A Population of Extremely Long Duration Gamma-Ray Bursts

Gamma-ray bursts (GRBs) are among the most energetic events in the universe. There are three classes of GRBs - short gamma-ray bursts (SGRBs) with durations less than ~2 seconds, long gamma-ray bursts (LGRBs) with durations extending to ~1,000 seconds and ultra-long gamma-ray bursts (ULGRBs) with durations of ~10,000 seconds. SGRBs originate from the mergers of neutron star binaries or neutron star-black hole binaries, while LGRBs are created when the cores of massive stars collapse to form black holes.

ULGRBs have such long burst durations that their progenitors are unlikely to be the same as those for LGRBs. Instead, ULGRBs probably originate from the collapse of giant or supergiant stars into black holes. These stars are orders of magnitude larger than the progenitors of LGRBs, resulting in much longer collapse times. Additionally, these stars have lower densities, resulting in lower mass in-fall rates. The continuous in-fall of material into the nascent black hole drives a GRB with an extremely long duration, leading to an ULGRB. Alternatively, ULGRBs can also be created when white dwarfs get tidally shredded by intermediate mass black holes (IMBH).

Figure 1: Artist’s impression of a gamma-ray burst.

Figure 2: Parameter space for GRBs and other high energy transient phenomena plotted as a function of burst duration versus average luminosity. The classes of events are - soft gamma repeaters (SGRs), short gamma-ray bursts (SGRBs), low-luminosity and long gamma-ray bursts (LLGRBs), long gamma-ray bursts (LGRBs), ultra-long gamma-ray bursts (ULGRBs) and tidal disruption events (TDEs). Andrew Levan (2015).

Recent observations by NASA’s Swift Gamma-Ray Burst Mission have revealed what could be a new population of GRBs with extremely long durations that exceed 100,000 seconds. These extremely long duration GRBs may represent the tidal disruption of main sequence stars by supermassive black holes (SMBHs). When a main sequence star comes too close to a SMBH, the gravitational pull on the star’s outer layers from the SMBH can be stronger than the star’s own gravity. This can cause the star to be completely or partly disrupted.

Material stripped from the star forms an accretion disk around the SMBH and a small fraction of the material may be expelled at relativistic velocities, driving a tidal disruption flare (TDF) that is observed as an extremely long duration GRB. Tidal disruption events (TDEs) involving the tidal shredding of main sequence stars by SMBHs are likely to be the progenitors of extremely long duration GRBs.

Reference:
Andrew Levan (2015), “Swift discoveries of new populations of extremely long duration high energy transient”, arXiv:1506.03960 [astro-ph.HE] 

Ultra-Diffuse Galaxies in the Virgo Cluster

Recently, a large number of low surface brightness (LSB) galaxies were discovered in the Coma Cluster - a large cluster of galaxies located over 300 million light years away. Galaxies do not possess well defined boundaries. They simply get fainter and fainter towards their outer regions. As a result, the size of a galaxy is defined by its effective radius, also known as the half-light radius. The effective radius is the radius within which half of the galaxy’s luminosity is emitted.

Many of the LSB galaxies in the Coma Cluster are large, with effective radii between 2 to 5 kpc. One kpc (kiloparsec) is equivalent to 3,260 light years, and for comparison, the effective radius of the Milky Way galaxy is estimated to be ~3.6 kpc. They are also extremely diffuse galaxies with central surface brightnesses between 24 to 26 mag/arcsec². LSB galaxies are vulnerable to tidal perturbations as they move through the cluster and interact with other galaxies. Tidal disruption of a LSB galaxy can strip the galaxy down to only its dense nucleus, leading to the formation of an ultra-compact dwarf (UCD) galaxy.

Mihos et al. (2015) present the discovery of three LSB galaxies in the Virgo Cluster - a much nearer cluster of galaxies located ~50 million light years away. The three LSB galaxies are dubbed VLSB-A, VLSB-B and VLSB-C. They are extremely diffuse galaxies with central surface brightnesses around 27 mag/arcsec² and effective radii between 3 to 10 kpc. All three LSB galaxies appear quite diverse in their physical properties.

VLSB-A appears as a nucleated LSB galaxy with a tidal stream extending off it. This indicates that VLSB-A is presently experiencing tidal perturbations and its diffuse component is currently being stripped away. Since the nucleus of VLSB-A shares the same structural properties as UCD galaxies, VLSB-A will most likely become a new UCD galaxy after its diffuse component is stripped away.

The properties of VLSB-B and VLSB-C are not as clear compared to VLSB-A. However, both VLSB-B and VLSB-C do show a lack of obvious tidal distortion. This suggests they may lie in the outskirts of the Virgo Cluster or may be “falling” into the Virgo Cluster for the first time. Alternatively, they may be highly dominated by dark matter, making them less susceptible to tidal stripping. Interestingly, VLSB-B is appears to host a small population of globular clusters. These globular clusters may indicate the presence of a massive halo of dark matter around the galaxy which means that the galaxy has stronger self-gravity and is therefore more protected against tidal stripping.

Surface brightness of the three LSB galaxies in the Virgo Cluster. Mihos et al. (2015).

Structural properties of the three LSB galaxies in the Virgo Cluster compared with other stellar systems, including early type galaxies in the Virgo Cluster and Fornax Cluster, and in the Local Group, as well as globular clusters and UCD galaxies in the Virgo Cluster, and the LSB galaxies found in the Coma Cluster. The dashed orange lines show the globular cluster selection box, while lines of constant surface brightness are shown in green. Mihos et al. (2015).

Reference:
Mihos et al. (2015), “Galaxies at the extremes: Ultra-diffuse galaxies in the Virgo Cluster”, arXiv:1507.02270 [astro-ph.GA]

WOH G64 is a Stellar Behemoth with a Thick Disk

Artist’s impression of WOH G64. Credit: ESO.

WOH G64 is a remarkable red supergiant (RSG) star located 163,000 light years away in the Large Magellanic Cloud (LMC), a satellite galaxy of the Milky Way. The “WHO” in the star’s name comes from the initials of its discoverers - Westerlund, Olander and Hedin; and “G64” indicates that it is the 64th entry in the catalogue published in 1981. The physical properties of WOH G64 are extreme. The star has a relatively cool surface temperature of 3,400 K, but it shines roughly 300,000 times more brightly than does the Sun. With that, WOH G64 is estimated to be 1,540 times larger in size than the Sun, making it one of the largest known RSG stars. If placed at the position of the Sun in the Solar System, the surface of WOH G64 would extend almost to the orbit of Saturn.

Observations of WHO G64 made with the Very Large Telescope Interferometer (VLTI) operated by the European Southern Observatory (ESO) in Chile revealed the presence of an enormous and thick disk of gas and dust around the star. The inner edge of the disk is at ~15 stellar radii from WOH G64, about 120 times the distance of Earth from the Sun. The rest of the disk extends far out from WHO G64, reaching almost one light year in total size. Between 3 to 9 solar masses worth of material is estimated to be in the disk of gas and dust. WHO G64 is at an advance stage of its evolution and is experiencing violent, unstable mass loss. WHO G64 started out with ~25 times the Sun’s mass, but it has since lost between one tenth and a third of its original mass. This stellar behemoth is nearing its final fate as a supernova.

References:
- Levesque et al. (2009), “The Physical Properties of the Red Supergiant WOH G64: The Largest Star Known?”, arXiv:0903.2260 [astro-ph.SR]
- Ohnaka et al. (2009), “Spatially resolved dusty torus toward the red supergiant WOH G64 in the Large Magellanic Cloud”, arXiv:0803.3823 [astro-ph]
- Westerlund, Olander & Hedin (1981), “Supergiant and giant M type stars in the Large Magellanic Cloud”, Astronomy & Astrophysics Supplement Series 43: 267-295

Properties of a Newly Discovered Super-Neptune

Bakos et al. (2015) present the discovery of HATS-7b, a transiting Super-Neptune with an orbital period of 3.185 days around a K dwarf star. The host star of HATS-7b has an effective temperature of 4,990 K, 85 percent the Sun’s mass, 82 percent the Sun’s diameter and shines with 37 percent the luminosity of the Sun. Being a K dwarf star, it is somewhat cooler and less luminous than the Sun - a G dwarf star. The detection of HATS-7b was made using the HATSouth network, comprised of a number of fully automated telescopes in the Southern Hemisphere. The primary goal of the HATSouth network is to search for transiting exoplanets.

Figure 1: Artist’s impression of a Neptune-like planet.

By measuring how much HATS-7b dims its host star when it transits in front, the planet is estimated to be 0.563 times the size of Jupiter. As HATS-7b circles its host star, it also gravitationally perturbs its host star, causing its host star to wobble back and forth. The magnitude of wobbling depends on the planet’s mass. Radial velocity measurements of the host star’s wobbling motion indicate that HATS-7b has 0.120 times the mass of Jupiter, placing it in the mass regime of super-Neptunes. For comparison, the planet Neptune has 0.054 times the mass of Jupiter, or 17.147 times the mass of Earth.

Figure 2: Transit light curve of HATS-7b phase folded to the planet’s orbital period of 3.185 days. The lower panel zooms in on the transit. Bakos et al. (2015).

Figure 3: Radial velocity measurements for the host star of HATS-7b. The gravitational perturbation from HATS-7b induces radial velocity semi-amplitude of 18.4 ± 1.9 m/s on its host star. Bakos et al. (2015).

Knowing the size and mass of HATS-7b allows the planet’s bulk composition to be constrained. Interior models of HATS-7b indicate a hydrogen-helium (H2-He) mass fraction of 18 ± 4 percent if the planet has a rock-iron core and a H2-He envelope, or a H2-He mass fraction of 9 ± 4 percent if the planet has an ice core and a H2-He envelope. If HATS-7b has a rock-iron core and a hydrogen-helium envelope, the best fit models give a core mass of 31 ± 4 Earth-masses and an envelope mass of 7 ± 1.5 Earth-masses. If instead HATS-7b has a ice core and a hydrogen-helium envelope, the best fit models give a core mass of 34.5 ± 4 Earth-masses and an envelope mass of 3.5 ± 1.5 Earth-masses.

The composition of HATS-7b is broadly similar to that of Uranus and Neptune, but quite different from Jupiter and Saturn, which are both predominantly comprised of hydrogen and helium. Super-Neptunes like HATS-7b and the recently discovered HATS-8b (also by the HATSouth network) are important for understanding the transition from ice giants (i.e. Uranus and Neptune) to gas giants (i.e. Jupiter and Saturn). HATS-7b circles in a close-in orbit around its host star at a distance of only 6 million km. This is 25 times closer than Earth is from the Sun. The dayside of HATS-7b is heated to a temperature of over 1,000 K.

Figure 4: Mass-radius diagram of super-Neptunes (planets with less than 0.18 times Jupiter’s mass) and super-Earths with accurately measured masses and radii (less than 20 percent uncertainties). Colour indicates equilibrium temperature. HATS-7b is marked with a box and Neptune is marked with a blue triangle. Abbreviations are: K: Kepler, H: HAT, HS: HATSouth, C: Corot. Bakos et al. (2015).

References:
- Bakos et al. (2015), “HATS-7b: A Hot Super Neptune Transiting a Quiet K Dwarf Star”, arXiv:1507.01024 [astro-ph.EP]
- Bayliss et al. (2015), “HATS-8b: A Low-Density Transiting Super-Neptune”, arXiv:1506.01334 [astro-ph.EP]

Properties of Oxygen Sequence Wolf-Rayet Stars

Oxygen sequence Wolf-Rayet (WO) stars are some of the rarest stars in the universe. To date, only nine WO stars are known, two of which are in binary systems. WO stars are remarkably hot, with surface temperatures ranging from 150,000 K to 210,000 K. For comparison, the Sun’s surface temperature is only 5,778 K. Core-helium burning is the process where helium in a star’s core undergoes nuclear fusion, leading to the production of oxygen and carbon. WO stars are massive stars that have passed the end of core-helium burning and have shed their outer envelopes through powerful stellar winds to reveal hotter underlying layers. Also, after a massive star exhausts the helium in its core, it naturally contracts and heats up. This is consistent with the high surface temperatures observed for WO stars.

WO stars are extremely luminous objects with hundreds of thousands of times the Sun’s luminosity. Despite their extreme luminosities, WO stars are generally smaller in size than the Sun. For example, the WO star WR102 is almost 300,000 times more luminous than the Sun, but its diameter is only 40 percent of the Sun’s. The intensely hot surfaces of WO stars are observed to be enriched with the products from core-helium burning, particularly carbon and oxygen. The helium surface mass fraction on WO stars is usually between 20 to 30 percent, but ranges from 44 percent for the least hot WO star to as low as 14 percent for the hottest known WO star. Observations have shown that the oxygen and carbon surface mass fractions can be as high as 24 and 62 percent, respectively, as in the case for the WO star WR102.

WO stars represent a very brief stage in the evolution of massive stars, predicted to be the final evolutionary stage of massive stars with initial masses between 40 to 60 times the Sun’s mass. These remarkable stars are likely to explode as type Ic supernovae in ~1,000 to 10,000 years. Type Ic supernovae are a class of stellar explosions caused by the core collapse of massive stars that have shed their outer envelopes of hydrogen and helium. As a result, type Ic supernovae do not contain hydrogen and helium. For comparison, type Ib supernovae are another class of stellar explosions involving massive stars that have only shed their outer envelope of hydrogen.

Locations of several WO stars on the Hertzsprung-Russell diagram. Also indicated are several WC stars (i.e. carbon sequence Wolf-Rayet stars). Tramper et al. (2015).

Model showing the evolution of the surface mass fractions of the WO star WR102 since the onset of core-helium burning. Tramper et al. (2015).

Reference:
Tramper et al. (2015), “Massive stars on the verge of exploding: the properties of oxygen sequence Wolf-Rayet stars”, arXiv:1507.00839 [astro-ph.SR]

A Saturn-Mass Planet beyond the Snowline of an M Dwarf Star

Gravitational microlensing is a powerful technique for detecting exoplanets that orbit their host stars beyond the snowline (i.e. the distance from a star where temperatures are cool enough for water-ice and other volatiles to condense). Beyond the snowline, a protoplanetary disk around a star is expected to contain more material, allowing for the formation of more massive planets. When a foreground star crosses the line-of-sight to a background star, the gravity of the foreground star can act as a lens, magnifying the light from the background star. The change in the brightness of the background star with time is measured in the form of a light curve. If the background star hosts a planet, the planet’s own gravity can induce a “spike” in the light curve of the background star.

Using the technique of gravitational microlensing, Fukui et al. (2015) present the discovery of a Saturn-mass planet with ~0.34 times the mass of Jupiter orbiting an M dwarf star with ~0.39 times the Sun’s mass at a projected separation of either ~0.74 AU (close model) or ~4.3 AU (wide model). The planet is identified as OGLE-2012-BLG-0563Lb and it is the 5th sub-Jupiter-mass (i.e. between 0.2 to 1.0 times the mass of Jupiter) to be found around an M dwarf star through gravitational microlensing. Although it is clear that there is a population of sub-Jupiter-mass planets around M dwarf stars, there appears to be a paucity of Jupiter-mass planets (i.e. planets with ~1 to 2 times the mass of Jupiter) around the same type of star. This suggests that planet formation via the core-accretion mechanism rarely produces Jupiter-mass planets around M dwarf stars due to the lack of material in the protoplanetary disk.

Reference:
Fukui et al. (2015), “OGLE-2012-BLG-0563Lb: a Saturn-mass Planet around an M Dwarf with the Mass Constrained by Subaru AO imaging”, arXiv:1506.08850 [astro-ph.EP

Two Record-Breaking Compact Stellar Systems

Following the discovery of the densest known galaxy, M60-UCD1, Sandoval et al. (2015) present the discovery of M59-UCD3 and M85-HCC1 - two record-breaking compact stellar systems. Since the density of stars in a galaxy decreases gradually away from the center, one way to express the size of a galaxy is by its half-light radius. The half-light radius is basically the size of the volume of space that is contributing to half the galaxy’s brightness and it can also apply to other stellar systems.

M59-UCD3 is an ultracompact dwarf (UCD) galaxy similar in size to M60-UCD1. It has a half-light radius of roughly 70 light years, but it is 40 percent more luminous than M60-UCD1. This makes M59-UCD3 the new densest known galaxy. M59-UCD3 is estimated to contain roughly ~180 million times the Sun’s mass and this means an average density of roughly 30 solar masses per volume of space one light year across. For comparison, an observer on Earth can see ~6,000 stars with unaided eyes under “typical” dark sky conditions. However, an observer in the core of M59-UCD3 would see roughly a million stars in the sky.

M85-HCC1 is an extremely compact stellar system with a half-light radius of roughly 6 light years, similar in size to a typical globular cluster. It is estimated to contain ~12 million times the Sun’s mass and has an average density of roughly 3,000 solar masses per volume of space one light year across. This is ~100 times the density of M59-UCD3. For comparison, the nearest star to the Sun is Proxima Centauri, located 4.24 light years away.

M59-UCD3 and M85-HCC1 are estimated to be ~9 billion and ~3 billion years old, respectively. They are most likely the remnant cores of what were once much larger galaxies that got tidally-stripped. This scenario can be tested for M59-UCD3 by looking to see if it contains an “overweight” supermassive black hole (SMBH) since almost all massive galaxies host SMBHs. M60-UCD1, the previous record holder for the densest known galaxy, is known to host an “overweight” SMBH comprising a whopping ~15 percent of the galaxy’s total mass.

References:
- Sandoval et al. (2015), “Hiding in plain sight: record-breaking compact stellar systems in the Sloan Digital Sky Survey”, arXiv:1506.08828 [astro-ph.GA]
- Strader et al. (2014), “The Densest Galaxy”, arXiv:1307.7707 [astro-ph.CO]
- Seth et al. (2014), “A Supermassive Black Hole in an Ultracompact Dwarf Galaxy”, arXiv:1409.4769 [astro-ph.GA]

Properties of a Pair of Juvenile Brown Dwarfs

Observations of DE0823-49 indicated that it is a system composed of a pair of brown dwarfs identified as DE0823-49A (primary) and DE0823-49B (secondary) with spectral types L1.5 ± 0.6 and L5.5 ± 1.1, respectively. Both objects orbit each other every ~248 days. The estimated effective temperatures of DE0823-49A and DE0823-49B are 2,150 ± 100 K and 1,670 ± 140 K, respectively. Models predict that both objects have masses between 0.028 to 0.063 times the Sun’s mass for DE0823-49A and between 0.018 to 0.045 times the Sun’s mass for DE0823-49B, with a mass ratio (i.e. mass of DE0823-49A relative to DE0823-49B) of 0.64 to 0.74. This places both objects in the brown dwarf mass regime and also below the lithium-burning mass limit of 0.065 times the Sun’s mass.

The age DE0823-49 is estimated to be between 80 million to 500 million years. This is consistant with the presence of lithium detected on DE0823-49A which implies it has a mass of less than 0.065 times the Sun’s mass. An object more massive that that is expected to burn away its lithium. Given that DE0823-49A is less than 0.065 times the Sun’s mass and still relatively hot, and that brown dwarfs cool gradually as they age, the age of DE0823-49A cannot exceed 500 million years. Since both DE0823-49A and DE0823-49B formed together, they should have the same age. With an age of less than 500 million years, DE0823-49 is a pair of juvenile brown dwarfs. It is also relatively nearby, located only 67.5 light years away.

Reference:
Sahlmann et al. (2015), “DE0823-49 is a juvenile binary brown dwarf at 20.7 pc”, arXiv:1505.07972 [astro-ph.SR]

Hot-Neptune GJ 436b has a Comet-Like Tail

GJ 436b is a Neptune-mass exoplanet in a close-in orbit around its parent star. The planet is so intensely irradiated by its parent star that its own gravity is not strong enough to hold on to the hydrogen in its atmosphere, causing the hydrogen to escape into space to form a large exospheric cloud which surrounds and trails the planet. This discovery was made in the ultraviolet band with the Space Telescope Imaging Spectrograph (STIS) on the Hubble Space Telescope (HST).

Artist’s impression of GJ 436b. Image credit: Mark Garlick/University of Warwick.

GJ 436b transits in front of its parent star every 2.64 days. The transit depth in the ultraviolet band is 56.3 ± 3.5 percent (1σ), far deeper than the transit depth of 0.69 percent in the optical band. This means that over half the disk of the star is eclipsed in the ultraviolet band. Furthermore, in the ultraviolet band, the transit of GJ 436b starts ~2 hours before and ends more than 3 hours after the ~1 hour optical transit. Such a transit signature is believed to be due to the passage of a large cloud of hydrogen surround and trailing the planet.

The mass-loss rate of GJ 436b is estimated to be between ~10² to 10³ tons per second. At this rate, the planet loses only ~0.1 percent of its atmosphere per billion years, far too small to deplete the planet’s atmosphere over the lifetime of its parent star. However, planets similar to GJ 436b that orbit closer to their host stars can experience much more dramatic mass-loss, possibly even causing them to erode down to their rocky cores.

Reference:
Ehrenreich et al., “A giant comet-like cloud of hydrogen escaping the warm Neptune-mass exoplanet GJ 436b”, Nature 522, 459-461 (25 June 2015)

Illuminating the Dark Sides of Tidally-Locked Planets

Figure 1: Artist’s impression of a habitable planet.

Red dwarf stars are by far the most common stars in the Universe. Observations by NASA’s Kepler space telescope have shown that Earth-sized planets are ubiquitous around red dwarf stars. Red dwarf stars are much smaller and much less luminous than stars like the Sun. For a planet around a red dwarf star to receive as much insolation as Earth gets from the Sun, the planet needs to orbit much closer to the star. As a result, the planet is likely tidally-locked, with a permanent dayside hemisphere and a permanent nightside hemisphere.

If an extraterrestrial civilisation evolved from such a planet or migrated in from elsewhere, it may wish to make better use of the planet’s surface by illuminating the permanent nightside. The most obvious way is to place mirrors in space to illuminate the planet’s dark side. Either a single large mirror could be stationed at the L2 Lagrange point or a fleet of smaller mirrors could be placed into orbit around the planet (Figure 2). From a feasibility standpoint, a fleet of smaller mirrors is less technically challenging than a single large mirror at the L2 Lagrange point.

Figure 2: Schematic illustration of three methods of dark-side illumination (not to scale). Planetary grayscale bands indicate different levels in stellar illumination. In the three cross-sectional drawings, (a) shows a large circular or annular mirror stationed at the L2 Lagrange point, (b) shows multiple small mirrors in circular orbits, (c) shows multiple small mirrors in elliptical orbits designed to maximise the duty cycle of the mirrors.

Illuminating the dark side of a tidally-locked planet with a fleet of orbiting mirrors will warm up the planet’s dark side and lead to a smaller day-night temperature contrast on the planet. This will slow the transfer of heat from the dayside to the nightside, causing the dayside to warm up. However, the mirrors can obscure some of the starlight falling on the planet’s dayside and this can keep the dayside from warming up excessively.

To provide sufficient illumination, a fleet of orbiting mirrors is expected to have a total reflective surface area comparable to the cross-sectional area of the planet. As a result, if the planet were to transit its parent star, the presence of such a fleet of orbiting mirrors could be detected from the unique transit light curve. In order to model and understand such a transit light curve, the fleet of mirrors orbiting the planet can be approximated as a translucent annulus around the planet.

During the initial phases of the transit (i.e. ingress), only the mirrors block starlight, so the light curve decreases relatively gradually. When the planet itself starts to transit, the light curve begins to decrease more steeply. The light curve decreases gradually again when the planet is completely in front of the star and the remaining fleet of mirrors, approximated as a translucent annulus around the planet, blocks more starlight. Once the fleet of mirrors are completely in front of the star, the transit light curve resembles that of a larger planet. 

A fleet of orbiting mirrors basically lengthens the transit duration and deepens the transit depth. If the fleet of orbiting mirrors is designed for maximum efficiency such that all the reflected starlight is directed to illuminate the planet’s dark side, then the eclipse duration (i.e. duration where the planet passes behind the star) will be smaller than the transit duration. A fleet of orbiting mirrors around a planet, if present, could be detectable by the James Webb Space Telescope (JWST).

Figure 3: Transit light curves for a planet with 2 Earth radii orbiting an M5 red dwarf star (dashed line: planet alone), the same planet surrounded by a fleet of orbiting mirrors extending to 3 planetary radii (solid line: mirror + planet), and a planet without mirrors but large enough to produce a light curve with the same depth of transit (dotted line: large planet). Korpela, Sallmen & Greene (2015).

Figure 4: Transit light curves that result when a planet with 2 Earth radii, located in the middle of an M5 red dwarf star’s habitable zone, transits in front of the star. In all cases, the planet is surrounded by a fleet of orbiting mirrors extending to 3 planetary radii (solid), 2 planetary radii (dotted) and 10 planetary radii (dashed). Korpela, Sallmen & Greene (2015).

Reference:
Korpela, Sallmen & Greene (2015), “Modeling Indications of Technology in Planetary Transit Light Curves -- Dark-side illumination”, arXiv:1505.07399 [astro-ph.EP]

Forming Graphite-Like Carbon in the Atmosphere of GJ 436b

GJ 436b is a Neptune-sized planet that orbits around a relatively small M-dwarf star located only 33 light years away. The planet is ~10 times closer to its parent star than Mercury is to the Sun and it completes one orbit around its host star in just 2.64 days. Being so close to its parent star, the planet is referred to as a hot-Neptune and its dayside is heated to a temperature of ~800 K. GJ 436b is estimated to have 22 times the mass of Earth and 4.3 times the radius of Earth. A rocky core is predicted to exist at the center of GJ 436b. Surrounding the rocky core is a mantle of exotic high pressure water-ice. Finally, an envelope of hydrogen and helium gas forms the outermost layer of GJ 436b.

By observing the thermal emission from the dayside of GJ 436b, Stevenson et al. (2010) found that the methane (CH4) to carbon monoxide (CO) ratio in the planet’s atmosphere is at least ~10,000 times smaller than predicted. This is puzzling because in a hydrogen-dominated atmosphere at the temperatures found on GJ 436b, carbon in the atmosphere should prefer CH4 over CO. A likely reason for the lower-than-expected abundance of atmospheric CH4 is that the process of photodissociation can decompose CH4 into fragments which then react with each other or with other CH4 molecules to form C2 hydrocarbons (i.e. hydrocarbon molecules containing 2 carbon atoms). Successive reactions can lead to the formation of C3 and C4 hydrocarbons.

Using laboratory experiments to mimic the chemical processes in hydrogen-dominated atmospheres containing C3 and C4 hydrocarbons at elevated temperatures, Dangi et al. (2015) show that the surfaces of silicon grains can act as a catalyst, allowing the conversion of C3 and C4 hydrocarbons to carbonaceous refractory matter with carbon content greater than 90 percent. On GJ 436b, silicon grains can be supplied by micrometeoroids. This explains the low CH4 to CO ratio in the atmosphere of GJ 436b, whereby CH4 photochemically converts to higher order hydrocarbons such as C3 and C4 hydrocarbons. The catalytic surfaces of micrometeoroids then convert these hydrocarbons to refractory graphite-like carbon.

References:
- Stevenson et al. (2010), “Possible thermochemical disequilibrium in the atmosphere of the exoplanet GJ 436b”, arXiv:1010.4591 [astro-ph.EP]
- Dangi et al. (2015), “Toward the Formation of Carbonaceous Refractory Matter in High Temperature Hydrocarbon-rich Atmospheres of Exoplanets Upon Micrometeoroid Impact”, ApJ 805:76 (7pp)

Saturn’s Phoebe Ring is Even Larger than Thought

Figure 1: Artist’s conception of Saturn’s giant Phoebe ring as seen in infrared light. Credit: NASA/JPL/Space Science Institute.

Saturn’s giant Phoebe ring, discovered in 2009, is believed to be formed by particles thrown off from Phoebe, a distant moon of Saturn whose surface is as dark as coal. When it was discovered, the Phoebe ring was observed to lie between 128 and 207 Saturn radii from the planet, with a vertical extent of 40 Saturn radii. A recent study published by Hamilton et al. (2015) using new observations from NASA’s Wide-field Infrared Survey Explorer (WISE) mission found that the Phoebe ring is much larger than previously measured.

WISE acquired an infrared image of the entire Phoebe ring which shows the ring spanning between 100 and 270 Saturn radii from the planet. This is well beyond the orbit of Phoebe which goes around Saturn between 180 and 270 Saturn radii from the planet. The Phoebe ring covers an area of sky ~7,000 times larger than Saturn itself. It is likely that small moons with Phoebe-like orbital inclinations that orbit Saturn more distant than Phoebe itself are additional contributors of ring material. If many of these small moons turn out to be undiscovered 100 m sized or kilometre-sized objects, then they are likely to be an important additional source of ring material.

Figure 2: Measured profile of Saturn’s giant Phoebe ring. Ring flux is clearly detectable to at least 270 Saturn radii, well beyond Phoebe’s maximum distance of 250 Saturn radii from Saturn. All ring profiles agree well inward to about 100 Saturn radii at which point scattered light from Saturn becomes problematic. Hamilton et al. (2015).

Reference:
Hamilton et al., “Small particles dominate Saturn’s Phoebe ring to surprisingly large distances”, Nature 522, 185-187 (11 June 2015)

Birth of a Supermassive Black Hole

The existence of supermassive black holes (SMBHs) with ~1 billion times the Sun’s mass in the early Universe is a puzzle because there is insufficient time for stellar mass black holes with ~10 times the Sun’s mass to grow into SMBHs. A study by Matsumoto et al. (2015) propose that the conditions present in the early Universe can allow protostars to quickly grow to form supermassive stars (SMSs) with ~100,000 times the Sun’s mass. When these SMSs exhaust their nuclear fuel, they collapse directly to form massive black holes with ~10,000 times the Sun’s mass. These massive black holes can seed the formation of SMBHs in the early Universe.

A SMS is expected to have a radius of roughly 1 billion kilometres and a lifespan of only ~1 million years. Once the SMS exhausts its nuclear fuel, it starts collapsing into a massive black hole. During the collapse, relativistic jets can be launched from the accretion disk around the growing, nascent black hole. The relativistic jets can break through the star’s surface to produce an ultra-long gamma ray burst (ULGRB) with duration of ~10,000 to 1,000,000 seconds. Gamma-ray bursts (GRBs) are generally classified into short gamma-ray bursts (SGRBs) with durations under 2 seconds and long gamma-ray bursts (LGRBs) with duration greater than 2 seconds. ULGRBs are a third class of GRBs with durations of ~10,000 seconds or more. The collapse of SMSs into massive black holes may be observed as ULGRBs.

Reference:
Matsumoto et al. (2015), “Direct Collapse Black Holes Can Launch Gamma-Ray Bursts and Get Fat to Supermassive Black Holes?”, arXiv:1506.05802 [astro-ph.HE]

Detection of a Brown Dwarf Orbiting a Sun-Like Star

Using a method known as gravitational microlensing, where the gravitational field of a foreground object distorts and magnifies the light from a background star, Ranc et al. (2015) present the discovery of MOA-2007-BLG-197Lb, a brown dwarf in orbit around a Sun-like star. Brown dwarfs are substellar objects that occupy the mass range between gas giant planets and the least massive stars. In fact, MOA-2007-BLG-197Lb is the first brown dwarf to be detected around a Sun-like star through gravitational microlensing. Previous detections of brown dwarf companions to stars through gravitational microlensing have all been detections of brown dwarfs around diminutive red dwarf stars. MOA-2007-BLG-197Lb is a brown dwarf with 41 ± 2 times the mass of Jupiter and it orbits a Sun-like star with 0.82 ± 0.04 times the Sun’s mass at an observed projected separation of 4.3 ± 0.1 AU (i.e. one AU is the distance of Earth from the Sun).

Both MOA-2007-BLG-197Lb and its host star are located at a distance of roughly 14,000 light years from Earth. MOA-2007-BLG-197Lb is a valuable find because there appears to be a lack of brown dwarfs around Sun-like stars. One reason could be because brown dwarfs orbiting Sun-like stars lose angular momentum due to tidal interactions with their host stars, causing the brown dwarfs to spiral in and fall into their host stars. Nonetheless, gravitational microlensing is a great technique for detecting free-floating brown dwarfs, brown dwarf binaries and brown dwarf companions to stars. Future detections of brown dwarfs around Sun-like stars will offer better insight on the population distribution of brown dwarfs.

Light curve of the gravitational microlensing event that led to the detection of MOA-2007-BLG-197Lb and the best-fit model (solid line). The insert on the right shows a zoom-in of the “spike” in the light curve generated by the presence of MOA-2007-BLG-197Lb. Ranc et al. (2015).

Reference:
Ranc et al. (2015), “MOA-2007-BLG-197: Exploring the brown dwarf desert”, arXiv:1505.06037 [astro-ph.EP]

The Three Planets of Kepler-138

Figure 1: Artist’s impression of a rocky planet.

Located roughly 200 light years from Earth, Kepler-138 is a red dwarf star that is much fainter and cooler than the Sun. Observations by NASA’s Kepler space telescope have led to the detection of three planets circling Kepler-138. The three planets were detected using the transit technique and are identified as Kepler-138 b, Kepler-138 c and Kepler-138 d. When a planet passes in front of (i.e. transits) its parent star, it causes a slight drop in the observed brightness of the star, allowing the size of the planet to be measured. Additionally, the interval between consecutive transits is the planet’s orbital period. The transit technique allows the sizes and orbital periods of the planets around Kepler-138 to be determined.

In a multi-planet system, gravitational perturbations between neighbouring planets can cause the transit timing for each planet to vary. This effect is particularly sensitive to planets that are closely spaced or near orbital resonances. The three planets around Kepler-138 are in orbital resonances and exhibit observable transit timing variations (TTV). Kepler-138b and Kepler-138c orbit near the 4:3 resonance (i.e. Kepler-138 b completes 4 orbits for every 3 orbits of Kepler-138 c); while Kepler-138 c and Kepler-138 d orbit near the 5:3 resonance (i.e. Kepler-138 c completes 5 orbits for every 3 orbits of Kepler-138 d).

Figure 2: TTV signals of the three planets orbiting Kepler-138. (a) TTV of Kepler-138 b; (b) TTV of Kepler-138 c; (c) TTV of Kepler-138 d. Jontof-Hutter et al. (2015).

Figure 3: Mass-radius diagram of well characterized planets smaller than 2.1 Earth radii. Prior exoplanet characterizations and 1σ uncertainties are shown as grey points. Black points from left to right are Mercury, Mars, Venus and Earth. Red data points are the results for the planets of Kepler-138. Open circles mark previously measured masses for Kepler-138 c and Kepler-138 d. Error bars mark published 1σ uncertainties for the planets of Kepler-138, and masses and radii of all other characterized exoplanets within this size range. The curves mark bulk densities of 1 g/cm³, 3 g/cm³ and 10 g/cm³. Jontof-Hutter et al. (2015).

Measuring the TTV signals allows the masses of the three planets around Kepler-138 to be estimated. By knowing the mass and size of a planet, it allows the planet’s density and its bulk composition to be determined to see whether the planet is predominantly made up of rock, water or gas. Here are the properties of the three planets around Kepler-138:
- Kepler-138 b has an orbital period of 10.3 days, 0.522 times the Earth’s radius, 0.066 times the Earth’s mass, a density of 2.6 g/cm³ and it receives an incident stellar flux that is 6.81 times what Earth gets from the Sun.
- Kepler-138 c has an orbital period of 13.8 days, 1.197 times the Earth’s radius, 1.970 times the Earth’s mass, a density of 6.2 g/cm³ and it receives an incident stellar flux that is 4.63 times what Earth gets from the Sun.
- Kepler-138 d has an orbital period of 23.1 days, 1.212 times the Earth’s radius, 0.640 times the Earth’s mass, a density of 2.1 g/cm³ and it receives an incident stellar flux that is 2.32 times what Earth gets from the Sun.

Kepler-138 b, the innermost of the three planets around Kepler-138, is roughly the same size as Mars and its density is also consistent with a rocky composition like Mars. With less than one-tenth the Earth’s mass, Kepler-138 b is currently the lightest known exoplanet with a measured mass. The two outer planets, Kepler-138 c and Kepler-138 d are both similar in size to Earth. However, Kepler-138 c and Kepler-138 d have very different densities. Kepler-138 c has a high density which indicates a rocky composition, while the density of Kepler-138 d is less than half of Earth’s and suggests that a large portion of its composition is in the form of low density materials such as water or hydrogen. All three planets around Kepler-138 receive much more incident stellar flux than Earth gets from the Sun and are too hot to be habitable.

Reference:
Jontof-Hutter et al., “The mass of the Mars-sized exoplanet Kepler-138 b from transit timing”, Nature 522, 321-323 (18 June 2015)

Classifying Planets, Brown Dwarfs & Stars

Figure 1: Artist’s impression of a giant planet.

Stars are objects with sufficient mass to sustain hydrogen fusion in their cores, while sub-stellar objects (i.e. brown dwarfs and planets) have masses below what is required to sustain hydrogen fusion. The mass boundary between stars and sub-stellar objects lie at about 80 times Jupiter’s mass and seems to be quite clear. However, the mass boundary between what is considered a planet and what is considered a brown dwarf is unclear. One definition is that if a planet is defined as a sub-stellar object that has not undergone deuterium burning at any point in its life; then the mass boundary between planets and brown dwarfs is about 13 times Jupiter’s mass. Unfortunately, this distinction is weak because deuterium burning only occurs for a brief period and has a negligible impact on the future evolution of a brown dwarf.

Hatzes & Rauer (2015) propose a new definition for planets, brown dwarfs and stars by presenting the mass-density relationship for objects ranging from planets with ~0.01 times Jupiter’s mass through stars with more than 80 times Jupiter’s mass. The mass-density diagram shows 3 distinct regions. Objects below 0.3 times Jupiter’s mass show considerable scatter, objects between 0.3 to 60 times Jupiter’s mass have a positive slope and objects over 60 times Jupiter’s mass follow a negative slope. Objects below 0.3 times Jupiter’s mass are simply referred to as low-mass planets. For objects between 0.3 to 60 times Jupiter’s mass, there appears to be no distinguishing characteristic to separate the lower mass members which are clearly giant planets from the higher mass members which are generally considered to be brown dwarfs. As a result, all objects between 0.3 to 60 times Jupiter’s mass are simply defined by Hatzes & Rauer (2015) as giant planets.

The slope becomes negative for objects over 60 times Jupiter’s mass. However, at 60 times Jupiter’s mass, an object is still not massive enough to sustain hydrogen fusion in its core. As a result, Hatzes & Rauer (2015) suggest that objects with 60 to 80 times Jupiter’s mass are considered bona fide brown dwarfs and objects over 80 times Jupiter’s mass are stars. The change in slope at 60 times Jupiter’s mass is a good mass boundary between what constitutes a planet and what constitutes a brown dwarf. This is because the change in slope indicates a change in the interior structure between planets and brown dwarfs. To summarise, this new definition goes like this - objects less than 0.3 times Jupiter’s mass are low-mass planets, objects between 0.3 to 60 times Jupiter’s mass are giant planets and objects between 60 to 80 times Jupiter’s mass are brown dwarfs.

Figure 2: The densities and masses of stars (red squares), giant planets and brown dwarfs, and low mass planets. Triangles represent Kepler discoveries and dots are CoRoT exoplanets. Ground-based discoveries for high mass giant planets are shown by pentagons. Hatzes & Rauer (2015).

Figure 3: The points from Figure 2 shown in the mass-radius plane. Hatzes & Rauer (2015).

Reference:
Hatzes & Rauer (2015), “A Definition for Giant Planets Based on the Mass-Density Relationship”, arXiv:1506.05097 [astro-ph.EP]

A Venus-Mass Planet Orbiting a Brown Dwarf

Brown dwarfs are objects with masses that fill the gap between stars and planets. Like stars, brown dwarfs are also known to host planets around them. Most of these planets are gas giants with more than 1/10th the mass of their host brown dwarfs. These systems tend to resemble scaled down versions of binary star systems rather than normal planetary systems around stars.

Nevertheless, an object identified as OGLE-2012-BLG-0358Lb is known to be a planetary-mass object (i.e. secondary object) with 1.9 ± 0.2 times the mass of Jupiter. It orbits a brown dwarf (i.e. primary object) with 0.022 times the Sun’s mass at a projection separation of roughly 0.87 AU. This system has a low secondary-to-primary mass ratio of roughly 0.08, making it more analogous to normal planetary systems around stars than scaled down versions of binary systems.

Udalski et al. (2015) present the discovery of a system consisting of a brown dwarf orbited by a Venus-mass planet, with both objects orbiting a very low-mass star. This system was discovered via a technique known as gravitational microlensing and the event is dubbed OGLE-2013-BLG-0723. Gravitational microlensing occurs when a foreground object crosses the line of sight to a background star and the gravitational field of the foreground object acts like a lens, magnifying the brightness of the background star. In this case, the foreground object is the newly discovered system.

The three objects in the system are identified as OGLE-2013-BLG-0723Bb for the Venus-mass planet, OGLE-2013-BLG-0723B for the brown dwarf and OGLE-2013-BLG-0723A for the very low-mass star. The Venus-mass planet has 0.69 ± 0.06 the mass of Earth. It orbits at a projected separation of 0.34 ± 0.03 AU from a brown dwarf with 0.031 ± 0.003 times the Sun’s mass.

The brown dwarf itself forms a binary system with a very low-mass star with 0.097±0.009 times the Sun’s mass. Both the brown dwarf and the very low-mass star have a projected separation of 1.74 ± 0.15 AU. The Venus-mass planet and brown dwarf in the OGLE-2013-BLG-0723 system can be considered either as a scaled down version of a planet plus star system or as a scaled up version of a moon plus planet system.

Light curve of the microlensing event OGLE-2013-BLG-0723, including models with and without the planet. Left inset shows the planetary anomaly, which includes not just the obvious spike, but also a more extended low level depression. Udalski at al. (2015).

Reference:
Udalski at al. (2015), “A Venus-Mass Planet Orbiting a Brown Dwarf: Missing Link between Planets and Moons”, arXiv:1507.02388 [astro-ph.EP]

Identification of a Young Planetary-Mass Brown Dwarf

Young, planetary-mass brown dwarfs provide a good proxy for the study of gas giant planets around stars because these objects exist in isolation and are not overwhelmed by the glare of a parent star. Brown dwarfs cool as they age. As a result, young brown dwarfs are hotter and more luminous, making them easier to observe. Brown dwarfs are identified by the spectral types - M, L, T and Y, in order of decreasing effective temperature. The spectral type of a brown dwarf changes as it cools. More massive brown dwarfs take longer to cool.

Planetary-mass brown dwarfs that are younger than ~120 million years old are expected to be T dwarfs (i.e. T-type spectral class). These objects, if located in the Sun's neighbourhood (i.e. nearer than ~65 light years away) are bright enough to be studied by next generation observatories such as the James Webb Space Telescope (JWST) and the Giant Magellan Telescope (GMT). So far, only one isolated planetary-mass T dwarf candidate is known. This object is identified as CFBDSIR J214947.2-040308.9 and it is estimated to have between 4 to 7 times the mass of Jupiter.

Gagné et al. (2015) present the discovery of SDSS J111010.01+011613.1 (hereafter SDSS J1110+0116), a young, planetary-mass T dwarf. Measurements indicate that SDSS J1110+0116 is located at a distance of roughly 60 light years and it is a member of the AB Doradus moving group - a group of objects that formed in the same natal cluster and have since dispersed. J1110+0116 is estimated to have roughly 10 to 12 times the mass and 1.18 ± 0.02 times the radius of Jupiter. The mass of J1110+0116 is well in the planetary-mass regime. J1110+0116 is a young object with an estimated age of only 110 to 130 million years old. It is still glowing hot from the heat acquired during its formation and its effective temperature is estimated to be 940 ± 20 K.

Reference:
Gagné et al. (2015), "SDSS J111010.01+011613.1: A New Planetary-Mass T Dwarf Member of the AB Doradus Moving Group", arXiv:1506.04195 [astro-ph.SR

Deuterium Fusion in the Cores of Inflated Hot-Jupiters

Hot-Jupiters are a class of Jupiter-like exoplanets that reside in very close-in orbits around their parent stars. These planets are strongly irradiated and have intensely hot daysides. Like Jupiter, they are gas giant planets primarily composed of hydrogen and helium. Observations have shown that a large proportion of hot-Jupiters are inflated in size. Their radii appear too large even after accounting for the strong irradiation from the parent star. A number of mechanisms have been proposed to explain the inflated radii of hot-Jupiters. These include tidal heating and Ohmic heating. Basically, the inflated radii of hot-Jupiters require mechanisms that can deposit additional sources of energy within the bulk of the planet.

A study by Ouyed & Jaikumar (2015) suggests yet another mechanism that could account for the inflated radii of hot-Jupiters. Deuterium is an isotope of hydrogen and deuterium-deuterium (DD) fusion in the deep interior of hot-Jupiters can provide an extra source of energy. A problem with this process is that it requires extremely high temperatures (~100,000 K) in a layer of deuterium around the planet’s core. This an order of magnitude larger than the core temperatures typically found in hot-Jupiters. Nevertheless, DD fusion becomes more plausible if it instead occurs in the solid deuterated core of the planet. This is because the likelihood for DD fusion is significantly enhanced in a solid deuterated substrate and the temperatures typically found in the deep interiors of hot-Jupiter (~10,000 K) are sufficient to sustain the fusion process.

In the interior of a hot-Jupiter, core erosion takes place at a temperature of roughly 10,000 K or more. The high surface temperature on a hot-Jupiter induces the conditions needed for core erosion to occur. Erosion frees the deuterium needed to react with the non-eroded part of the core. This process can continue for billions of years, supplying the deuterium needed to sustain DD fusion and maintain the core region at ~10,000 K. The energy produced from DD fusion can keep the hot-Jupiter inflated, and in some causes, even over-inflating the hot-Jupiter.

Reference:
Ouyed & Jaikumar (2015), “Nuclear Fusion in the Deuterated cores of inflated hot Jupiters”, arXiv:1506.03793 [astro-ph.EP

Capturing Planetesimals from a Passing Star

A population of far-flung objects orbit the Sun in the outer regions of the Solar System, between the Kuiper belt and the more distant Oort cloud. These objects can be referred to as Sednoids, after 90377 Sedna - the first of its kind to be discovered in 2003. Roughly a dozen or so Sednoids have been detected so far. These objects orbit the Sun at distances of between 150 to 1,500 AU and come no closer than 30 AU to the Sun. They are too far to be Kuiper belt objects, but not far enough to be Oort cloud objects. At present, there is no explanation for how these objects came to be.

Figure 1: Artist’s impression of a planetesimal orbiting far from its host star.

Sednoids are also intriguing because their orbital inclination with respect to the ecliptic range between 10° to 30°, and their argument of perihelion cluster around 340° ± 55°. The terms “orbital inclination” and “argument of perihelion” are parameters used to describe an object’s orbit around the Sun. Basically, the ecliptic is defined as a plane of reference that is coplanar with Earth’s orbit around the Sun, while the argument of perihelion is the angle between an object’s ascending node (i.e. point where the object’s orbit around the Sun crosses the ecliptic from south to north) and its perihelion (i.e. point along the object’s orbit where it is closest to the Sun) in the direction of the object’s orbital motion.

The clustered distribution of the argument of perihelion of the Sednoids is odd because the effect of precession should have scattered the argument of perihelion within a span of several million years. As a consequence, either the clustering happened recently (i.e. within several million years), which is unlikely, or the presence of at least one massive object in the outer Solar System is keeping the Sednoids in the currently observed clustered distribution of their argument of perihelion. In fact, one or more massive objects with several times the mass of Earth could be orbiting the Sun between 200 and 300 AU.

Figure 2: Orbital distributions of planetesimals for the Sun (left) and for the encountering star (right). The top panels give inclination as a function of semi-major axis; the bottom panels give the orbital eccentricity. The red bullets give the orbital distributions of the planetesimals native to the Sun (assuming its disk extended to 90 AU), the light blue bullets are native to the passing star. Both initial planetesimal disks are strongly perturbed beyond about 30 AU, but within this distance they are hardly affected. Jilkova et al. (2015).

A study by Jilkova et al. (2015) propose that the Sednoids are actually a captured population of planetoids from another star that came close to the young Sun when the Sun will still in its crowded birth cluster. Using the orbital parameters of the Sednoids to reconstruct the trajectory of the passing star, it was found that the star had roughly 1.8 times the Sun’s mass and came as close as ~340 AU to the Sun with a relative velocity of roughly 4.3 km/s. During the encounter event, the Sun acquired planetesimals from the passing star and the passing star also acquired planetesimals from the Sun. The region of space where the Sednoids orbit the Sun is estimated to contain ~930 planetesimals.

Since more massive stars have shorter lives, the star that once passed close to the Sun should have already evolved into a carbon-oxygen white dwarf with ~0.6 times the Sun’s mass a few billion years ago. As the star evolved into a white dwarf, it lost a significant fraction of its mass which weakened its gravitational grip. As a result, the planetesimals it acquired from the Sun were expelled into interstellar space as free-floating planetesimals.

Reference:
Jilkova et al. (2015), “How Sedna and family were captured in a close encounter with a solar sibling”, arXiv:1506.03105 [astro-ph.EP]

An Intensely Irradiated Hot-Jupiter in a Polar Orbit

Delrez et al. (2015) present the discovery of WASP-121b, a hot-Jupiter in a short-period polar orbit around its host star. The orbit of WASP-121b causes it to transit its host star every 1.27 days. The host star of WASP-121b is an F-type main sequence star with 1.35 times the mass and 1.46 times the radius of the Sun. It shines with 3.3 times the Sun’s luminosity. Spectroscopic and photometric observations show that WASP-121b has 1.18 times the mass and 1.87 times the radius of Jupiter. This gives WASP-121b a mean density of only 1/5th the density of Jupiter.

Figure 1: Artist’s impression of an exoplanet obscuring its host star. Image credit: Pauline Moss.

The large radius of WASP-121b indicates that it is significantly inflated. WASP-121b orbits so close to its host star, it is just ~1.15 times the minimum distance from its host star where it will start to become tidally disrupted. Due to the proximity to its host star, WASP-121b receives an extreme amount of irradiation which heats the planet to an estimated temperature of roughly 2,360 K. WASP-121b joins a handful of intensely irradiated planets with super-inflated radii.

By measuring a phenomenon known as the Rossiter-McLaughlin (RM) effect, WASP-121b is found to be in a polar orbit around its host star. The orbit of WASP-121b passes almost directly over the poles of its host star. Being so close to its host star, WASP-121b is expected to be significantly deformed due to the intense tidal force it is subjected to. WASP-121b is predicted to be deformed into the shape of a triaxial ellipsoid, with the longest axis pointed towards its host star. The longest axis is estimated to be twice Jupiter’s diameter, while the planet’s shortest axis (i.e. the planet’s polar axis) is estimated to be 1.79 times Jupiter’s diameter.

Figure 2: Light curve showing the transit of WASP-121b in front of its host star. The transit depth indicates that the planet is 1.87 times the radius of Jupiter. Delrez et al. (2015).

Figure 3: Radial velocity curve for the host star of WASP-121b. The amplitude of the radial velocity curve indicates that the planet is 1.18 times the mass of Jupiter. Delrez et al. (2015).

Figure 4: Observed Rossiter-McLaughlin (RM) effect as WASP-121b transits its host star. Delrez et al. (2015).

Reference:
Delrez et al. (2015), “WASP-121 b: a hot Jupiter in a polar orbit and close to tidal disruption”, arXiv:1506.02471 [astro-ph.EP]