Possible Belt of Objects between Uranus and Neptune

Numerical simulations have revealed a region of space between the orbits of Uranus and Neptune where it is stable enough for a belt of small objects residing there to survive over the age of the Solar System without being ejected by gravitational interactions with the giant planets - Jupiter, Saturn, Uranus and Neptune. In the study, test particles were used to identify possible regions of stability. The results show that in the region between 24 to 27 AU from the Sun, about 0.3 percent of an initial population of test particles moving on low-eccentricity, low-inclination orbits could survive over the 4.5 billion year age of the Solar System.

Figure 1: A size comparison of Neptune and Earth.

 Figure 2: Orbits of Mars, Jupiter, Saturn, Uranus and Neptune shown to scale. Credit: Lunar and Planetary Institute.

Figure 3: Simulation results of test particles in the region between Uranus and Neptune. The test particles were initially on low-eccentricity, low-inclination orbits. The red points indicate test particles that survived the full 4.5 billion year integration, while the green points indicate test particles that did not survive. (Holman M.J., 1997)

Although members of this hypothetical belt of objects between Uranus and Neptune have yet to be detected, observations that have been carried out so far do not preclude its existence. Instead, the non-detections help constrain the size distribution of objects within this belt. At a very high level of confidence, there can no more than 7 objects larger than 640 km and 2,900 objects larger than 50 km. The total mass of this hypothetical belt is estimated to be comparable to the mass of the asteroid belt between Mars and Jupiter. Survey telescopes such as the proposed Large Synoptic Survey Telescope (LSST) may be able to detect member objects of this belt.

The discovery of objects in the 24 to 27 AU region between Uranus and Neptune would imply that processes happening during planet formation left material in this region undisturbed or that some subsequent process repopulated the region. Nevertheless, an absence of detectable objects would place more significant limits on the population of objects in this region and would suggest that processes such as planet migration have disrupted this fragile population.

Reference:

Holman M.J., “A possible long-lived belt of objects between Uranus and Neptune”, Nature, 387, 785-788 (19 June 1997)

An Isolated Giant

Very massive stars with ~100 times the Sun’s mass are exceedingly rare. It is still not known if such massive stars can form in isolation or only in star clusters. The centre region of the Milky Way galaxy hosts a unique environment where the process of star formation may differ from elsewhere in the galaxy. This makes the Galactic Centre an important test-bed for determining if very massive stars can indeed form in isolation. In the Galactic Centre are 3 very massive star clusters. The central Cluster surrounds the galaxy’s supermassive black hole, while 2 other massive star clusters, the Arches cluster and the Quintuplet cluster, are located no more than ~100 light years away.

Figure 1: Artist’s impression of the Arches cluster.

Images of the Galactic Centre acquired by the Hubble Space Telescope (HST) reveal a number of isolated massive stars located outside of the 3 known massive star clusters. Four scenarios that may explain the origin of these isolated massive stars are that these stars (1) were formed in isolation; (2) were formed within clusters that have already dispersed; (3) were ejected from one of the 3 known massive star clusters; (4) belong to clusters that have yet to be discovered.

WR 102ka is one of the most massive and most luminous stars known in the galaxy. It also happens to be one of those isolated massive stars in the Galactic Centre. The current mass of WR 102ka is estimated to be ~110 times the Sun’s mass. Since a massive star like WR 102ka has a very high mass loss rate, the initial mass of WR 102ka is believed to be ~150 times the Sun’s mass. WR 102ka blazes with ~3 million times the Sun’s luminosity. Its estimated age is ~2 million years. Very massive stars like WR 102ka live fast and die young.

Figure 2: Artist impression of a massive star.

 Figure 3: A massive star with ~20 times the Sun’s mass is shown next to a stack of lighter, Sun-like stars. For every one such massive star, there could be 500 to as many as 2000 smaller stars. Such a proportion of small to big stars is also expected for WR 102ka if it formed in a massive star cluster where it would be accompanied by a large number of less massive counterparts. Credit: NASA/JPL-Caltech.

In 2009, a team of astronomers used the European Southern Observatory’s (ESO) Very large Telescope (VLT) in Chile to study WR 102ka and its surroundings. If WR 102ka belongs in a star cluster that has yet to be discovered, simulations predict that in such a massive star cluster initially containing a single star with the mass of WR 102ka, there should be ~300 stars exceeding 20 times the Sun’s mass and ~10 of these stars may even exceed 100 times the Sun’s mass. Nevertheless, observations reveal no massive star cluster is associated with WR 102ka, even though such a cluster would have been clearly detectable in the observations.

It is also unlikely that WR 102ka formed within such a massive star cluster ~2 million years ago and that the cluster has since dispersed. This is because ~2 million years is too short a time for such a massive star cluster to disperse. For example, the Quintuplet cluster is an older massive star cluster with an estimated age of 3 to 5 million years and it is still very much intact.

It may be possible that WR 102ka was ejected from one of the 3 massive star clusters. There are a number of mechanisms that can eject a star from it natal cluster. Measurements show that WR 102ka has a radial velocity of 60 km/s. Even with a slow runaway velocity, say equal to its radial velocity of 60 km/s, WR 102ka could easily have arrived at its current location, far from any of the 3 massive star clusters in the Galactic Centre within its ~2 million year lifetime.

However, a number of observational evidences do not support an ejection scenario for WR 102ka. Firstly, it is unlikely that a very massive star like WR 102ka (possible the most massive star were it part of a cluster) can be ejected from a cluster while many less massive stars remain bound. Secondly, infrared observations using the Spitzer space telescope show that a bow shock in the direction of motion is absent around WR 102ka, while the same observations detected bow shocks around two other isolated massive stars in the Galactic Centre. For these two stars, one was ejected from the Central Cluster while the other was ejected from the Quintuplet cluster. Each star has a bow shock in its direction of motion as it ploughs through the interstellar medium.

Thirdly, images from the Spitzer space telescope and from the Wide-field Infrared Survey Explorer (WISE) show the presence of a dusty circumstellar nebula around WR 102ka. The nebula is probably stellar material thrown out by WR 102ka during previous evolutionary stages. The central position of WR 102ka within its nebula suggests that the star has remained at the same location.

All these observations show that WR 102ka, one of the most massive and most luminous stars in the galaxy, may have formed in isolation. In fact, the majority of isolated massive stars in the Galactic Centre do not display obvious bow shocks, suggesting that the massive star population in the Galactic Centre consists of stars formed in clusters, stars ejected from clusters and stars that formed in relative isolation.

Reference:

L. M. Oskinova et al. (2013), “One of the most massive stars in the Galaxy may have formed in isolation”, arXiv:1309.7651 [astro-ph.SR]

Spiral Structure around R Sculptoris

R Sculptoris is a dying red giant star located ~1,000 light-years from Earth. The star is surrounded by a detached shell of dust and gas that was created during a thermal pulse event where the star underwent a brief period of increased mass loss. A team of astronomers using the Atacama Large Millimeter/Submillimeter Array (ALMA) in the Atacama Desert of northern Chile discovered the presence of a spiral structure within the shell of material around R Sculptoris. The spiral structure extends from the central star outwards to the shell. “We’ve seen shells around this kind of star before, but this is the first time we’ve ever seen a spiral of material coming out from a star, together with a surrounding shell,” says Matthias Maercker, the lead author on the paper presenting the results.

Figure 1: A visualization of the spiral structure around R Sculptoris. Credit: ALMA (ESO/NAOJ/NRAO)

Towards the end of their lives, low-mass and intermediate-mass stars, such as the Sun, become red giant stars and start to lose a large amount of their mass. During the red giant phase of a star, it also periodically experiences thermal pulses. A thermal pulse occurs when a thin shell of helium surrounding the star’s core fuses the helium into carbon in an explosive fashion. The thermonuclear convulsion dumps a huge amount of energy into the star’s interior which causes the star to blast off a significant amount of material. For this reason, red giant stars like R Sculptoris are major contributors to the bulk of raw materials such as carbon and oxygen that are incorporated into the formation of future generations of stars and planets.

Figure 2: The detached shell and spiral structure observed in different velocity channels. Numbers in the top right corners indicate the velocity in km/s with respect to R Sculptoris. The detached shell is most pronounced at lower velocities while the spiral structure can be traced through all velocity channels. (M. Maercker et al., 2012)

 Figure 3: The green curve outlines the spiral structure around R Sculptoris. (M. Maercker et al., 2012)

Observational data combined with hydrodynamic simulations suggests that the shell of dust and gas around R Sculptoris was created when the star underwent a thermal pulse ~1,800 years ago, lasting for ~200 years. The spiral structure observed around R Sculptoris is most likely caused by the presence of a companion star that is shaping the stellar wind into a spiral pattern, like a rotating garden sprinkler. Extending from the central star out to the shell, the spiral structure can be followed over about 5 windings. Based on the spacing of the windings and the present-day expansion velocity of the stellar wind streaming from R Sculptoris, the companion star around R Sculptoris is estimated to have an orbital period of 350 years.

The amount of material present in the shell of dust and gas around R Sculptoris is estimated to have a mass of ~0.003 times the Sun’s mass (~1000 times the Earth’s mass). All these material was blasted outward at 50,000 kilometres per hour. Since the thermal pulse lasted for ~200 years, the mass-loss rate of R Sculptoris during the thermal pulse is ~5 Earth masses per year. This is about 30 times lower than the present-day mass-loss rate of ~0.15 Earth masses per year. “It’s a real challenge to describe theoretically all the observed details coming from ALMA, but our computer models show that we really are on the right track. ALMA is giving us new insight into what’s happening in these stars and what might happen to the Sun in a few billion years from now,” says Shazrene Mohamed, a co-author of the study.

Reference:

M. Maercker et al., “Unexpectedly large mass loss during the thermal pulse cycle of the red giant star R Sculptoris”, Nature 490, 232-234 (11 October 2012)

Moonlets in Saturn’s A Ring

Figure 1: An image of Saturn’s main rings taken by the Cassini spacecraft on 21 June 2004, a few days before the spacecraft entered orbit around Saturn. Credit: NASA/JPL/Space Science Institute.

Figure 2: An annotated version of an image of Saturn’s A Ring. This image was taken by the Cassini spacecraft on 9 May 2007, at a distance of approximately 1.1 million km from Saturn. Credit: NASA/JPL/Space Science Institute.

Saturn’s main rings are comprised primarily of water-ice particles ranging between ~1 cm and ~10 m in radius. In 2006, four ‘propeller’-shaped features were discovered in Saturn’s A Ring from images previous taken by the Cassini spacecraft. Simulations have show that these ‘propeller’-shaped features are formed by gravitational perturbations of ring particles from moonlets measuring tens to hundreds of metres in size.

Figure 3: Images of 4 ‘propeller’-shaped features discovered in Saturn’s A Ring. These images were taken by the Cassini spacecraft in 2004 and the nominal image resolution is 52 m per pixel. (M. S. Tiscareno et al., 2004)

These moonlets are embedded within the A Ring and are too small to be seen directly by the Cassini spacecraft. Instead, the Cassini spacecraft sees the ‘propeller’-shaped disturbances created by these moonlets. By 2008, ~150 features have been found in Saturn’s A Ring that are associated with moonlets embedded within the ring. About half of these features are sufficiently well resolved to reveal a characteristic ‘propeller’ shape.

As the moonlets orbit Saturn within the A Ring, they partially sweep up the ring particles around them to create the observed ‘propeller’-shaped features. However, they are not large enough to sweep clean their entire orbit around Saturn, unlike the moons Pan and Daphnis. Both these moons reside within gaps in Saturn’s A Ring. Pan orbits inside the 325 km wide Encke Gap, while Daphnis orbits inside the 42 km wide Keeler Gap.

The ‘propeller’-shaped features associated with the moonlets lie primarily in three ~1000 km wide belts in the middle section of Saturn’s A Ring, between 126,750 km and 132,000 km from Saturn’s centre. Also, relatively larger moonlets do exist and they can create bigger ‘propeller’-shaped features that can be tracked for years. On such feature is even nicknamed Bleriot, after a French aviator named Louis Bleriot. Bleriot is believed to be associated with a moonlet measuring ~1 km in size, significantly larger than many other moonlets.

Figure 4: In this image taken by the Cassini spacecraft, the central ‘propeller’ structure of Bleriot is estimated to be ~100 km in length. Credit: NASA/JPL-Caltech/Space Science Institute.

References:

- M. S. Tiscareno et al., “100-metre-diameter moonlets in Saturn’s A ring from observations of ‘propeller’ structures”, Nature 440, 648-650 (30 March 2006)

- M. Seiß, F. Spahn, M. Sremčević and H. Salo, “Structures induced by small moonlets in Saturn’s rings: Implications for the Cassini Mission”, Geophysical Research Letters Vol. 32 L11205 (2005)

- M. Sremčević et al., “A belt of moonlets in Saturn’s A ring”, Nature 449, 1019-1021 (25 October 2007)

- M. S. Tiscareno et al., “The Population of Propellers in Saturn’s A Ring”, Astronomical Journal, Volume 135, Page 1083-1091 (2008)

Ancient Nova Sell around Z Camelopardalis

Z Camelopardalis is a dwarf nova located at a distance of about 530 light years. It consists of a white dwarf accreting hydrogen-rich matter from a companion red dwarf star. The accreted matter forms an accretion disk around the white dwarf where instability in the accretion disk causes it to episodically dump much of itself onto the white dwarf. The hydrogen-rich matter gradually accumulates on the white dwarf’s surface. Each dump liberates a large amount of gravitational potential energy and causes Z Camelopardalis to brighten by up to a factor of 40. This process repeats itself every 20 days or so.

Figure 1: Artist’s concept of Z Camelopardalis, a stellar system featuring a white dwarf accreting hydrogen-rich matter from a companion star. Credit: NASA/JPL-Caltech.

It is predicted that white dwarfs like Z Camelopardalis will eventually accumulate sufficient hydrogen-rich matter to undergo classical nova eruptions. During a classical nova, the accumulated hydrogen-rich matter on the white dwarf ignites and fuses hydrogen into other heavier elements in a thermonuclear runaway process. The surface of the white dwarf explodes and creates clearly visible shells of ejected material in the aftermath. A classical nova is thousands of times more luminous than a dwarf nova.

In January 2004, an image of Z Camelopardalis and its surroundings was acquired by NASA’s Galaxy Evolution Explorer (GALEX) satellite, a space-based ultraviolet (UV) telescope. The image shows an arc of UV emitting material centred on, and located southwest (SW) of Z Camelopardalis. Linear nebulosities to the northeast (NE) and southeast (SE) are also visible. These features appear to extend out to a distance of up to ~2 light years around Z Camelopardalis and are believed to be part of a shell of ejected material that was formed when Z Camelopardalis underwent a classical nova eruption in the past.

Figure 2: Image of the dwarf nova Z Camelopardalis and its surroundings. Z Camelopardalis is circled. (M. Shara et al., 2012)

Observations of other classical novae show ejection velocities in the range 300 to 3000 km/s. Given the physical size of the classical nova ejecta around Z Camelopardalis, it allows the age range of the ejecta to be set at 2,400 to 240 years. A larger age is favoured due to the “snowplough effect” observed with other classical novae where the ejecta sweeps up the interstellar medium and decelerates. In fact, an age of more than 1,300 years is most likely for the ejecta around Z Camelopardalis. This means Z Camelopardalis underwent a classical nova eruption over 1,300 years ago. During such an event, Z Camelopardalis would have been one of the brightest stars in the sky for at least a few days. It is possible that historical observations of such an event might exist.

References:

- M. Shara et al., “An ancient nova shell around the dwarf nova Z Camelopardalis”, Nature 446, 159-162 (8 March 2007)

- M. Shara et al., “The Inter-Eruption Timescale of Classical Novae from Expansion of the Z Camelopardalis Shell”, ApJ 756 (2012) 107 arXiv:1205.3531 [astro-ph.SR]

A Pair of Very Young Brown Dwarfs

Brown dwarfs bridge the gap between the least massive stars (~0.075 Sun’s mass) and the most massive planets (~0.013 Sun’s mass). Although they form as stars do and resemble gas giant planets, a brown dwarf is neither star nor planet. In 2006, a team of astronomers announced the discovery of a pair of young brown dwarfs located ~1400 light years away in the Orion Nebula. The pair is identified as 2MASS J05352184-0546085 and it is the first example of an eclipsing binary system comprising two brown dwarfs. Here, the two brown dwarfs mutually eclipse each other as they orbit around their common centre of mass.

Figure 1: This artist’s impression shows an eclipsing binary system. As the two components orbit each other, they pass in front of one another and create periodic dips in their combined brightness. Credit: ESO/L. Calçada.

The higher mass ‘primary’ component has a mass of 0.054 ± 0.005 solar mass and a size of 0.669 ± 0.034 solar radii, while the lower mass ‘secondary’ component has a mass of 0.034 ± 0.003 solar mass and a size of 0.511 ± 0.026 solar radii. Since the Orion Nebula star cluster is extremely young, 2MASS J05352184-0546085 is believed to be no more than a few million years old. As such, the large radii of its two brown dwarfs (~5 times larger than older brown dwarfs) are consistent with theoretical predictions for young brown dwarfs in the earliest stages of gravitational contraction. In comparison, older brown dwarfs at ~1 billion years of age have sizes of ~0.1 solar radii.

Figure 2: Light curve of 2MASS J05352184-0546085. The ratio of eclipse depths provides a direct measure of the ratio of surface temperatures, with the deeper eclipse corresponding to the eclipse of the hotter ‘secondary’ component by the cooler ‘primary’ component. Both components orbit each other with an orbital period of 9.78 days. (Stassun et al., 2006)

Light curve data of 2MASS J05352184-0546085 indicates that the lower mass ‘secondary’ component has a higher temperature than the higher mass ‘primary’ component. The estimated temperatures of the ‘primary’ and ‘secondary’ components are 2,650 ± 100 K and 2,790 ± 105 K respectively. Finding that the higher mass ‘primary’ brown dwarf has a lower temperature than its ‘secondary’ companion is puzzling because theoretical models predict that a brown dwarf of a given mass will at all times be warmer than a lower mass brown dwarf of the same age.

Two explanations have been proposed that may explain the reversal of component temperatures with mass in 2MASS J05352184-0546085. The first is that the two brown dwarfs are mildly non-coeval, with the higher mass ‘primary’ component being ~0.5 million years younger than the ‘secondary’ component. The second explanation is that strong magnetic activity on the ‘primary’ component is suppressing the transport of energy from interior to surface and causing a lowering of its surface temperature.

References:

- Stassun, K. G., Mathieu, R. D., & Valenti, J. A., “Discovery of two young brown dwarfs in an eclipsing binary system”, Nature 440, 311-314 (16 March 2006)

- Stassun, K. G., Mathieu, R. D., & Valenti, J. A., “A Surprising Reversal of Temperatures in the Brown Dwarf Eclipsing Binary 2MASS J05352184-0546085”, ApJ 664: 1154-1166, 2007 August 1

An Extremely Cool White Dwarf Star

White dwarfs are dense stellar remnants that mark the final evolutionary state of intermediate-mass and low-mass stars. Right after formation, a white dwarf is very hot and will appear blue-white in colour. Since material in a white dwarf no longer undergoes fusion reactions like in a normal star, a white dwarf does not have a source of energy and will gradually cool as it radiates away its energy. As a white dwarf cools, the radiation emitted by it will shift from the hot blue-white colour of an O-type star (> 30,000 K) to the cool red colour of an M-type red dwarf star (< 4,000 K). Given sufficient time, a white dwarf will cool until it no longer emits significant heat or light, and become what is known as a black dwarf. Nonetheless, no black dwarfs are expected to exist in the present Universe because the length of time required for a white dwarf to cool to this state far exceeds the current age of the Universe.

Artist’s impression of a white dwarf shown to scale with the Earth.

In 1997, a white dwarf identified as WD 0346+246 was serendipitously discovered during re-examination of images taken for a survey to detect brown dwarfs in the Pleiades cluster. The images were acquired by the UK Schmidt Telescope at Siding Spring Observatory in New South Wales, Australia. WD 0346+246 is observed to have an estimated temperature of 3,900 K, making it one of coolest white dwarfs currently known. A subsequent estimate gives WD 0346+246 a slightly lower temperature of 3,750 K. At that temperature, it will have a colour temperature similar to a red dwarf star. Parallax measurements yield an estimated distance of approximately 90 light years for WD 0346+246. This places WD 0346+246 within the Sun’s neighbourhood.

WD 0346+246 also shows unusually large velocity components with respect to the disk of the Milky Way galaxy. This indicates WD 0346+246 is a member of the galactic halo and is currently passing through the Sun’s neighbourhood. Most stars in the Milky Way galaxy exist within a single plane known as the galactic plane. However, a population of stars inhabit a spherical region around the galaxy known as the galactic halo. Stars in the galactic halo tend to be much older and less enrich in elements heavier than hydrogen and helium than stars in the galactic disk. Assuming WD 0346+246 is a carbon white dwarf with a pure hydrogen atmosphere, the length of time required for it to cool to its current state is at least 10 billion years. The existence of WD 0346+246 suggests that a significant number of such cool white dwarfs reside in the solar neighbourhood and have yet to be identified.

References:

- Hambly, N. C., Smartt, S. J., & Hodgkin, S. T., “WD 0346+246: A Very Low Luminosity, Cool Degenerate in Taurus”, ApJ 489: L157-L160, 1997 November 10

- N. C. Hambly et al., “On the parallax of WD 03461246: a halo white dwarf candidate”, MNRAS (1999) 309 (4): L33-L36.

- P. Bergeron (2001), “The Halo White Dwarf WD 0346+246 Revisited”, arXiv:astro-ph/0105333

Searching for Far-Flung Objects

Trans-Neptunian objects (TNOs) are objects that orbit the Sun beyond Neptune. They mainly are found either as Kuiper belt objects (KBOs) or Oort cloud objects (OCOs). KBOs orbit the Sun within a region extending from the orbit of Neptune out to ~50 AU. Pluto is an example of a KBO and one of its largest members. Detecting KBOs, especially smaller bodies, is very challenging due to the huge distance and faintness of these objects. As a result, small KBOs with sizes on the order of a few km or less still elude detection even though they are expected to be far more numerous than larger KBOs. Beyond the KBOs are the OCOs. These objects comprise the Oort cloud, an enormous spherical cloud of icy bodies extending as far as 50,000 AU from the Sun, or nearly a quarter of the distance to Proxima Centauri, the nearest star to the Sun.

Figure 1: Artist’s concept of the Solar System in perspective. The scale bar is in astronomical units, with each set distance beyond 1 AU representing 10 times the previous distance. Credit: NASA/JPL-Caltech.

The Oort cloud is thought to be made of two separate regions: a spherical outer Oort cloud and a disc-shaped inner Oort cloud. No direct observations of the Oort cloud have yet been made. Nevertheless, a handful of presently known TNOs are possible members of the inner Oort cloud. One such object is Sedna - a large inner Oort cloud object about 1000 km in diameter. Sedna orbits around the Sun in an extremely elliptical orbit, coming as close as 76 AU from the Sun and receding out as far as 937 AU. Each orbit takes Sedna ~11,400 years. When Sedna was discovered in 2003, it was close to its minimum distance from the Sun. Given that an object like Sedna spends the vast majority of its time much further from the Sun; its discovery indicates there may be a large population of Sedna-sized objects out there.

Doressoundiram et al. (2013) propose a technique known as the serendipitous stellar occultation method to detect objects residing in the Kuiper belt and Oort cloud. It involves detecting such objects when they happen to pass in front of a star and occult it. This is well suited for detecting tiny and invisible objects that are otherwise too faint for current telescopes to photograph. In fact, the smallest KBO ever found is less than 1 km in size and it was found by the Hubble Space Telescope using just such a technique. Although the object is too faint to be photographed by the Hubble Space Telescope, it was detected when it passed in front of a background star, temporarily disrupting the starlight.

Figure 2: Artist’s impression of a small KBO in the process of occulting a background star. Credit: NASA, ESA and G. Bacon (STScI).

The instrument envisioned for the detection of TNOs by serendipitous stellar occultations is known as a Fast Multi-Object Photometer (F-MOP). It consists of several 3×3 bundles of optical fibres feeding a CCD. Each 3×3 fibre bundle would monitor a single star and the whole instrument would monitor a relatively wide field of view containing 50 to 100 stars that are suitable for detecting TNOs by stellar occultations. The F-MOP instrument is mounted onto an 8 m diameter telescope. The stars would be monitored with a cadence of at least 50 Hz to catch the occultation events. This is because, for a small TNO, an occultation event is expected to last for only about one second.

Figure 3: Synthetic light curve profile of an occultation event involving a 4 km object at 1300 AU. (Doressoundiram et al., 2013)

Consider a target star with an angular size of 0.015 milli-arcseconds. For such a star, the F-MOP instrument can detect an occultation event caused by an object 300 m or larger at 50 AU, or an object 840 m or larger at 200 AU. If the angular size of the target star is smaller, then the minimum size of an object that can be detected by occultation is also smaller. For an object at 5000 AU, it can be detected if it is 30 km or larger and occults a target star with an angular size of 0.014 milli-arcseconds. If the angular size of the target star is reduced by a factor of ten to 0.0014 milli-arcseconds, an object 4 km or larger can be detected if it occults. Hence, serendipitous stellar occultation can serve as a powerful and unique tool to probe the Kuiper belt and Oort cloud.

Reference:

Doressoundiram et al., “Ground-based exploration of the outer Solar system by serendipitous stellar occultations”, MNRAS 428, 2661-2667 (2013)

Ozone Layers on Alien Earths

Figure 1: Artist’s impression of an Earth-like planet. Credit: Scott Richard.

In the search for Earth-like planets around other stars, the presence of life on these worlds can be determined by looking for various biomarker gases in the planet’s atmosphere. Two promising biomarker gases are oxygen, which is produced almost entirely by photosynthesis on Earth, and ozone, which is produced in the Earth’s stratosphere when ultraviolet (UV) light splits oxygen into individual oxygen atoms where they then combine with other oxygen molecules to form ozone. Ozone is a good indicator of photosynthetic life because even a small amount of atmospheric oxygen can result in a significant concentration of ozone.

Segura et al. (2003) developed photochemical and radiative/convective atmospheric models of Earth-like planets around 3 different types of stars: F-type, G-type (Sun) and K-type. This is to see how an Earth-like planet might differ from a planet circling our Sun. The models assume a present-day Earth-analogue planet with an atmospheric oxygen concentration at the present atmospheric level (PAL). Also, the planet’s distance from its host star is scaled according to the star’s luminosity such that the planet’s average surface temperature is 288 K, which is similar to the average surface temperature of present-day Earth.

Figure 2: A comparison of the types of stars from M-type to B-type.

On Earth, the stratosphere is a layer of the atmosphere where temperature increases as altitude increases, reaching a peak at 45 km altitude. This heating is caused by absorption of UV flux by ozone. In the atmospheric models by Segura et al. (2003), stratospheric temperatures are warmest for the planet around an F-type star and coolest for the planet around a K-type star (Figure 3). This is because an F-type star is hotter and produces a higher UV flux, resulting in the greater absorption of UV flux by ozone. Furthermore, the higher UV flux also causes the planet around the F-type star to have a thicker ozone layer (Figure 4).

Figure 3: Atmospheric temperature profile results for an Earth-like planet with 1 PAL of oxygen and around different types of stars. (Segura et al., 2003)

Figure 4: Atmospheric ozone concentration results for an Earth-like planet with 1 PAL of oxygen and around different types of stars. (Segura et al., 2003)

The amount of UV flux reaching a planet’s surface is of concern due to its ability to damage cells and even DNA. On Earth, most damage caused by UV radiation is from UV-A (315-400 nm) and UV-B (280-315 nm), with UV-B being somewhat more dangerous. Fortunately, almost no UV-C (< 280 nm) penetrates the Earth’s atmosphere. UV-C is the most energetic and most dangerous form of UV radiation. The atmospheric models for an Earth-like planet with 1 PAL of oxygen show, in the UV range of 200 to 400 nm, the surface of the planet around a K-type star receives 0.44 times the UV flux that the Earth’s surface receives, while the surface of the planet around an F-type star receives 1.61 times the UV flux.

For UV-B alone, the planet around a K-type star receives 0.43 times the amount Earth receives, while the planet around an F-type star receives 0.68 times the amount. This shows that planets around K-type and F-type stars exhibit significantly better UV protection than Earth at 1 PAL of oxygen, despite an F-type star being hotter than our Sun. For a K-type star, it is simply due to it being cooler and emitting less UV flux. For the F-type star, it is a consequence of its higher UV flux which results in the formation of a much thicker ozone layer. Planets around all 3 types of stars also show negligible amounts of UV-C reaching the surface.

Figure 5: Incoming and surface UV fluxes for Earth-like planets with different UV fluxes and orbiting around different stars. (Segura et al., 2003)

Figure 6: Normalized surface UV dose rates relative to present-day Earth for skin cancer (erythema) and DNA damage on Earth-like planets with different UV fluxes and orbiting around different stars. (Segura et al., 2003)

The atmospheric models by Segura et al. (2003) can be extended to oxygen levels lower than 1 PAL. For 0.1 PAL of oxygen and in the UV range of 200 to 400 nm, the planet around a K-type star receives 0.45 times the UV flux that present-day Earth receives, while the planet around an F-type star receives 1.62 times the UV flux. For UV-B alone, the values are 0.61 (planet around K-type star) and 0.85 (planet around F-type star) times the UV flux that present-day Earth gets. It is also clear that below ~0.01 PAL of oxygen, the ozone layer becomes too thin to provide significant UV shielding regardless of the type of star the planet circles around.

Reference:

Segura et al., “Ozone Concentrations and Ultraviolet Fluxes on Earth-Like Planets Around Other Stars”, Astrobiology Volume 3, Number 4, 2003

An Extremely Low Density Super-Earth

The main objective of NASA’s Kepler space telescope is to determine the frequency of Earth-like planets around Sun-like stars. Analysis of the data collected by Kepler has already revealed a wide diversity of Earth mass and super-Earth mass planets. These planets have masses ranging from ~1 to ~10 times Earth’s mass. Using data collected by Kepler together with follow-up observations by ground-based telescopes, Aviv Ofir et al. (2013) report the discovery of a gas giant and a very low density super-Earth around the star Kepler-87. The gas giant is identified as Kepler-87b while the low density super-Earth is identified as Kepler-87c.

Kepler detects planets by looking for the tell-tale signature when a planet passes in front of its parent star and blocks some of the star’s light. It allows the planet’s size to be directly measured since a larger planet will block more of the star’s light than a smaller planet. Because Kepler-87b is a gas giant, it has a deeper transit depth (i.e. blocks more of the light from its star) than Kepler-87c. Based on the transit depths, Kepler-87b measures 13.49 ± 0.55 Earth radii and Kepler-87c measures 6.14 ± 0.29 Earth radii.

Figure 1: Artist’s impression of a super-Earth with a substantial outer envelop of hydrogen and helium. Credit: Paul A. Kempton.

Figure 2: Transit light curves of the planets and planet candidates in the Kepler-87 system. From top to bottom: transiting exoplanets Kepler-87b and Kepler-87c, and planetary candidates KOI 1574.03 and KOI 1574.04. Above each light curve are the model residuals. (Aviv Ofir et al., 2013)

Figure 3: Mass-radius relation for all known planets with masses below 30 times the mass of Earth. It is obvious that Kepler-87c occupies a unique position on this parameter space as the lowest density planet for its super-Earth mass range. (Aviv Ofir et al., 2013)

Kepler-87, the parent star of Kepler-87b and Kepler-87c, is somewhat more luminous than the Sun. This is because, although Kepler-87 has about the same surface temperature as the Sun, it spans 1.82 times the Sun’s diameter. Kepler-87 also has an estimated age of 7 to 8 billion years.  The time intervals between successive transits show orbital periods of 114 days for Kepler-87b and 191 days for Kepler-87c. As such, the distances of both planets from Kepler-87 can be determined and the estimated equilibrium temperatures of Kepler-87b and Kepler-87c are 205°C and 130°C, respectively. In addition, the data from Kepler also shows the presence of two more short-period super-Earth sized planet candidates - KOI 1574.03 and KOI 1574.04. Unlike Kepler-87c, both planet candidates are each less than 2 Earth radii in size. Follow-up observations will be required to confirm their planetary status.

Using the transit timing variations (TTVs) method, the masses of Kepler-87b and Kepler-87c are determined to be 324.2 ± 8.8 Earth masses and 6.4 ± 0.8 Earth masses, respectively. The TTVs method involves measuring the timing variations between consecutive transits when Kepler-87b and Kepler-87c transit their parent star. These timing variations are caused by gravitational perturbations on the outer planet (Kepler-87c) by the inner planet (Kepler-87b) and vice versa. With their sizes and masses known, the bulk densities of Kepler-87b and Kepler-87c are determined to be 73 percent and 15 percent the density of liquid water, respectively. As a result, Kepler-87b is a rather typical Jupiter mass planet with a Saturn-like density while Kepler-87c is an extremely low density planet in the super-Earth mass regime.

Kepler-87c is the lowest density planet currently known in the super-Earth mass regime and its density is even similar to the least dense hot-Jupiters such as TrES-4b and WASP-17b. However, unlike these hot-Jupiters, Kepler-87c is not strongly irradiated by its parent star. Kepler-87c’s equilibrium temperature of 130°C pales in comparison with the temperatures of hot-Jupiters that can get up to ~2000°C. The extremely low density of Kepler-87c is consistent with ~20 percent or more of its mass in the form of a low density gaseous outer envelope comprised of hydrogen and helium. For comparison, the super-Earth Kepler-10b has an estimated bulk density of 8.8 times the density of liquid water. Such a high density implies that Kepler-10b must be predominantly made of rock and iron. The almost two orders of magnitude density difference between Kepler-87c and Kepler-10b demonstrates the enormous compositional variety for planets with masses between ~1 to ~10 times Earth’s mass.

Reference:

Aviv Ofir et al. (2013), “An Independent Planet Search In The Kepler Dataset. II. An extremely low-density super-Earth mass planet around Kepler-87”, arXiv:1310.2064 [astro-ph.EP]