Thursday, July 24, 2014

Ultra-Dense Ocean on a Neutron Star

A neutron star is an ultra-dense remnant core leftover from the violent demise of a massive star. It packs roughly as much mass as the Sun in an incredibly tiny volume measuring just several kilometres across. A spoonful of its material would contain a mass of roughly a billion tons. If the neutron star has a sufficiently close stellar companion, it can strip material from the companion in a process known as accretion. The accreted material can lead to the formation of an ocean on the neutron star. This ultra-dense and exotic ocean is comprised of elements with atomic number Z = 6 and larger. Most of these elements are formed from nuclear burning of the accreted hydrogen and helium from the companion star. Here, the ions behave like a liquid, hence the term “ocean”. Nonetheless, it is in no way like the oceans on Earth. The densities, pressures and temperatures are so extreme that they are only comprehensible numerically.

Figure 1: Artist’s impression of an accreting neutron star. Material stripped from the companion star forms an accretion disk around the neutron star. Image credit: NASA / Goddard Space Flight Centre / Dana Berry.

The ability to observe the sky in X-rays using space-based instruments has led to the discovery of superbursts. These energetic outbursts recur on timescales of years and are believed to be driven by the unstable ignition of a carbon-enriched layer on a neutron star. To ignite a superburst, a carbon-enriched layer needs to contain a carbon mass fraction of roughly 20 percent. However, such a carbon-enriched layer is difficult to produce in most theoretical models. Besides requiring enough carbon, models for superbursts also require large ocean temperatures of roughly 600 million K. Such high temperatures are difficult to attain from standing heating models of neutron stars.

A study by Medin & Cumming (2011) suggests that the preferential freezing of heavier elements at the base of the ocean on an accreting neutron star can substantially enrich the ocean with lighter elements such as oxygen and carbon. At the base of the ocean, the increasing pressure from the continuous accretion of material onto the neutron star forces the preferential freezing of heavier elements. The separation of lighter elements from heavier elements releases energy and provides an additional source of heating for the ocean. After the preferential freeze-out of heavier elements, the remaining fluid becomes lighter than the fluid immediately above it and acts as a source of buoyancy which drives convective mixing of the ocean. Convection distributes the heat throughout the ocean in the form of a convective flux. The extra heat input can raise the temperature of the ocean up to the required ignition temperature of around 500 to 600 million K to produce a superburst.

In the study, a 300 million K ocean consisting of a mixture of iron (Z = 26) and selenium (Z = 34), and a mixture of oxygen (Z = 8) and selenium (Z = 34) is examined. At the base of the ocean, the preferential freezing of heavier elements enhances the abundances of lighter elements in the ocean. For example, a mixture of oxygen and selenium with initial 2 percent oxygen by mass can be enriched to almost 40 percent oxygen by mass. Although oxygen was chosen as the light element in this study, models with carbon (Z = 6) were also investigated and shown to yield similar enrichment results. The carbon mass fraction can be brought up by enrichment to the required ~20 percent for superburst ignition.

Figure 2: Phase diagram for crystallization of an iron/selenium mixture (top panel) and an oxygen/selenium mixture (bottom panel) in a 300 million K ocean on a neutron star. The stable liquid region of each phase diagram is labelled as “L”, the stable solid region(s) are labelled as “S” or “S1” and “S2”, and the unstable region is filled with plus symbols. Additionally, in each panel the composition at the top of the ocean is marked by a vertical dashed line, the ocean-crust boundary is marked by a horizontal dotted line, the composition of the liquid at the base of the ocean is marked by a filled square, and the composition of the solid(s) in the outer crust are marked by filled circles. Medin & Cumming (2011).

Figure 3: Thermal profile of an ocean on an accreting neutron star. The ocean is composed of a mixture of oxygen and selenium. The solid line represents the thermal profile when the convective flux (i.e. energy released at the base of the ocean from the separation of lighter elements from heavier elements) is included in the total heat flux. The dashed line represents the thermal profile when the convective flux is ignored (i.e. the total heat flux is due only to the heat emanating from the neutron star’s interior). Medin & Cumming (2011).

Reference:
Medin & Cumming, “Compositionally Driven Convection in the Oceans of Accreting Neutron Stars”, ApJ 730:97 (10pp), 2011 April 1.

Wednesday, July 23, 2014

Could it be a “Q-Star” instead of a Black Hole?

Compact objects fall under two categories - neutron stars or black holes. Neutron stars are the ultra-dense, compact remnant cores of massive stars. They are made almost entire of neutrons and have densities comparable to the density of an atomic nucleus. These neutrons are held together and kept from transmuting back into normal matter by the neutron star’s intense gravity which arises from its extraordinary compactness. A teaspoon of neutron star material would contain a mass of roughly a billion tons. The minimum and maximum mass possible for any neutron star is between ~0.1 and ~3 times the Sun’s mass. Below the minimum mass, the neutron star’s gravity is too weak to hold the star together and the star “decompresses” into normal matter. Above the maximum mass, the neutron star’s gravity becomes sufficiently strong to crush it into a black hole.

Figure 1: Artist’s impression of a neutron star whose intense gravity is lensing light from the background.

Nevertheless, the physics of matter at ultra-high densities remains poorly understood. Bahcall, Lynn & Selipsky (1990) propose that the same type of matter found in a neutron star could be stably confined by an alternative means other than gravity. Such a form of matter, though still considered ultra-dense, would have densities far below what is found in a neutron star. The outcome is that a compact object made of such a form of matter could exceed 3 times the Sun’s mass and would not collapse into a black hole under its own gravity since it is not as compact as a neutron star. These objects are termed “Q-stars”.

Theoretical models by Miller, Shahbaz & Nolan (1997) show Q-stars can be up to several times the Sun’s mass, far above the maximum mass for neutron stars. Furthermore, Q-stars that are several times the Sun’s mass can have radii less than 1.5 times the event horizon radius of a black hole of corresponding mass. Basically, a black hole’s event horizon is a non-physical boundary around a black hole, and within it, gravity is strong enough to keep even light from escaping. Since a black hole does not have a true surface, its event horizon could be regarded as its “surface”.

Figure 2: Radius of a Q-star plotted as a function of its mass. Miller, Shahbaz & Nolan (1997).

A non-rotating Q-star with 12 times the Sun’s mass can have a radius as small as ~52 km. In comparison, a black hole of the same mass would have an event horizon radius of 36 km. This difference is less than a factor of 1.5 and shows that a Q-star can be comparable in size to the event horizon of a black hole of corresponding mass. As a consequence, it may be difficult to observationally determine whether a high-mass compact object with several times the Sun’s mass is a black hole or a Q-star.

One possible method to distinguish a black hole from a Q-star would be to observe the accretion of material by the high-mass compact object. If the object were a Q-star, the accretion flow would eventually intersect the surface. If the accretion flow extends further inwards, closer than what would otherwise be the surface of the Q-star, it would be good evidence that the high-mass compact object is a black hole rather than a Q-star. An example of a known high-mass compact object that could turn out to be a Q-star is V404 Cygni - an object currently thought to be a black hole with ~12 times the Sun’s mass. Even so, one should be mindful that Q-stars are purely theoretical constructs and they may not exist at all.

References:
- Bahcall, Lynn & Selipsky, “New Models for Neutron Stars”, ApJ (1990) 362, 251.
- Miller, Shahbaz & Nolan, “Are Q-stars a serious threat for stellar-mass black hole candidates”, MNRAS (1990) 294: L25-L29.

Tuesday, July 22, 2014

Formation of Binary Giant Planets

Giant planets seem to be ubiquitous around Sun-like stars. Our Solar System has two giant planets - Jupiter and Saturn. Both planets are primarily composed of hydrogen and helium. Jupiter and Saturn have 318 and 95 times the mass of Earth, respectively. Beyond Saturn, the planets Uranus and Neptune are generally classified as “ice giants” because they have much smaller masses and differ considerably in composition compared to Jupiter and Saturn. The orbits of Jupiter and Saturn form a 5:2 orbital resonance. For every five times Jupiter circles the Sun, Saturn would circle the Sun twice. On the whole, the orbits of Jupiter and Saturn are stable over the entire age of our Solar System.

In a planetary system with two giant planets, such as our Solar System, energy and angular momentum are conserved between the two giant planets, and the planetary system is stable. Instability only occurs if the orbits of the two giant planets bring them very close to one another. Exoplanet discoveries over the years have revealed a remarkable diversity of planetary systems. A number of studies have shown that planetary systems with three or more giant planets tend to be unstable. For such a planetary system, perturbations by the additional giant planet(s) tend to destabilise the system.

Figure 1: Artist’s impression of a pair of binary giant planets.

Figure 2: Artist’s impression of a giant planet.

When a planetary system consisting of three or more giant planets is destabilised, it can lead to a number of interesting outcomes. Ochiai et al. (2014) show that gravitationally bounded pairs of giant planets (i.e. binary giant planets) can form via planet-planet scattering during the destabilisation of a planetary system with three giant planets. In their study, N-body simulations of planetary systems with three Jupiter-mass giant planets were performed. The N-body simulations show that as much as ~10 percent of the planetary systems result in the formation of binary giant planets.

During the destabilization of a planetary system with three giant planets, the possible outcomes are - ejection of a planet, planet-planet collision, planet-star collision, formation of a hot-Jupiter and formation of a pair of binary giant planets. A hot-Jupiter forms when a giant planet is thrown inwards to its star whereby planet-star tidal interactions can circularise the orbit of the giant planet into a close-in orbit around the star, leading to the formation of a hot-Jupiter. As for binary giant planets, such a pair could form when two giant planets pass sufficiently close to one another that enough tidal dissipation occurs between them to form a gravitationally bound pair.

In their N-body simulations of planetary systems with three giant planets, Ochiai et al. (2014) used four sets of 100 simulation runs corresponding to the four different initial stellarcentric semimajor axes - 1, 3, 5 and 10 AU for the innermost giant planet. In the nomenclature, “stellarcentric semimajor axis” refers to the average distance of the giant planet from its host star and 1 AU is a unit of measurement equal to the average Earth-Sun separation distance. For the two outer giant plants, their semimajor axes are, respectively, factors of 1.45 and 1.9 times the semimajor axis of the innermost giant planet. The four sets of 100 runs follow the evolution of the planetary system over a period of 10 million years.

The results from the 400 simulation runs show that the formation rate of binary giant planets is ~10 percent and nearly independent of the stellarcentric semimajor axis. Binary giant planets generally form near their initial orbits because the period when they form is normally during the early stages of orbital instability. Regardless of the initial stellarcentric semimajor axes, the distribution of the semimajor axes of the binary giant planets (i.e. average distance between the two giant planets in the binary) show a peak at 2 to 4 times the combined planetary radii of the two giant planets in the binary. Also, the 400 simulation runs show that ejection rates increase and collision rates decrease as stellarcentric semimajor axis increases.

Figure 3: Distribution of the semimajor axes of the binary giant planets obtained from the 400 simulation runs. For each pair of binary giant planets, the semimajor axis is expressed as a ratio to the combined planetary radii of the two giant planets in the binary. Ochiai et al. (2014).

Figure 4: Results obtained from the 400 simulation runs for the four different initial stellarcentric semimajor axes - 1, 3, 5 and 10 AU. The colours represent binary giant planets (red), planet-planet or planet-star collisions (light green), hot-Jupiters (blue), ejections (magenta), and three giant planets still remaining after 10 million years (light blue). Ochiai et al. (2014).

Binary giant planets are expected to be stable over the long-term. If the stellarcentric semimajor axis of a pair of binary giant planets is larger than ~0.3 AU, the system is stable for ~10 billion years, which is similar in duration to the main-sequence lifespan of a Sun-like star. Interestingly, binary giant planets can have moons with wide orbits that circumscribe both planets. A loosely bound moon around one of the two giant planets has a roughly 20 percent chance of surviving the formation process leading to a pair of binary giant planets. Additionally, binary giant planets can also capture large moons into orbit around them, much like how Neptune captured its large moon Triton. Current planet detection methods might be able to detect binary giant planets.

Reference:
Ochiai et al., “Extrasolar Binary Planets. I. Formation by Tidal Capture during Planet-Planet Scattering”, ApJ 790:92 (10pp), 2014 August 1

Monday, July 21, 2014

Kepler-421b: A Uranus-Sized Planet near the Snow-Line

“In future, children won’t perceive the stars as mere twinkling points of light: they’ll learn that each is a ‘Sun’, orbited by planets fully as interesting as those in our Solar System.”
- Martin Rees

A protoplanetary disk is a circumstellar disk of material around a young star in which the formation of planets occurs. The snow-line marks the distance from the central star where the protoplanetary disk becomes cool enough for volatiles such as water to condense into solid ice grains. By analysing publicly available data from NASA’s Kepler space telescope, Kipping et al. (2014) present the discovery of a cold transiting planet near the snow-line. This planet, identified as Kepler-421b, is the first of its kind to be discovered. It is similar in size to Uranus and it circles a star that is slightly cooler than the Sun in a nearly-circular orbit with an orbital period of 704.2 days. Kepler-421b is the longest period transiting planet discovered to date.

Figure 1: Artist’s impression of a Uranus-like planet with a large moon in orbit around it.

Figure 2: Transit light curve of Kepler-421b. Based on how much light it blocks when it passes in front its parent star, Kepler-421b is estimated to be ~4 times the Earth’s diameter, roughly the size of Uranus. Kipping et al. (2014).

“Finding Kepler-421b was a stroke of luck,” says lead author David Kipping of the Harvard-Smithsonian Center for Astrophysics (CfA). “The farther a planet is from its star, the less likely it is to transit the star from Earth’s point of view. It has to line up just right.” Kepler-421b is ~1.2 AU from its parent star. At that distance, the planet is closer to its parent star than Mars is from the Sun. Since its parents star is only ~40 percent as luminous as the Sun, Kepler-421b receives only ~64 percent of the insolation Mars gets from the Sun, or ~28 percent of the insolation Earth gets from the Sun. Kepler-421b receives the same amount of insolation as an object at ~2 AU from the Sun. If Kepler421b has a Uranus-like albedo, the planet’s effective temperature would be ~180 K. For comparison, Earth has a mean surface temperature of 288 K, or 15°C.

Assuming Kepler-421b has a Uranus-like composition (i.e. an ice giant), the planet probably formed at its current distance from its parent star (i.e. in situ formation). At that distance, it is cool enough for icy planetesimals to form in the protoplanetary disk, eventually leading to the creation of an ice giant. The large orbital period of 704.2 days means that transits of Kepler-421b are relatively infrequent. In fact, only two transits have been observed so far and those were sufficient to result in its initial detection. Unfortunately, the 3rd transit occurred in March 2014, after Kepler’s primary mission. Nevertheless, the 4th transit opportunity is in February 2016. Kepler-412b is the first known transiting Uranus-sized planet in a long-period orbit. Determining its mass and finding more of its kind would be the next logical steps. Finally, the large distance of Kepler421b from its parent star makes it an appealing target in the search for exomoons.

Reference:
Kipping et al. (2014), “Discovery of a Transiting Planet Near the Snow-Line”, arXiv:1407.4807 [astro-ph.EP]

Hot-Jupiters around Red Giant Stars

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

Figure 2: Artist’s impression of a gas giant planet.

After ~6 billion years or so, the Sun will start running out of hydrogen in its core and being to enter its post-main-sequence phase of evolution characterised by a large increase in its luminosity. All the planets circling the Sun will receive much greater insolation than they do now. Presently, Jupiter orbits the Sun at a distance of roughly 5 AU, where 1 AU is the average Earth-Sun separation distance. When the Sun enters post-main-sequence evolution, Jupiter might become so intensely irradiated that it becomes a “hot-Jupiter”.

This occurs because the Sun’s luminosity will increase by a factor of several thousand during two stages in its post-main-sequence evolution - the red giant branch (RGB) stage, followed by the asymptotic giant branch (AGB) stage. During the RGB stage, the Sun’s interior is characterised by an inert helium core surrounded by a hydrogen-burning shell (i.e. hydrogen fusing into helium). For the subsequent AGB stage, the Sun’s interior is characterised by an inert carbon core surrounded by a helium-burning shell (i.e. helium fusing into carbon), and a hydrogen-burning shell.

A study by Spiegel & Madhusudhan (2012) show that Jupiter’s atmosphere can be transiently heated to temperatures of up to ~1000 K or more when the Sun goes through its RGB and AGB stages. Many of the currently known Jupiter-mass planets in wide, several-AU orbits around Sun-like stars (i.e. stars between 1 to 3 times the Sun’s mass) will also experience such a temperature increase when their host stars evolve off the main-sequence. The authors term such planets “red giant hot-Jupiters” (RGHJs) to distinguish them from typical hot-Jupiters that circle in short-period, close-in orbits around main-sequence stars.

Figure 3: Artist’s impression of a gas giant planet.

Figure 4: Orbital separations where RGHJs can be found around a Sun-like star. The first rise in temperature corresponds to the RGB phase and the second rise corresponds to the AGB phase. Spiegel & Madhusudhan (2012).

Gas giant planets like Jupiter start out warm and gradually cool over time. However, the intense heating a RGHJ receives from its post-main-sequence host star could “reset” its evolutionary clock. As a result, RGHJs or post-RGHJs could appear younger than they actually are. Nevertheless, there is a major difference between RGHJs and typical hot-Jupiters around main-sequence stars. Typical hot-Jupiters trap some fraction of incident stellar irradiation to produce bulk heating (i.e. internal heating) on timescales spanning tens of millions of years, allowing these objects to settle into a quasi-steady thermal state. By comparison, RGHJs do not have the luxury of time since the RGB and AGB phases last nowhere as long. A RGHJ would not be intensely irradiated for a long enough time to produce any significant bulk heating, even though their outer layers might appear strongly heated.

When a star enters the RGB stage, and subsequently, the AGB stage, it will undergo a huge increase in mass loss. The star loses mass in the form of a continuous stellar wind streaming away from the star. Since stellar wind speeds are typically a few times a star’s surface escape velocity, the stellar wind from a post-main-sequence star would travel much slower compared to the stellar wind from a main-sequence star. This is because a post-main-sequence star has puffed up so greatly in size that its surface escape velocity is small.

In fact, measurements of the stellar wind speeds of AGB stars by Zuckerman & Dyck (1989) show velocities less than 40 km/s, with a clustering around 5 to 25 km/s. For comparison, stellar winds from main-sequence stars, such as the present-day Sun, have speeds of a few 100 km/s. A Jupiter-mass planet around a post-main-sequence star would have sufficient gravity to capture and accrete the slow stellar wind flowing pass it. The total accreted mass is estimated to be of order ~1/10,000th of the planet’s mass for a Jupiter-mass planet.

The stellar wind streaming away from a post-main-sequence star can impose enough drag to change the orbits of small objects circling the star. These objects could then impact the RGHJ and enhance the abundance of heavy elements in the planet’s atmosphere. Furthermore, a post-main-sequence can exhibit a high carbon-to-oxygen ratio due to carbon dredged-up from the star’s interior. A RGHJ accreting the stellar wind from such a star could acquire a carbon-rich atmosphere.

Figure 5: Artist’s impression of a gas giant planet.

 Figure 6: Artist’s impression of a gas giant planet with a few of its moons appearing as points of light.

The intense stellar irradiation experienced by a RGHJ means that wind speeds in its atmosphere are expected to be faster than on present-day Jupiter. However, wind speeds on a RGHJ would not be as strong as on a typical hot-Jupiter around a main-sequence star because a RGHJ is not tidally-locked, and consequently, would not have a large enough day-night temperature contrast to drive strong winds. Roughly estimating, the expected wind speeds for present-day Jupiter, a RGHJ and a typical hot-Jupiter are 40 m/s, ~100 m/s and ~1000 m/s respectively.

Changes in the incident stellar irradiation around an evolving post-main-sequence star can cause interesting changes in the atmospheric chemical properties of a RGHJ. The present-day Jupiter has an atmospheric temperature of 165 K at the 1 bar level. As the Sun’s luminosity increases considerably during its post-main-sequence phase, an important change for a soon-to-be RGHJ would be the enhancement of the H2O abundance in the planet’s atmosphere as it becomes warm enough (i.e. ~300 K) for water-ice to sublimate. The abundance of H2O drops slights for a brief period during the RGB stage when atmospheric temperatures on the RGHJ exceed ~600 K. At such temperatures, some of the oxygen in H2O becomes bounded in silicates. The same drop in H2O abundance might also occur when temperatures rise again during the AGB stage.

Figure 7: Post-main-sequence evolution of Jupiter’s equilibrium temperature and atmospheric chemical composition. Spiegel & Madhusudhan (2012).

Figure 8: Example spectra of Jupiter as a function of equilibrium temperature, as seen from a distance of 5 AU. Spiegel & Madhusudhan (2012).

References:
- Spiegel & Madhusudhan (2012), “Jupiter will become a hot Jupiter: Consequences of Post-Main-Sequence Stellar Evolution on Gas Giant Planets”, arXiv:1207.2770 [astro-ph.EP]
- Zuckerman & Dyck (1989), “Outflow Velocities from Carbon Stars”, Astronomy and Astrophysics, 209, 119-125

Sunday, July 20, 2014

Life and the Formation of Continents

On Earth, the presence of life plays a major role in determining the chemistry of the atmosphere and oceans. A study by D. Höning et al. (2014) suggests that the presence of life may play an even deeper role in influencing the planet’s evolution. In particular, the presence of life can enhance continental weathering rates, thereby increasing the rates at which sediments wash into and settle on the bottom of the oceans. These sedimentary layers hold within them a significant amount of water and hydrated minerals. Along convergent plate boundaries, the oceanic crust gets subducted into the Earth’s mantle, bringing along the water-rich sedimentary layers. As the subducting oceanic crust dives deeper, the increasing lithostatic pressure squeezes free water out from the sedimentary layers in a process known as shallow dewatering.

The enhanced continental weathering rates due to the presence of life would lead to thicker sedimentary layers, and consequently, increase the amount of water being subducted. In addition, the enhanced continental weathering rates would reduce the amount of shallow dewatering due to a greater abundance of low-permeability deposits such as clay-rich deposits in the sedimentary layers. These low-permeability deposits effectively ‘seal off’ water entailed in the sedimentary layers from being squeezed out, and in doing so, reduces the amount of shallow dewatering, allowing more water to be transported by subduction deeper into the Earth’s mantle. Water that is not squeezed out becomes bound in stable hydrated minerals as it is dragged further down by the subducting oceanic crust.

Figure 1: Artist’s impression of an Earth-like planet hosting a system of rings.

Figure 2: Schematic cartoon depicting Earth’s global water cycle, where water is represented by large and small dots, its path by black arrows, and movement of the oceanic plate by white arrows. Initial water uptake occurs within the submarine oceanic crust and sediments. Water loss first occurs after the subduction trench through dewatering, followed by the formation of the water-rich partial melt. The partial melt drives arc volcanism and continental crust formation. However, a fraction of the water contained in the subducting plate is regassed into the mantle. The water leaves the convecting mantle at mid oceanic ridges (MOR) as free volatiles or becomes part of the newly formed oceanic crust. D. Höning et al. (2014).

At a depth of roughly 100 km, the hydrated minerals brought down by the subducting oceanic crust become unstable and releases water into the surrounding mantle. This lowers the melting temperature of the surrounding mantle and leads to partial melting. Buoyancy drives the partial melt towards the surface, causing surface volcanism and the formation of new continental crust. The amount of newly formed continental crust is directly proportional to the amount of water released to produce partial melting. A larger amount of water driven down by the subduing ocean crust and released to form partial melting would enhance the rate of production of continental crust.

Also, not all the water in the form of hydrated minerals is released to form partial melting. Some of it continues deeper into the Earth’s mantle where it dissolves, hydrating the mantle. This lowers the effective viscosity of the mantle (i.e. makes the mantle more fluid), which has the effect of stabilizing plate tectonics. Process such as subduction, volcanic activity and the formation of new continental crust depends a lot on plate tectonics. If the mantle was dry, plate tectonics might not occur due to the high mantle viscosity. In a way, plate tectonics needs water to operate.

On Earth, the oceans cover 70 percent of the planet’s surface, with land covering the remaining 30 percent. This study by D. Höning et al. (2014) show that the presence of life does play an important role in the formation of continents on Earth and should be applicable to the evolution of other Earth-like planets as well. In the study, a model depicting Earth with present day continental weathering and erosion rates show that after roughly 4 billion years, the planet reaches a steady state with continental area covering 40 percent of the planet’s surface and an upper mantle water concentration of 300 parts per million (ppm). The same model is then ran with 60 percent of present day continental weathering and erosion rates, which one might expect on an abiotic Earth (i.e. a lifeless Earth). After 4 billion years, the planet attains a steady state with continents covering ~5 percent of its surface and an upper mantle water concentration of 40 ppm.

Figure 3: Artist’s impression of an Earth-like world. In this case, it is a moon of a gas giant planet.

Figure 4: Artist’s impression of an Earth-like planet. Image credit: Adrian Thomassen.

These findings suggest that the difference between a life-filled and a lifeless Earth can showup as a significant difference in the extent of continental coverage. On a biotic world (i.e. a life-filled planet), the presence of life enhances the formation of continents and stabilizes plate tectonics. In contrast, on an abiotic world (i.e. a lifeless planet), continental coverage is expected to be less, and the occurrence of plate tectonics would be less likely. If future studies support these notions, the detection of large continental coverage and/or plate tectonics on Earth-like exoplanets could serve as a form of biosignature in the search for life beyond Earth. “If we find a planet somewhere in the universe with a continental coverage similar to the Earth, it may be a good place to search for life,” said lead author of the study, Dennis Höning, a planetary scientist at the German Aerospace Centre’s Institute of Planetary Research in Berlin.

Reference:
D. Höning et al., “Biotic vs. abiotic Earth: A model for mantle hydration and continental coverage”, Planetary and Space Science 98 (2014) 5-13.

Saturday, July 19, 2014

Hellacious Superrotating Winds on Hot-Jupiters

Hot-Jupiters are a class of exoplanets that share similar characteristics to Jupiter (i.e. they are all gas giant planets), but have extraordinarily high surface temperatures because they orbit very close to their parents stars. On a hot-Jupiter, the intense heating on the planet’s dayside drives powerful winds that tear continually around the planet, transporting heat from the dayside and dumping it on the nightside. These hellacious winds whip around the planet from west to east, generating what is known as superrotation. The winds extend from the planet’s equator to latitudes of typically 20° to 60°. Because a hot-Jupiter orbits so close to its parent star, the planet is most likely tidally-locked and presents the same hemisphere towards its parent star all the time. Superrotation on a tidally-locked hot-Jupiter tends to produce an eastward displacement of the planet’s hottest region from the substellar point by up to 10° to 60° of longitude.

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

S. Faigler & T. Mazeh (2014) analysed the Kepler light-curves of four transiting hot-Jupiters - KOI-13b, HAT-P-7b, TrES-2b and Kepler-76b. The light-curves of these four planets show beaming, ellipsoidal and reflection/emission (BEER) phase modulations. As a hot-Jupiter circles its parent star, it gravitationally tugs at the star, causing the star to “wobble”. From an observer’s point of view, the star would appear to approach and recede in a periodic fashion as the hot-Jupiter orbits around it. In the BEER phase modulations, the beaming effect, also know as Doppler boasting, is caused by an increase (decrease) in the brightness of the parent star as it approaches (recedes from) the observer. The ellipsoidal effect is caused by the tidal distortion of the star by the hot-Jupiter. Both the beaming and ellipsoidal phase modulations are proportional to the hot-Jupiter’s mass. Finally, the reflection/emission phase modulations are the result of a combination of light reflected from the planet’s dayside (reflection) and thermal radiation that is re-emitted by the planet (emission).

The back and forth motion of a star induced by the gravitational tugging of an orbiting hot-Jupiter also causes the star’s spectrum to be blue-shifted (red-shifted) when the star is observed to approach (recede from) the observer. This results in a radial velocity signature that is proportional to the planet’s mass and it serves as an independent measure of the planet’s mass in addition to the BEER phase modulations. Radial velocity measurements available for the hot-Jupiters - HAT-P-7b, TrES-2b and Kepler-76b show that the planetary-mass derived from the beaming amplitude is noticeably larger than from the radial velocity measurements. Also, for all four hot-Jupiters - KOI-13b, HAT-P-7b, TrES-2b and Kepler-76b, the planetary-mass derived from the beaming amplitude is larger than from the ellipsoidal amplitude.

S. Faigler & T. Mazeh (2014) suggests that the apparent planetary-mass discrepancy is found to be caused by superrotation, whereby the eastward displacement of the planet’s hottest spot from the substellar point induces an angle shift in the planet’s reflection/emission phase modulation which “leaks” into the beaming modulation and artificially boosts its observed amplitude. As a consequence, the planetary-mass estimated from the beaming amplitude is somewhat larger than the real one. When the effect of superrotation is included in a modified “BEER model”, the apparent planetary-mass discrepancies disappear. This study shows that hot-Jupiter superrotation may be a rather common phenomenon that can be identified in the Kepler light-curves of hot-Jupiters that exhibit considerable BEER phase modulations. For each of the four hot-Jupiters in this study, the hottest spot is estimated to be displaced eastwards from the substellar point by 0.8° ± 0.9° for KOI-13b, 5.4° ± 1.5° for HAT-P-7b, 38° ± 18° for TrES-2b and 9.2° ± 1.3° Kepler-76b.

 
Figure 2: The four panels correspond to the four hot-Jupiters - KOI-13b, HAT-P-7b, TrES-2b and Kepler-76b. For each panel, the solid line represents the best-fit superrotating BEER model of the corresponding hot-Jupiter’s light-curve. The residuals are shown below each panel. On each panel, the dashed, dash-dot and dotted lines represent the shifted reflection/emission, beaming and ellipsoidal models, respectively. The vertical red dashed line marks the phase of maximum reflection/emission. Notice that the phase of maximum reflection/emission of each planet comes before phase 0.5. This is consistent a superrotation-induced eastward displacement of the planet’s hottest spot from the planet’s substellar point. S. Faigler & T. Mazeh (2014).

Reference:
S. Faigler & T. Mazeh (2014), “BEER analysis of Kepler and CoRoT light curves: II. Evidence for emission phase shift due to superrotation in four Kepler hot Jupiters”, arXiv:1407.2361 [astro-ph.EP]

Friday, July 18, 2014

Two Tight Pairs of Low-Mass Binary Brown Dwarfs

Figure 1: Artist’s impression of a brown dwarf. Heat from its warm interior “leaks” out through gaps in its cloud coverage. A cool brown dwarf would resemble Jupiter more than it would resemble a star.

Brown dwarfs are substellar objects that span the gap between the most massive planets and the least massive stars. Like stars, brown dwarfs can also come in pairs. Using a technique known as gravitational microlensing, Choi et al. (2013) reported the discovery of two pairs of very low-mass binary brown dwarfs identified as OGLE-2009-BLG-151 and OGLE-2011-BLG-0420. Gravitational microlensing is observed when the gravity of an intervening object (lens) magnifies the light from a background star (source). It happens as the intervening object crosses the line-of-sight between the observer and the background star.

OGLE-2009-BLG-151 was first detected by the Optical Gravitational Lensing Experiment (OGLE) group and then independently detected by the Microlensing Observations in Astrophysics (MOA) group in 2009, hence its alternate designation - “MOA-2009-BLG-232”. The other gravitational microlensing event, OGLE-2011-BLG-0420, was detected by the OGLE group in 2011. A number of ground-based telescopes also provided follow-up observations for both gravitational microlensing events.

Figure 2: Light-curves of the binary brown dwarf gravitational microlensing events OGLE-2009-BLG-151 and OGLE-2011-BLG-0420. Choi et al. (2013).

The light-curve of OGLE-2009-BLG-151 is characterised by two prominent spikes, consistent with a binary-lens model. Based on the light-curve, OGLE-2009-BLG-151 is inferred to be a tightly-bound pair of very low-mass brown dwarfs with masses 0.018 and 0.0075 times the Sun’s mass. The two brown dwarfs are projected to be 0.31 AU, or 46 million km apart from each other. That is equal to Mercury’s closest distance to the Sun. For comparison, the average Earth-Sun separation distance is 1 AU, or 149.6 million km.

As for OGLE-2011-BLG-0420, its light-curve appears smooth and symmetric. However, upon careful observations, the light-curve shows noticeable deviations indicative of a binary-lens model rather than a single-lens model. Similar to OGLE-2009-BLG-151, OGLE-2011-BLG-0420 consists of two very low-mass brown dwarfs in a tight binary system. The two brown dwarfs are 0.025 and 0.0094 times the Sun’s mass, and are spaced only 0.19 AU, or 28 million km apart. In fact, the two brown dwarfs of OGLE-2011-BLG-0420 are spaced as far apart as Jupiter’s outermost moons are from Jupiter.

The total system masses of OGLE-2009-BLG-151 and OGLE-2011-BLG-0420 are 0.025 and 0.034 times the Sun’s mass, respectively, placing them well below the hydrogen-burning limit of ~0.08 times the Sun’s mass. What makes the discovery of OGLE-2009-BLG-151 and OGLE-2011-BLG-0420 interesting is that the two systems have among the lowest total system masses known for brown dwarf binaries. Additionally, OGLE-2009-BLG-151 and OGLE-2011-BLG-0420 are also the tightest known brown dwarf binaries. The discovery of these two systems among the relatively small sample of gravitational microlensing events involving binary systems shows that tightly-bound pairs of very low-mass brown dwarfs are not uncommon.

Figure 3: Projected separation versus total system mass for a compilation of binaries. Grey circles indicate old field binaries, whereas blue squares indicate young (< 500 million year old) systems. The size of the symbols is proportional to the square root of the mass ratio (i.e. the ratio of the less massive component to the more massive component in the binary system). The red stars correspond to OGLE-2009-BLG-151 and OGLE-2011-BLG-0420. Choi et al. (2013).

Figure 4: Binding energy versus total system mass for the same binaries as shown in Figure 3. Choi et al. (2013).

Reference:
Choi et al. (2013), “Microlensing Discovery of a Population of Very Tight, Very Low-mass Binary Brown Dwarfs”, arXiv:1302.4169 [astro-ph.SR]

Thursday, July 17, 2014

Circumsubstellar Disk around a Young Brown Dwarf

Brown dwarfs are substellar objects that span the gap between the most massive planets and the least massive stars. These objects are believed to form in the same way stars do, and in their infancy, can possess disks of material from which planets might form. Hannah Broekhoven-Fiene et al. (2014) report on the discovery of a circumsubstellar disk of material around a young brown dwarf identified as KPNO Tau 3. The discovery was based on submillimeter observations of KPNO Tau 3 using the Submillimetre Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT). Submillimeter astronomy is a branch of observational astronomy that involves the part of the electromagnetic spectrum between the far-infrared and microwave wavebands.

Artist’s impression of how a planetary system around a brown dwarf might look like. Image credit: Drew Taylor.

KPNO Tau 3 is situated in the relatively nearby Taurus star-forming region, ~450 light-years away. The circumsubstellar disk detected around KPNO Tau 3 is estimated to contain ~130 Earth-masses worth of material, assuming a gas to dust ratio of 100:1. A planetary system consisting of a few sub-Earth-mass or Earth-mass planets might eventually coalesce out from this circumsubstellar disk. Furthermore, the detection of cold, ~20 K dust grains implies that a significant fraction of dust in the circumsubstellar disk is at a large enough distance from KPNO Tau 3 where the radiated energy from the young brown dwarf is too feeble to have sufficiently warmth the dust grains by much.

The presence of cold dust in the circumsubstellar disk around KPNO Tau 3 is consistent with the belief that brown dwarfs, at least some fraction of them, form in the same manner as low-mass stars. An alternate brown dwarf formation mechanism involves the ejection of a stellar embryo from its place of birth. The ejection process ‘starves’ the stellar embryo such that is no long able to accrete enough matter to form a full-fledge star and instead settles as a brown dwarf. A formation scenario like this would truncated the brown dwarf’s circumsubstellar disk and result in the absence of cold dust grains far from the brown dwarf.

Reference:
Hannah Broekhoven-Fiene et al. (2014), “The Disk around the Brown Dwarf KPNO Tau 3”, arXiv:1407.0700 [astro-ph.SR]

Wednesday, July 16, 2014

Swinging on Highly Eccentric Orbits

A number of exoplanets are known to orbit their host stars in wildly eccentric orbits. These exoplanets have “comet-like” orbits where they come close to their host stars and then recede far out. One example is HD 80606 b - a Jupiter-like planet which orbits its host star in a highly elongated orbit with eccentricity 0.9336. The planet’s distance from its host star varies from 0.03 to 0.88 AU, where 1 AU is the average Earth-Sun distance. At closest approach, HD 80606 b receives over 800 times more insolation than when it is farthest from its host star. Near closest approach, temperatures on HD 80606 b can rise from 800 K to 1500 K in a mere 6 hours. HD 80606 b attains a maximum orbital velocity of 227 km/s when it is closest to its host star. At that speed, a sufficiently large meteoroid barrelling into HD 80606 b could produce a truly spectacular meteor on the planet’s nightside.

Figure 1: A model of HD 80606 b showing the stormy response of the planet’s atmosphere right after closest approach to its host star. Credit: NASA/JPL-Caltech/G. Laughlin et al.

In the Solar System, all known planets go around the Sun in relatively circular orbits. Nonetheless, besides comets, a number of asteroids are known to orbit the Sun in highly eccentric orbits. These asteroids have “comet-like” orbits that bring them from a far out locale, to a close swing around the Sun, and out again. One notable asteroid in this category is 2006 HY51 - a 1.2 km asteroid in an extremely eccentric orbit around the Sun. 2006 HY51 has a remarkable orbital eccentricity of 0.9688. It comes within 0.081 AU of the Sun (1/4 of Mercury’s closest distance to the Sun) and recedes as far as 5.118 AU from the Sun (grazing Jupiter’s orbit).

Figure 2: Artist’s impression of an asteroid.

2006 HY51 goes around the Sun ever 1530 days, spending the vast majority of its time further than Earth is from the Sun. It is believed to be an asteroid with a stony composition and not an inactive comet. When 2006 HY51 is closest to the Sun, it receives nearly 4000 times more insolation from the Sun compared to when it is at its farthest. The insolation it receives at closest approach is 150 times more intense than the insolation Earth gets from the Sun. From 2006 HY51, the Sun at closest approach would appear over 12 times larger than it would from Earth. The orbital velocity of 2006 HY51 reaches almost 150 km/s when it is swinging by the Sun at closest approach. In comparison, Earth orbits the Sun with an average orbital speed of 29.8 km/s. A few asteroids such as 2005 HC4 and 2008 FF5 have highly eccentric orbits similar to 2006 HY51 that take them very near the Sun. However, both are much smaller than 2006 HY51 and do not swing out as far.

References:
- Moutou C. et al. (2009), “Photometric and spectroscopic detection of the primary transit of the 111-day-period planet HD 80606 b”, Astronomy and Astrophysics 498 (5): L5-L8
- Fossey S.J., Waldman I.P., and Kipping D.M. (2009), “Detection of a transit by the planetary companion of HD 80606”, Monthly Notices of the Royal Astronomical Society: Letters 396: L16-L20
- G. Laughlin et al., “Rapid heating of the atmosphere of an extrasolar planet”, Nature 457, 562-564 (29 January 2009)
- Yoonyoung Kim et al. (2014), “Physical Properties of Asteroids in Comet-like Orbits in Infrared Asteroid Survey Catalogs”, arXiv:1405.2989 [astro-ph.EP]