is an unusually low density planet, probably of gaseous composition.

KIC 9832895b is the 10th circumbinary planet discovered using data collected from NASA’s Kepler space telescope. In addition, KIC 9832895b is also in the circumbinary habitable zone where temperatures are relatively clement. The time-averaged insolation that KIC 9832895b receives is estimated to be 94 percent the intensity of insolation Earth receives from the Sun. Although KIC 9832895b is itself unlikely to harbour life, it could host a large moon capable of supporting life. Of the 10 circumbinary planets known so far, KIC 9832895b is the third found to lie within the circumbinary habitable zone.

Interestingly, the inclination of KIC 9832895b oscillates with a 102.8 year period. As a result, transits only occur ~8 percent of the time. This explains why the three detected transits of KIC 9832895b were only found in the later portion of the Kepler dataset. The transits will not be observable after 2015 and will only return on 2066. Since the transits do not always happen, for every system like the one hosting KIC 9832895b, there are ~12 similar systems where planetary transits are not observed.

Reference:

Welsh et al. (2014), “KIC 9632895 - The 10th Kepler Transiting Circumbinary Planet”, arXiv:1409.1605 [astro-ph.EP]

Identifying Alien Planets with Clear Skies

When a planet transits in front of its host star, a tiny fraction of the starlight passes through the planet’s atmosphere and carries with it signatures of the planet’s atmospheric constituents. This can allow the planet’s atmosphere to be characterised using an observational technique known as transmission spectroscopy. However, the atmospheres of planets can be cloudy, hazy or clear-sky (i.e. free of clouds and hazes). The presence of clouds or hazes can obscure the lower layers of the atmosphere and make the planet less desirable for characterisation. As a result, it is worth identifying whether a planet has clear skies before a large amount of telescope time is dedicated to characterising its atmosphere.

Misra & Meadows (2014) propose a method to readily distinguish cloudy, hazy and clear-sky planets. This involves measuring the amount of starlight being refracted through the atmospheres of transiting planets using upcoming large collecting area space and ground-based telescopes such as the James Webb Space Telescope (JWST) and the European Extremely Large Telescope (E-ELT). The refraction of starlight by a planet’s atmosphere can lead to an increase of flux both prior to ingress (i.e. before the start of a transit) and subsequent to egress (i.e. after the end of a transit).

The presence of a global cloud or haze coverage tends to obscure layers of a planet’s atmosphere that refract light. As a result, the detection of refracted light pre-ingress and post-egress would strongly suggest the absence of a global cloud or haze layer, making the planet a promising candidate for follow-up observations to characterise its atmosphere. In the models, the atmospheric pressure cut-offs are at 1 mbar (hazy case), 0.1 bars (cloudy case) and 1 bar (clear-sky case). A higher pressure cut-off indicates a greater depth of measurable atmosphere.

Results from the study show detecting refracted light requires less than 10 hours of total observing time for Jupiter-sized planets with JWST and for Super-Earths/Mini-Neptunes with E-ELT. Since the increase in flux due to refraction prior to ingress and subsequent to egress can be readily detected for clear-sky planets, it can quickly identify whether a planet is a good candidate for extended follow-up observations. Characterising a planet’s atmosphere is a very time consuming process, making it important to select good candidates (i.e. clear-sky planets) prior to characterisation.

Reference:

Misra & Meadows (2014), “Discriminating Between Cloudy, Hazy and Clearsky Exoplanets Using Refraction”, arXiv:1409.7072 [astro-ph.EP]

An Oblate Giant Planet

Kepler-39b is a gas-giant planet in orbit around an F-type star. It is 18 times Jupiter’s mass, 1.22 times Jupiter’s radius and it transits its host star every 21.09 days. A study by Wei Zhu et al. (2014) using data from NASA’s Kepler space telescope found that Kepler-39b has an oblateness of 0.22 ± 0.11. In fact, this is the first tentative detection of oblateness for a planet outside the Solar System. When an oblate planet transits its host star, the transit light curve will exhibit small differences from that of a purely spherical planet.  

In the Solar System, the gas-giant planets Jupiter and Saturn are oblate in shape due to their rapid rotations. The oblateness of an object is expressed as the ratio of its equatorial-polar radius difference to its equatorial radius. The equatorial radius is larger than the polar radius by 7 percent for Jupiter and by 10 percent for Saturn. As such, Jupiter’s oblateness is 0.07 and Saturn’s oblateness is 0.1. With an oblateness of 0.22 ± 0.11, Kepler-39b is substantially more oblate than any planet in the Solar System.

The large oblateness of Kepler-39b is most likely rotationally induced. With that, its rotation period is estimated to be 1.6 ± 0.4 hours. For comparison, the rotation periods of Jupiter and Saturn are 9.9 and 10.6 hours, respectively. Although the rotation of Kepler-39b is remarkably fast, it is lower than its estimated break-up rotation period of ~0.9 hours. In addition to its large oblateness, Kepler-39b is also inflated in size. Its close proximity to its host star and its estimated equilibrium temperature of around 900 K is insufficient to account for its inflated size.

Reference:

Wei Zhu et al. (2014), “Constraining the Oblateness of Kepler Planets”, arXiv:1410.0361 [astro-ph.EP]

When Supermassive Stars Explode in the Early Universe

Supermassive stars with ~10,000 to ~100,000 times the Sun’s mass are believed to have formed in the very early universe. These are the first generation of stars in the universe and are entirely comprised of hydrogen and helium. They live very short lives before collapsing directly to form black holes. A team of astrophysicists ran a number of supercomputer simulations and found that some of these supermassive stars die in a rather unusual way. Instead of collapsing to form black holes, supermassive stars in a narrow mass range between 55,000 to 56,000 times the Sun’s mass explode as highly energetic thermonuclear supernovae, leaving nothing behind.

This image is a slice through the interior of a supermassive star of 55,500 solar masses along the axis of symmetry. It shows the inner helium core in which nuclear burning is converting helium to oxygen, powering various fluid instabilities (swirling lines). This snapshot shows a moment one day after the onset of the explosion, when the radius of the outer circle would be slightly larger than that of the orbit of the Earth around the Sun. (Credit: Ken Chen, UC Santa Cruz)

A supermassive star with 55,500 times the Sun’s mass lives for about 1.69 million years before it becomes unstable and starts to collapse. During its pre-collapse phase, the size of the star is slightly larger than the diameter of Earth’s orbit around the Sun and the star has an effective surface temperature of about 70,000 K. The star is also remarkably luminous, with ~1.5 billion times the Sun’s luminosity. With the onset of helium burning in the star’s core, the prodigious amount of thermal photons being generated in the core affects the star’s gravitational field by becoming an additional source of gravity. As a consequence, the core begins to contract, causing the temperature and density in the core to rise rapidly, accelerating nuclear burning.

As the core contracts, helium begins to burn explosively, fusing to carbon, and then to oxygen, neon, magnesium and silicon. The explosive nuclear burning occurs within a span of only several hours and releases ~10 times more energy than the binding energy of the star. This causes the star to halt its collapse and unbind completely in a massive explosion known as a general relativistic supernova (GSN). The amount of energy produced in such an event is ~10,000 times the energy released by a typical supernova. In total, about half the mass of the star is ejected in the form of elements heavier than hydrogen and helium. Mostly elements between carbon and silicon are produced, with only trance amounts of iron group elements.

After getting blasted out into the cosmos, these heavy elements are incorporated in the formation of subsequent generations of stars and planets. The energetic demise of these supermassive stars can be detected by upcoming space-based observatories such as ESA’s Euclid and NASA’s Wide-Field Infrared Survey Telescope (WFIRST). Additionally, indirect observational signatures of GSN explosions might be found by looking for early galaxies that are iron deficient but enhanced with elements from carbon to silicon.

Reference:

Ke-Jung Chen et al. (2014), “General Relativistic Instability Supernova of a Supermassive Population III Star”, arXiv:1402.4777 [astro-ph.HE]