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HD 100453

Editor’s note: Astrobites is a graduate-student-run organization that digests astrophysical literature for undergraduate students. As part of the partnership between the AAS and astrobites, we occasionally repost astrobites content here at AAS Nova. We hope you enjoy this post from astrobites; the original can be viewed at astrobites.org!

Title: The Orbit of the Companion to HD 100453A: Binary-Driven Spiral Arms in a Protoplanetary Disk
Author: Kevin Wagner, Ruobing Dong, Patrick Sheehan, et al.
First Author’s Institution: Steward Observatory, University of Arizona
Status: Published in ApJ

Today’s paper combines a wide range of data sets — spanning the radio to near-infrared — and analysis techniques — orbit fitting and hydrodynamic simulations — to connect a binary companion to intriguing features seen in the protoplanetary disk around the primary star.

Using the Spectro-Polarimetric High Contrast Exoplanet Research instrument (SPHERE) in 2015, astronomers discovered a two-armed spiral structure in the disk around HD 100453 A (see the cover image). This structure is very different from the gaps seen in images of protoplanetary disks from the Atacama Large Millimeter/submillimeter Array (ALMA) such as HL Tau and TW Hya. The spiral arms seen in the disk around HD 100453 A and two other disks (SAO 206462 and MWC 758) could be caused by a massive companion (planet or star) orbiting outside the disk or processes within the disk such as self-gravity or dead zones. The HD 100453 system is unique in that it has a known M-dwarf companion of about 0.2 solar masses (HD 100453 B). The authors of this paper show that this companion is the cause of the spiral arms seen in the disk, without invoking other driving mechanisms.

The first step in connecting the companion star to the spiral arms of the disk was to determine the companion’s orbit. The authors used six observations with SPHERE and the Nasmyth Adaptive Optics System and Near-Infrared Imager (NACO) cameras on the Very Large Telescope and the Magellan Adaptive Optics system taken over a span of 14 years. The authors took care to minimize systematic errors in the astrometry which could be introduced by errors in the plate scale, orientation of the telescope (which direction is north on the camera) and using a coronagraph. With six pairs of separations and position angles, the authors were able to fit the orbital parameters of the companion M-dwarf. Most important for determining the origin of the spiral arms are the semi-major axis (109 ± 9 au), eccentricity (0.17 ± 0.07), and inclination (32.5 ± 6.5 degrees). This semi-major axis and eccentricity are consistent with the companion truncating the disk at 40 AU, much smaller than a typical disk around a single star.

Since the mutual inclination between the companion and disk has a significant effect on the evolution of the system, the authors needed to determine the inclination of the protoplanetary disk. They used publicly available ALMA observations of carbon monoxide in the disk. Fitting a simple smooth disk profile to the Keplerian orbits of the gas gave a disk inclination of 28 degrees, consistent within 1σ with the inclination of the companion.

Image of the HD 100453 system (top) compared with hydrodynamic and radiative transfer simulation viewed from an inclination of 30 degrees (bottom). [Wagner et al. 2018]

The final step was to run a hydrodynamic simulation of the entire system, including the effects of the companion. The authors evolved an initially smooth disk for 100 orbits of the companion and produced synthetic observations using a radiative transfer code. A sample of simulation results is shown in the figure to the right. The separation of the spiral arms, their pitch angle, and the locations where they sprout from the central ring are all well reproduced by the model. The authors note that the disks in their simulations are ~30% larger than the observed disk, though they suspect this is likely due to the short amount of time for which the simulations were run (100 companion orbits) compared to the age of the system (~12,000 companion orbits). If the computer time were available to run the simulation longer, the authors speculate that the companion would truncate the disk further. The amount of truncation also depends on the scale height and viscosity of the disk which are likely not exactly correct in their models.

The agreement between the inclination of the M-dwarf companion and the disk suggest that the entire system formed from a single cloud rather than the companion later becoming bound to the primary star (and its disk). The likely inclination of HD 100453 A (determined by comparing the observed rotational velocity of the star with stars of similar mass) is also consistent with the disk and companion star. This rules out a possible scenario where the companion formed separately but torqued the disk to share its inclination while leaving the star untouched.

While the spiral arms in the HD 100453 A disk are clearly driven by HD 100453 B, it is hard to make the same conclusion for the other two disks hosting “grand design” spiral arms. This and other studies suggest where a companion could be located with respect to the spiral arms in those systems, but previous searches for such a companion in these systems have found nothing, setting strict limits on companion mass (or brightness). As always, more work is needed to determine the origin of the spiral arms in SAO 206462 and MWC 758.

About the author, Samuel Factor:

Sam Factor is a 3rd year Ph.D. candidate at The University of Texas at Austin studying direct imaging of extrasolar planets and low mass binary stars. He uses an interferometric post processing technique to allow the detection of companions below the diffraction limit of the telescope.

binary black holes

Editor’s note: Astrobites is a graduate-student-run organization that digests astrophysical literature for undergraduate students. As part of the partnership between the AAS and astrobites, we occasionally repost astrobites content here at AAS Nova. We hope you enjoy this post from astrobites; the original can be viewed at astrobites.org!

Title: Merger of Multiple Accreting Black Holes Concordant with Gravitational Wave Events
Author: Hiromichi Tagawa & Masayuki Umemura
First Author’s Institution: Eötvös University, Hungary; National Astronomical Observatory of Japan
Status: Accepted to ApJ

A love story that begins with a chance encounter between strangers might sound romantic, but for black holes, the resulting attachment is often inescapable. Today’s astrobite explores one of the many theory-oriented publications written in the wake of LIGO‘s six gravitational-wave (GW) events. We’ll see how the authors explored the ramifications of throwing several unassociated black hole (BH) “strangers” into the mix (it’s complicated — accretion, three-body interactions, and more are at play in mediating mergers), and what it could mean in the context of recent GW discoveries.

Though the LIGO and Virgo detectors have been on hiatus since last fall (the start of a year-long break between observing runs O2 and O3), the world of astrophysics continues to be bombarded with new GW results informed by O1 and O2 data. In October, for instance, the LIGO-Virgo team announced the detection of GWs from a binary neutron star merger (GW170817) accompanied by a gamma-ray burst (GRB 170817A). The timing couldn’t have been more impeccable: the event, which was glimpsed in LIGO, Virgo, and electromagnetic observations, occurred just days before the conclusion of O2. Just like that, the era of multi-messenger astronomy had finally begun.

GW observations of BH mergers yield some information about the properties of the objects themselves, but the question of how unassociated BHs end up close enough to merge (and what that environment looks like) remains unanswered. In today’s featured paper, the authors go about exploring these issues using N-body simulations of multiple-black-hole systems in gas-dense environments. Their simulations are sophisticated (post-Newtonian), with detailed general relativity and gas dynamics being taken into account.

Usually, BH mergers are simulated with binary evolution in mind; that is, systems with associated BHs are considered. In contrast, this team’s plan was to simulate the behavior of five unassociated accreting BHs in several gas-dense environments (we’ll see later why that helps unassociated holes come together) in order to determine what initial parameters could yield LIGO-like mergers. This means that initial BH masses were comparable to LIGO’s ~30 M☉ component BHs (the authors simulate 20, 25, and 30 M☉ equal-mass systems). Gas number density varied between 102 and 1010 cm-3, though the total amount of gas in the simulation stayed constant at 105 M☉.

Figure 1: A plot of the masses of the closest BHs (m2 vs. m1) right before they finally merge. Blue points are 20, black are 25, and red are 30 M☉. The component masses from the three most massive LIGO mergers (along with errors) are shown in the boxes. [Tagawa & Umemura 2018]

Figure 2 shows how various parameters change over the course of one simulation. In the beginning, BHs get closer due to gas dynamical friction: if a massive object is moving through a sea of particles (like a dense gas cloud), the small components get pulled gravitationally into the wake of the larger one, causing it to lose energy. The final binary merger is mediated by loss of energy through GW radiation. Between these two periods, the unassociated BHs become well acquainted, with interloping BHs taking the place of one of the binary components (twice!) and wreaking havoc on the system. The addition of accretion to the model is enlightening, too. Each BH gains ~10 M☉ through accretion near the beginning of the three-body interactions, but that quickly abates before the binary merger (in Figure 2 (c) and (a), a rapid increase in velocity during merging causes a significant drop in accretion rate). This is an interesting detail, as less gas accretion around the merger may cause electromagnetic counterparts to be dimmer than expected.

Figure 2: Several properties of black-hole mergers as a function of time. Panel (a) shows accretion rate, (b) shows mass, (c) shows velocity, and (d) shows distance between the closest two black holes. The simulation ends when the first two BHs merge. [Tagawa & Umemura 2018]

Through a thorough argument, the authors conclude that active galactic nuclei, or AGN, are the most likely environments for LIGO-esque mergers to take place. In short, this required estimating the expected merger time in both an AGN system and a giant molecular cloud (their estimate was between 30 and 100 Myr). Because of the importance of dynamical friction in causing the mergers, these timescales were possible only in environments with gas density > 106 cm-3, according to the simulations. This information, along with estimated event rates informed by LIGO detections, led the authors to conclude that AGN with high gas density provided the most fertile environments for unassociated BH strangers to merge.

The assumption of evenly distributed gas and the lack of a central, massive BH make these simulations imperfect. However, the scientific team’s ability to simulate multiple unassociated BHs is vital in expanding upon the classical model of systems with already-associated binaries. Further LIGO-Virgo detections will help us understand the environments in which the mergers occur in much greater detail. Still, these simulations are already incredible for elucidating the complicated dynamics of mergers with only a few GW event detections.

About the author, Thankful Cromartie:

I am a graduate student at the University of Virginia and completed my B.S. in Physics at UNC-Chapel Hill. As a member of the NANOGrav collaboration, my research focuses on millisecond pulsars and how we can use them as precise tools for detecting nanohertz-frequency gravitational waves. Additionally, I use the world’s largest radio telescopes to search for new millisecond pulsars. Outside of research, I enjoy video games, exploring the mountains, traveling to music festivals, and yoga.

warm Jupiter

Editor’s note: Astrobites is a graduate-student-run organization that digests astrophysical literature for undergraduate students. As part of the partnership between the AAS and astrobites, we occasionally repost astrobites content here at AAS Nova. We hope you enjoy this post from astrobites; the original can be viewed at astrobites.org!

Title: Models of Warm Jupiter Atmospheres: Observable Signatures of Obliquity
Author: Emily Rauscher
First Author’s Institution: University of Michigan
Status: Published in ApJ

Observing exoplanets is challenging! So how can we ever imagine learning something about their seasons? Over the past decade, astronomers have made extensive progress in understanding the atmospheres of hot Jupiters, including weather detections. But hot Jupiters are tidally locked and therefore experience no seasons. For observations of seasons, we need to push outwards to planets on longer orbits, where tidal interaction with the star is minimal. But a longer orbit means cooler planets. Cooler planets emit less thermal radiation, making them far dimmer than hot Jupiters, with blackbody spectra that peak at longer infrared wavelengths. This population of “warm Jupiters,” or Jupiter-sized planets with temperatures between 500–1000K, are out of reach for current telescopes. But with the James Webb Space Telescope’s 6.5-m mirror and its ability to observe out to longer wavelengths than current telescopes, astronomers will soon be studying the atmospheres of this new population of exoplanets!

How do these warm Jupiters differ from tidally locked hot Jupiters? Gravitational interactions between a star and a planet on a close orbit — like a hot Jupiter — will slow the rotation of the planet to the point where its rotation is the same as its orbital period. As a consequence, one side of the planet always faces its star, while the other side is never illuminated. Tidal locking also circularizes the orbit of a planet and removes any rotational tilt (obliquity). We therefore know the rotational period, eccentricity, and obliquity of a hot Jupiter without any required analysis. Warm Jupiters, on the other hand, are less affected by significant tidal effects, which means we have no intrinsic knowledge of these parameters. The author addresses the obliquity part of this problem in today’s astrobite by posing the following question: can we detect and determine the obliquity of a warm Jupiter and, in doing so, finally observe seasons on an exoplanet?

Wait! Time Out! Obliquities, Rotational Tilt, Seasons?

Figure 1: Earth’s rotational or axis tilt is the reason for our season(s). The hemisphere tilted towards the Sun experiences summer while the opposite hemisphere experiences winter. Spring and Fall occur when neither hemisphere is tilted towards the Sun leading to equal heating. [Golden Guide to Weather from St. Martin’s Press]

Besides being a word that scores you 23 points in Scrabble, the obliquity, or rotational tilt, of a planet controls the length and strength of that planet’s seasons. Figure 1 illustrates how Earth’s obliquity of 23 degrees creates seasonal changes over the course of an orbit. Summer or winter in one hemisphere depends on whether our rotational axis is pointing towards or away from the Sun, respectively. Now imagine Earth with no tilt. With no tilt, we wouldn’t have seasons. But with larger tilt, our seasons would be more extreme.

OK, Got It! Let’s Build a Planet!

The author creates a hypothetical warm Jupiter that has all the same properties of Jupiter, including same radius, mass, and rotational period. But instead of orbiting the Sun once every 5 years, this planet orbits a Sun-like star every 10 days, giving it a temperature of about 900 K. Using a global circulation model (GCM), the author simulates the atmosphere of this warm Jupiter at varying obliquities. What seasons look like on this warm Jupiter is plotted in Figure 2 for obliquities of 30 degrees (top panel), 60 degrees (middle panel), and 90 degrees (bottom panel). The fast rotational period (10 hours) of this planet compared to the orbital period of 10 days causes the atmosphere to smear out most of the day/night temperature contrast, allowing the author to average the temperature over longitude (east–west direction). The larger obliquity correlates with longer and more extreme seasons at higher latitudes (north–south direction). For obliquities greater than 60 degrees, the poles of the warm Jupiter become hotter than the equator, leading to larger temperature contrasts than the 30-degree (Earth-like obliquity) model.

Figure 2: Map of the longitudinally averaged temperature as a function of latitude and time over one orbit. Top panel is a warm Jupiter with 30 degrees obliquity, middle panel is 60 degrees obliquity, and bottom panel has 90 degrees obliquity. The black dashed line represents the location of the subsolar point over time. [Rauscher et al. 2017]

The Planet Now Has Seasons, Let’s “Observe” It

The paper first analyzes the phase curves of these hypothetical warm Jupiters. A phase curve is the light curve of a planet as it orbits around its star. At different points in its orbit, the planet will emit more or less light depending on what fraction of the day side we observe. As the planet has a 10-day orbit, the author notes that this would require continuous observations with JWST for those 10 days. From these phase curve models, the author noticed a degeneracy between the obliquity of a warm Jupiter and its viewing orientation. Figure 3 shows that a planet with the same obliquity can appear very differently depending on from the angle at which we observe it. By summing up the total flux of this planet at different locations in its orbit, we can create phase-curve observations. However, phase curves only provide a 1D total flux map of the planet. Even with the same obliquity, we will observe different amounts of flux simply due to the viewing angle. Phase curves alone don’t provide enough information to measure obliquity and viewing angle independently.

Figure 3: Models of the warm Jupiter at viewed with different orientations. The top panel shows orientations of the planet if we were to observe the planet directly above the equator. The second panel shows the same planet at obliquities 30, 60, and 90 degrees as in the top panel, but twisting the planet towards our line of sight by half of the obliquity value. The bottom panel twists the planet even more towards our line of sight where our viewing angle is equal to that of the planet’s obliquity. For example, the bottom right image twists the planet 90 degrees from its original orientation in the top right image. This adds a complication to the problem, we now have a degeneracy between obliquity and the viewing angle or orientation of the planet to our line of sight. A movie of this figure can be found here. [Rauscher et al. 2017]

In order to break this degeneracy, we will need more than just a measurement of the planet’s total flux. Eclipse mapping might be the solution, as it provides a 2D spatial map of the planet’s dayside. Figure 4 (from Majeau et al. 2012) illustrates the concept behind eclipse mapping. As the planet passes behind its star, slices of the planet are hidden over time corresponding to the shape of the eclipse. Rauscher concludes that by studying the shape of this eclipse, we will gain sufficient information to distinguish between the obliquity and viewing angles of the planet. And JWST should have high enough precision to detect these different shapes. 

Figure 4: The concept of eclipse mapping. As the planet passes behind its star, slices of the planet map to the slopes of the secondary eclipse. Combining this information with the eclipse depth should help observers distinguish the direction from which we are viewing the planet. [Majeau et al. 2012]

That said, this paper explores obliquities on a warm Jupiter assuming a known eccentricity and rotation rate. The reality is that these parameters will be unknown when observing an actual warm Jupiter. How this will affect the presented observations is currently being explored. This paper does stress that these unknowns will not wipe out our ability to measure obliquity; instead they will just make the data a little more “interesting” to analyze. With JWST, the future does appear to be hot, or uh, bright for warm Jupiters and exo-seasons!

About the author, Jessica Roberts:

I am a graduate student at the University of Colorado, Boulder, where I study extra-solar planets. My research is currently focused on understanding the atmospheres of the extremely low-mass low-density super-puffs. Out of the office, you will probably find me running, cross-stitching, or playing with my dog.

GW170817

Editor’s note: Astrobites is a graduate-student-run organization that digests astrophysical literature for undergraduate students. As part of the partnership between the AAS and astrobites, we occasionally repost astrobites content here at AAS Nova. We hope you enjoy this post from astrobites; the original can be viewed at astrobites.org!

Title: Brightening X-Ray Emission from GW170817/GRB 170817A: Further Evidence for an Outflow
Authors: John J. Ruan, Melania Nynka, Daryl Haggard, Vicky Kalogera, and Phil Evans
First Author’s Institution: McGill University
Status: Published in ApJL

In August 2017, gravitational waves from a binary neutron star merger (GW170817) were detected for the first time ever by LIGO and Virgo. Also detected were (deep breath recommended) — gamma-ray, X-ray, ultraviolet, optical, infrared, and radio waves — all from the same source. The era of multi-messenger astronomy was thus kicked off in spectacular fashion. The coincident short gamma-ray burst (GRB 170817A) confirmed that merging neutron stars are indeed one of the progenitors of short GRBs. Read this astrobite for more details about these observations.

Today’s bite further illustrates why multi-messengers are so very important and exciting. It is not entirely clear what happened after the merger, and complementary information from different channels can help construct an accurate picture. The optical signal, for example, may carry slightly different information about the source than the radio waves do. So, if we manage to detect both, we end up with an additional clue about whatever produced them! Here, we will discuss the X-ray observations of GW170817/GRB 170817A.

A Slight Delay…

When the Chandra X-ray observatory followed up the optical counterpart of GW170817 at ~2 days post-burst, no X-ray counterpart was detected. The X-rays did show up eventually, at ~9 days post-burst. This delayed brightening is quite unusual — standard GRB afterglows show dimming instead. Additional observations at ~15 and 16 days also showed no significant dimming. Then, because of sky proximity of the source to the Sun, sensitive X-ray monitoring was not possible until December 2017.

Getting Brighter!

Today’s paper presents Chandra observations at 109.2 days post-burst, immediately after constraints from the Sun were lifted. The X-ray flux of GRB 170817A has brightened significantly, with a count rate ~4 times larger than the previous detection at 15.6 days. Figure 1 shows the X-ray image at 109.2 days post-burst, alongside the one at 15.6 days post-burst, clearly showing the source get brighter. The spectra can be found in Figure 2 in the paper.

Figure 1: Chandra X-ray image of GRB 170817A at 15.6 days post-burst (left), compared with the image at 109.2 days post-burst (middle), along with the difference image (right). The host galaxy NGC 4993 and two other sources in the field can also be seen. [Ruan et al. 2018]

Multi-Messenger Magic

As it turns out, radio observations of this source have also reported monotonically brightening emission. Comparing the X-ray fluxes to radio reveals that they are brightening at a similar rate. This points to a common origin for these two messengers!

But what does all this mean for the post-merger scenario? Remember, what happened after the merger is not a settled question, but it’s really interesting if we want to study physics under some pretty extreme conditions. Let’s explore three post-merger narratives and how well they can predict the X-ray and radio light curves that were observed. You can also read more about GRBs, jets and afterglows here, here and here before diving in.

  • Model 1: A uniform (“top-hat”) jet, seen off-axis.
    Radio light curves do not support this model since the brightening occurs at a slower rate than predicted. The same goes for the X-ray observations.
  • Model 2: A structured jet, seen off-axis.
    A structured jet that breaks out of the surrounding material (ejected during the merger) can predict X-ray and radio light curves that are a good match. When the jet is observed off-axis, a slowly brightening afterglow is seen, coming from increasingly relativistic material closer to the jet axis.
  • Model 3: Afterglow of an outflow.
    In this model, an observable jet that breaks out is not required or favored. An outflow that injects energy into the shock continuously can produce a slow, monotonic rise of the afterglow emission. The outflow can either be a cocoon of material shocked by the jet, or a high-velocity tail of the dynamical ejecta from the merger.

Figure 2: Chandra X-ray light curve of GRB 170817A (black points), along with predicted X-ray light curves. The gray region is the period over which observations were not possible because of Sun constraints. [Ruan et al. 2018]

Figure 2 shows the Chandra X-ray light curve of GRB 170817A, along with predictions from structured-jet, cocoon and ejecta-outflow models. Note that the early-time X-ray flux is underpredicted. However, these models were fitted to the radio light curve, which does not cover early times. A complete fit would need to include the X-ray data as well.

Takeaways

The X-ray observations disfavor simple top-hat jets and support the scenario where both the X-ray and radio emissions are the afterglow of an outflow or structured jet. Continued monitoring will provide even more information for constraining post-merger models. Stay tuned!

About the author, Sanjana Curtis:

I’m a grad student at North Carolina State University. I’m interested in extreme astrophysical events like core-collapse supernovae and compact object mergers.

M dwarf

Editor’s note: Astrobites is a graduate-student-run organization that digests astrophysical literature for undergraduate students. As part of the partnership between the AAS and astrobites, we occasionally repost astrobites content here at AAS Nova. We hope you enjoy this post from astrobites; the original can be viewed at astrobites.org!

Title: Reduced Diversity of Life Around Proxima Centauri and TRAPPIST-1
Authors: Manasvi Lingam and Abraham Loeb
First Author’s Institution: Harvard Smithsonian Center for Astrophysics
Status: Published in ApJL

Are we alone in the universe? This question still baffles astronomers (along with the rest of humanity) today, despite the confirmed discovery of thousands of exoplanets. Sure, there may be a plethora of other space pebbles out there, but do any of them actually host life? And if they do, could it ever compare to the variety of species we see on Earth, from human beings to tardigrades?

Unfortunately, we don’t have the answers to these questions yet, but today’s authors take a step in the right direction. By creating a model to measure potential biodiversity, they explore which stars are most likely to host planets capable of supporting complex life. These findings may help future exoplanet habitability studies point in the right direction.

Why Is This Important?

Many habitability studies focus largely on M dwarfs, which readily reveal orbiting planets due to their low mass. Planets have indeed been found in the habitable zone (HZ) of these stars. But just because a planet resides in the HZ doesn’t mean that it’s exactly welcoming to any life forms. You see, the HZ of these M dwarfs doesn’t extend very far from the star itself. This means that any planets finding refuge in an M dwarf’s HZ are exposed to strong stellar winds, which can strip a planet of its atmosphere and destroy any chances for life.

Today’s authors postulate that the maximum timespan a planet can hold onto its atmosphere largely determines the biodiversity of life on that planet. If a planet has had a longer time to allow for speciation — the development of new species via evolution — then a more diverse biosphere would be present. This translates into a greater likelihood for intelligent life and increased chances for detection via habitability surveys.

The Process

The model used in this paper is quite simple, based only on the exponential growth of species richness, which is simply the number of species present at a given time. This growth is described by what the authors call the characteristic timescale of species diversification; i.e., the rate at which new species develop. By assuming that there are about 107 eukaryotic species present on Earth, the authors model the increased diversification of the Cambrian period and the time it took for life to emerge (known as abiogenesis). Since the model predicts these events nicely for our home planet, they next proceed into the exoplanetary realm.

First, they set constraints on the length of time an exoplanet is habitable. We would take an upper limit to be the lifetime of the host star (and therefore a function of mass) if we didn’t know better. But this is science, so of course there are other factors at play. Cue in that pesky stellar wind! The authors calculate just how long an atmosphere could withstand this phenomenon before being obliterated. They use this interval as the maximum timescale over which life could diversify (if less than the lifetime of the star).

With this time limit, they model the potential number of species on an exoplanet as a function of time. They find that the likelihood for life to develop is largely dependent on the mass of the host star. This relationship is shown in Figure 1, along with the relationship between mass and biological clock.

exoplanetary evolution timescales

Figure 1: (Left) the maximum timescale for exoplanetary evolution as a function of stellar mass. The top, solid line marks the solar lifetime. The dashed line marks the accepted timescale of eukaryogenesis on Earth, while the dotted line marks the abiogenesis timescale. (Right) the peak species diversity as a function of mass. The blue dashed line marks the current microbial species on Earth; the blue dotted line marks the current number of eukaryotic species. The black dashed and dotted lines denote the max number of species attainable for Proxima b and TRAPPIST-1 planets, respectively. [Lingam et al. 2017]

Results That May Excite You and Let You Down All at Once

According to the model, stars falling within the range of 0.28 to 1.74 solar masses are the only ones that could host life long enough for eukaryotic species to emerge. While K dwarfs and most G dwarfs fit these criteria, only higher mass M dwarfs make the cut.

Interestingly, a subset of this range (stars between 0.38 and 1 solar mass) could actually host life more complex than ours, since the timescale for speciation is even longer than that on Earth. Specifically, the maximum time for which a planet could support life occurs when the host star is 0.55 solar masses. This corresponds to a K-dwarf with a lifetime 6 times that of the Sun. So it seems like K dwarfs are the way to go if we want to find extraterrestrials.

On the downside, the model shows that stars with less than 0.17 solar masses would host planets incapable of supporting life. The generally accepted timescale of abiogenesis on Earth is about 200 Myr. If we assume that this timeframe applies to other worlds, then these low mass stars would have their atmospheres ripped away long before any semblance of life arose from the primordial soup. Sure, the abiogenesis timescale on other worlds may be different than our own, but hey, it’s the only comparison we’ve got.

Proxima Centauri, the closest star to our solar system (besides the Sun) and host to Proxima b. [2MASS/UMass/IPAC-Caltech/NASA/NSF]

Here’s the kicker. Some of the most infamous exoplanets, like Proxima b and the TRAPPIST-1 planets, orbit stars with very low masses — Proxima Centauri weighs in at only 0.12 solar masses and TRAPPIST-1 at a measly 0.08 solar masses. In light of this study, the exoplanets belonging to these star systems are very unlikely to host life, and if they do, it would be minimal at best.

There is still hope, however, for other well-known exoplanets like Kepler-186f and Kepler-1229b, whose host stars sit right in the sweet spot at 0.54 solar masses. This paper shows that the characteristics of the host star are crucial to the biodiversity of its planets. It also gives us a better idea of where to look for life in the universe. When we search for aliens, or more scientifically, habitable planets, we may want to direct our resources towards K and G type stars.

About the author, Lauren Sgro:

I am a PhD student at the University of Georgia and, as boring as it may sound, I study dust. This includes debris disk stars and other types of strange, dusty star systems. Despite the all-consuming nature of graduate school, I enjoy doing yoga and occasionally hiking up a mountain.

Gliese 1214b

Editor’s note: Astrobites is a graduate-student-run organization that digests astrophysical literature for undergraduate students. As part of the partnership between the AAS and astrobites, we occasionally repost astrobites content here at AAS Nova. We hope you enjoy this post from astrobites; the original can be viewed at astrobites.org!

Title: The Formation of Mini-Neptunes
Authors: Julia Venturini and Ravit Helled
First Author’s Institution: University of Zurich, Switzerland
Status: Published in ApJL

To be a master chef, one must have an incredible amount of culinary expertise and creativity to create a wide variety of dishes. To be a computational astrophysicist building planets in simulations, it feels a lot more like going to the store and buying a box of pancake mix. Instead of buying all the ingredients, you start with a pre-made mixture. From there, all you have to do is just add water!

Figure 1: Pancake mix for Mini-Neptunes. To serve, just add a gaseous atmosphere (between 10 and 25% of the planet’s total mass). [Amazon]

Simulating planets can be quite similar. In place of the pancake mix, you can start out with a ball of rock. From there, all you have to do is just add gas (see Figure 1). When making pancakes, one thing to be careful about is how much water to add. If you add hardly any water, you will just end up with the same dried-out bowl you started out with. If you add way too much water, you will end up with pancake soup.

Unlike in cooking, there’s no recipe to follow when building a planet; you do not get to decide how much gas to add! The size of the atmosphere a planet accumulates depends on its own properties and the properties of the surrounding protoplanetary disk in which it forms. Mini-Neptunes are the “perfect pancake” of planets — they ended up with just the right amount of gas. Any less and they would have stayed Earth-sized dried-out rock; any more and they would have grown to Jupiter-sized gaseous soup.

In today’s paper, Julia Venturini and Ravit Helled explore which planet and disk conditions are best for building planets that have just the right-sized atmosphere to be classified as mini-Neptunes — the most common type of exoplanet (see Figure 2), even though there are none in our solar system.

known transiting planet sizes

Figure 2: Histogram showing the frequency of known planets of different sizes. Mini-Neptunes are the most common, followed by the slightly smaller super-Earths. [NASA Ames/W. Stenzel]

Too Much Water for Your Pancakes

Mini-Neptunes are defined to be planets that are less than 10 times the Earth’s mass and have a heavy, hydrogen-dominated atmosphere that makes up between 10 and 25% of the planet’s total mass. (Most planets that are at least 1.6 times the Earth’s radius in size are likely mini-Neptunes or bigger. Any smaller planet is more likely to have a lighter atmosphere and be classified as a super-Earth or an Earth.) With mini-Neptunes being so common, we would naturally expect a wide variety of conditions to be favorable for forming them.

However, it seems easy for planets to avoid becoming mini-Neptunes. The smallest planets have so little mass that they struggle to accumulate a significant amount of gas. For comparison, the Earth’s atmosphere only makes up <0.0001% of its total mass. On the other end, bigger planets whose atmospheres reach the mass of their rocky cores become unstable and undergo what is called “runaway gas accretion” in which they can easily grow to Jupiter-sized (318 Earth masses) or larger.

This would not be a problem if mini-Neptunes could just stop growing once their atmospheres reached a fractional mass of 10–25%. Unfortunately, as long as there is a disk from which they can accrete gas, they will keep growing. This would be like if you poured the right amount of water into your bowl of pancake mix, but then were forced to keep pouring water until you ran out! Thus, the only way a planet can end up a mini-Neptune is if the gas in the disk dissipates at just the right time for the planet to have accreted just the right amount of gas for its atmosphere.

Simulating Planetary Evolution

The authors simulate the growth of a planetary core that starts out 100 times less massive than the Earth in the presence of a depleting gaseous disk enriched with small rocky material. In the first stage, only the rocky core grows. In the second stage, the planet becomes massive enough to gravitationally attract a sizable amount of gas — at which point both the core and the atmosphere continue to grow. They then repeat this process with different planet and disk conditions.

Forming Mini-Neptunes with Pebbles

Let’s start with one of the cases where the authors grow the planet with cm-sized pebbles:

Here, the planet is placed at 5 AU with a transparent atmosphere that is just hydrogen and helium (non-enriched). The planet’s core then grows very quickly and bursts past the upper limit of 10 Earth masses in a little over 1 Myr (Case A: dashed blue lines in Figure 3). Before 2 Myr, the atmosphere’s fractional mass already exceeds the upper limit of 25% and the planet is well on its way to growing to Jupiter-sized within 3 Myr — the average lifetime of a protoplanetary disk. With these conditions, mini-Neptunes should never form.

Let’s move the planet out to 20 AU. Being farther away from the star slows down the growth of the core (since there is less material to accrete at these distances, while the planet also orbits much slower). Unfortunately, it also slows down the growth of the atmosphere for the same reasons (Case B: dashed purple lines in Figure 3), again preventing mini-Neptunes from forming.

To speed up the atmosphere’s growth without speeding up the core’s growth, the authors enrich the atmosphere with water (which is more realistic). This makes the atmosphere heavier, allowing it to gravitationally contract faster — freeing up space to accumulate more gas into the planet’s atmospheric zone. With these conditions, the planet reaches the mini-Neptune phase, but too early! Before 2 Myr, the planet has once again accumulated too much of an atmosphere (Case C: solid purple lines in Figure 3).

simulations of planet formation

Figure 3. Left: Planet mass evolution over time for planets that do not become mini-Neptunes. The black components of each line refer to when each planet crosses the mini-Neptune stage. Right: Atmosphere fractional mass as the planet grows. Each line switches from thick to thin after 3 Myr (the average disk lifetime). [Venturini & Helled 2017]

To slow down the atmosphere’s growth, the authors make the atmosphere more opaque. With a high opacity, it is harder for heat to escape. Since the atmosphere cannot cool as quickly, it cannot contract — making it harder to accumulate more gas. As a result, the planet will become and stay a mini-Neptune from 2 to 4 Myr, right when the disk is expected to dissipate (Case D: solid purple lines in Figure 4)! With these optimal conditions, Venturini and Helled find that 40% of protoplanetary disks should be able to form mini-Neptunes depending on the exact lifetime of the disk.

simulations of mini-Neptune formation

Figure 4: Case D forms a mini-Neptune, as the average disk dissipates at just the right time (vertical red line at 3 Myr) for the planet to end up with just the right amount of gas (and total mass). [Venturini & Helled 2017]

Summary

Besides finding that mini-Neptunes can form from pebbles at 20 AU as described above, the authors also find that mini-Neptunes can form from km-sized planetesimals closer in at 5 AU in 83% of systems. Between these two types of accreted material, mini-Neptunes should be able to form easily in a wide range of locations for a variety of atmospheric properties. With these results, it is no surprise mini-Neptunes are so common in the “Kepler sector” of the galaxy.

About the author, Michael Hammer:

I am a 3rd-year graduate student at the University of Arizona, where I am working with Kaitlin Kratter on simulating planets, vortices, and other phenomena in protoplanetary disks. I am from Queens, NYC; but I’m not Spider-Man…

Boyajian's star

Editor’s note: Astrobites is a graduate-student-run organization that digests astrophysical literature for undergraduate students. As part of the partnership between the AAS and astrobites, we occasionally repost astrobites content here at AAS Nova. We hope you enjoy this post from astrobites; the original can be viewed at astrobites.org!

Title: The First Post-Kepler Brightness Dips of KIC 8462852
Authors: Tabetha S. Boyajian et al.
First Author’s Institution: West Virginia University
Status: Published in ApJL

We’ve posted before about Boyajian’s star, one of the great unsolved mysteries of the Kepler mission. Discovered by citizen scientists in 2015, this star has everything: deeply bizarre (and bizarrely deep) dips in flux, a hundred-year fade, intermittent brightening spells. Since Kepler, investigations of this star have been hampered by the lack of new data — it’s hard to tell, for example, whether the crazy flux dips repeat if you don’t stare at the star continuously, because you might have just missed them.

All of that changed in 2016, when Dr. Boyajian began monitoring her namesake star with the Las Cumbres Observatory Global Telescope Network (LCOGT). LCOGT has two completely independent telescopes in the northern hemisphere — one in Hawaii and one in the Canary Islands — so on May 18, 2017, when both telescopes reported that Boyajian’s star was dimming anew, Dr. Boyajian could immediately rule out instrumental effects as the cause.

That dimming turned out to be the first night of a very interesting summer for Boyajian’s star. Since May, the star has dimmed four separate times. In today’s paper, Dr. Boyajian presents the new data, and offers, for the first time, a hint at a solution to the mystery.

The four dimming events of May–September 2017, named the “Elsie family.” The y-axis represents the amount of light coming from Boyajian’s Star relative to its ordinary state, and the x-axis represents time in days. Each color represents a different telescope in the Las Cumbres network — their Texas observatory, in green, came online in November 2017. [Boyajian et al. 2018]

The Case of the Elsie Family

The first thing to notice about the four dips of the Elsie family are their wonderful names. “Elsie” comes from the initials “L.C.” of “light curve.” “Celeste” (inspired by the initials “C.L.”) is so named because it’s Elsie’s opposite — instead of dimming rapidly and then brightening slowly, it dimmed slowly and brightened quickly. “Skara Brae” is named after a neolithic town in Scotland, unexpectedly unearthed by a passing storm; “Angkor” after the great abandoned Cambodian city, obscured by forest for hundreds of years, but ultimately uncovered.

It’s a mark of the exceptional and inspiring level of public engagement in this research that these events were named at all, let alone so loftily — most astronomers are happy to stick with catalog numbers and Julian dates. (The dimming events observed in Boyajian’s star during the Kepler mission were given names like D1540, for comparison.) But this project owes everything to its citizen scientists. Not only did they discover the star in the first place, but they also crowd-funded the Las Cumbres observations that revealed the Elsie family.  

The second thing to notice is that all four dips are of similar depth (the star dims to ~98% of its ordinary brightness), but drastically different shapes. In other words, Boyajian’s star looks no more like a regular old exoplanet-hosting star than it did at the end of the Kepler mission, four years ago. Skara Brae bears some resemblance to one of the dips observed by Kepler, but we won’t know if it’s truly a repeat of that earlier dip until we’ve watched it for much longer and looked for further repeats.

An Answer…

The most important thing the Las Cumbres observations tell us, though, isn’t about the number or the shape of the new dips. It’s about their color — or rather, how they appear when viewed through filters of different colors. Behold, the first color information we have about the dips of Boyajian’s Star:

The “Elsie” event as observed in bandpasses of three different colors, by two of the Las Cumbres telescopes (plotted as circles and triangles, respectively). Elsie is deepest in the B band (the bluest of the three bands) and shallowest in the i’ band (the reddest of the three). The dependence of the dimming on color suggests that circumstellar dust is responsible for the dip. [Boyajian et al. 2018]

Elsie is deeper in the blue than it is in the red! From that, we can deduce that whatever is blocking the light from Boyajian’s Star is less amenable to letting blue light through than red. It’s tough to explain that behavior with an opaque object, like a planet, transiting in front of the star — rather, Dr. Boyajian and her team argue, it’s more likely that clouds of dust grains, smaller than a micrometer across, are responsible. Think of such grains as tiny glass spheres eclipsing the star, refracting starlight off its original course, scattering away short blue wavelengths and leaving longer red wavelengths less affected.

…Or Is It?

Of course, these proposed dust grains still have to come from somewhere, and that’s an entirely new puzzle. Micrometer-sized grains are so small that they get pushed around — or, more accurately, away — by starlight itself. If dust is the answer, it must be continuously resupplied or created around Boyajian’s Star. Dusty comets, planetesimals, or collisions between such objects are one possible source of dust, so the exo-comet hypothesis (Dr. Boyajian’s original explanation for the Kepler dips!) might be back in play. Luckily, analysis of the colors of the other three dips is in the works, and Las Cumbres is still looking. Stay tuned!

About the author, Emily Sandford:

I’m a PhD student in the Cool Worlds research group at Columbia University. I’m interested in exoplanet transit surveys. For my thesis project, I intend to eat the Kepler space telescope and absorb its strength.

double pulsar

Editor’s note: Astrobites is a graduate-student-run organization that digests astrophysical literature for undergraduate students. As part of the partnership between the AAS and astrobites, we occasionally repost astrobites content here at AAS Nova. We hope you enjoy this post from astrobites; the original can be viewed at astrobites.org!

Title: A Direct Measurement of Sense of Rotation of PSR J0737-3039A
Authors: Nihan Pol et al.
First Author’s Institution: West Virginia University
Status: Accepted in ApJ

Disclaimer: Several of the scientists on today’s featured paper are collaborators of this article’s author; however, she had nothing to do with this work.

In 2017, the field of pulsar astrophysics turned 50 years old. In those 50 years, we’ve learned a lot about these rapidly rotating neutron stars — and they’ve facilitated some of the most exciting scientific discoveries of the last century. It’s easy to assume that in the years since 1967 we’ve confirmed most of the basic assumptions we make about these compact objects (like their evolutionary history, motion, and composition). However, today’s astrobite reveals that we still have a lot to learn.

pulsar

Diagram of a pulsar, a rotating neutron star with a strong magnetic field. [NASA/Goddard Space Flight Center Conceptual Image Lab]

We use pulsars to do unfathomably cool science, from testing general relativity and detecting gravitational waves to enabling spacecraft navigation. With such an advanced understanding of these objects, you might expect us to be able to determine a property as basic as what direction a pulsar is rotating. Remarkably, we haven’t been able to — at least not with much certainty. This study from Pol et al. presents the first direct determination of a pulsar’s sense of rotation (i.e. if it’s rotating prograde or retrograde with respect to its orbit). It also provides support for the widely accepted rotating lighthouse model of pulsars, which describes pulsars as rotating neutron stars with radiation emitted from their magnetic poles. Their beams sweep across our line of sight, creating the rapid pulsing signals we detect.

The authors of this work exploited a one-of-a-kind system, called PSR J0737–3039 (the Double Pulsar). It is the only known binary that consists of two detectable radio pulsars, and its fast orbital period (~2.5 hours) makes it the most relativistic pulsar binary we know of. More subtle features (detailed below) make this system ideal for the determination of a pulsar’s sense of rotation.

The Double Pulsar system consists of a pulsar with a 22.7-millisecond period (pulsar A) in orbit with a pulsar of period 2.8 seconds (B). Magnetic dipole radiation (for a review, scroll to Energetics here) from pulsar A has been shown to introduce a new signal (the “modulation signal”) when A’s radiation hits B’s magnetosphere, which is the subtle feature that enables Pol et al.’s analysis. In general terms, they are able to deduce the direction in which pulsar A is spinning by looking at the modulation signal. If the modulation signal has a slightly longer period than if the pulsar were not rotating, they can conclude that it is rotating in the same direction (prograde) as its orbital motion. If the period is shorter, it implies that pulsar A’s rotation is retrograde relative to its orbit.

Brave astrobites readers will want to refer to the paper that describes the algorithm the authors used to resample and transform the data. After these transformations, the main analysis occurs by looking at Fourier power spectra. A Fourier transform splits a signal into its constituent frequencies, and the power spectrum shows how much signal (technically, the average signal squared, or power) is present at any given frequency. In short, the scientists perform three tests by looking at the Fourier power spectrum corresponding to prograde, retrograde, and no rotation. The three possibilities are denoted by a value s, which could be -1, 0, or 1 (retrograde, none, or prograde). The signal from A is corrected in three separate trials, once for each value of s, to see which sense of rotation pushes the most power into the frequency of A’s main signal (i.e. which correction works the best). They find that s = 1, indicating prograde motion, is the best correction (see Figure 1).

Figure 1. The Fourier power spectrum of the modulation signal for s = -1, 0, and 1. More power appears in the s = 1 trial (cyan) than the other two; therefore, the authors conclude that the pulsar is rotating in the same direction as its orbit (prograde). [Pol et al. 2018]

This unique analysis is the first time the sense of rotation of a pulsar has been directly determined, which is an exciting result in its own right. However, this information can be used in several applications. For example, knowing the direction of rotation of a pulsar with respect to its orbit can help inform theories describing the evolutionary history of these systems. Pulsar B is much younger than A, so it might be interesting to know how much the supernova that birthed pulsar B affected pulsar A. The results of today’s paper suggest that the “kick” from B’s supernova did very little to disturb A’s motion. Additionally, the direction of rotation can (along with other measured orbital parameters) help determine the moment of inertia of A. With this information and an accurate mass measurement, it is possible to deduce the radius of the pulsar. Very little is known about the structure and composition of neutron stars, and a radius measurement can help inform and rule out proposed theories.

In a field as strange and complicated as astronomy, it’s wise not to assume what we know and don’t know about the universe. Today’s astrobite reminds us that despite our solid understanding of many complex issues in astrophysics, innovation still continues in understanding basic traits of the systems we study.

About the author, Thankful Cromartie:

I am a graduate student at the University of Virginia and completed my B.S. in Physics at UNC-Chapel Hill. As a member of the NANOGrav collaboration, my research focuses on millisecond pulsars and how we can use them as precise tools for detecting nanohertz-frequency gravitational waves. Additionally, I use the world’s largest radio telescopes to search for new millisecond pulsars. Outside of research, I enjoy video games, exploring the mountains, traveling to music festivals, and yoga.

galaxy disruption

Editor’s note: Astrobites is a graduate-student-run organization that digests astrophysical literature for undergraduate students. As part of the partnership between the AAS and astrobites, we occasionally repost astrobites content here at AAS Nova. We hope you enjoy this post from astrobites; the original can be viewed at astrobites.org!

Title: Empirical Determination of Dark Matter Velocities using Metal-Poor Stars
Authors: Jonah Herzog-Arbeitman, Mariangela Lisanti, Piero Madau, Lina Necib
First Author’s Institution: Princeton University
Status: Submitted to ApJL

Our galaxy is embedded in a cloud of dark matter, thought to consist of tiny particles traveling along orbits through the halo. These dark matter particles permeate all regions of the galaxy, extending far beyond the edge of the bright central spiral, but also orbiting through our solar system, and even passing right through the Earth. This is why scientists build giant detectors, hoping to trap some of these dark matter particles as they pass by. So far, these experiments have not detected dark matter, but that lack of detection is actually quite interesting. Finding out what dark matter is not, and thereby narrowing down the possibilities, is an important step towards revealing the true nature of these mysterious particles.

In order to really understand what it means when a detecter does not see dark matter, it is important to have a clear prediction for how much dark matter should be detected. For example, if we expect very few dark matter particles to pass through the Earth in a given amount of time, then maybe the lack of detections over a few years doesn’t actually mean those particles don’t exist. One essential piece of information in this prediction is the velocity of dark matter particles as they orbit past our solar system.

So, how can we determine the speed of these particles that we haven’t even directly detected? Well, let’s look back at where these particles actually come from. Dark matter halos grow over time by consuming other dark matter halos. This process is called hierarchical structure formation. The Milky Way is continuously pulling in smaller galaxies and then tearing them apart, thoroughly mixing their stars and dark matter particles into the Milky Way halo (Figure 1).

This understanding of the origin of these particles reveals an important piece of information: when dark matter particles join the Milky Way, they are often accompanied by stars. This is great news, because stars, unlike dark matter particles, are not invisible, and we can directly measure their velocities. If we can confirm that dark matter particles tend to move at similar velocities to their stellar companions, then this problem of determining the local dark matter velocity is much simpler! Finding out if this is in fact the case is exactly the goal of today’s paper.

The tricky thing here is that the Milky Way is continuously forming new stars, so the authors need to find a way to distinguish between the stars formed within the Milky Way and stars that formed in smaller galaxies and were then consumed by the Milky Way along with the corresponding dark matter. This turns out to be fairly straightforward: stars that form in smaller galaxies tend to have a different chemical composition than stars that are currently being formed in the Milky Way. This is because Milky Way stars are forming from materials that have been enriched with heavier elements by generations of star formation, while the stars in smaller galaxies are not. The stars we are interested in are therefore what astronomers call “metal-poor.” The prediction is therefore that metal-poor stars and dark matter particles should have similar velocities.

Figure 2. A simulated Milky Way-like galaxy, from the ERIS simulation used in this paper. [Simone Callegari]

The authors use simulations of Milky Way-like galaxies (Figure 2) to compare the velocities of dark matter and stars, and find that this prediction holds up! Figure 3 shows the distributions of velocities for dark matter and different stellar populations. The black histogram is dark matter, the cyan histogram is all stars, and the orange histogram is only metal-poor stars. The black and orange histograms line up pretty well, meaning the velocity of metal-poor stars does tend to match that of the dark matter. This means that by observing the velocities of these stars near the Sun, we can improve our understanding of the dark matter velocity. This will improve our interpretation of the results of dark matter experiments. In particular, based on preliminary calculations, the authors show that the velocity is lower than previously thought. They suggest that this may weaken the significance of non-detections at smaller dark matter particle masses.

Figure 3. Velocity histograms of different components of the Milky Way, as seen in the ERIS simulation. The black histogram shows the velocity distribution of dark matter. The cyan histogram illustrates the velocity of all stars, and has a much larger central peak than the dark matter distribution. The orange histogram, however, which includes only metal-poor stars, is very similar to the dark matter velocity distribution. [Herzog-Arbeitman et al. 2018]

This is a really exciting result. Previous estimates of the dark matter velocity all came from simulations and theoretical predictions, so this new method, which uses observations of our actual galaxy, rather than a simplified model, should really improve the accuracy of these calculations. Furthermore, current experiments like Gaia are greatly improving our understanding of the local stellar velocity distribution, which will continue to increase the power of this method to determine the local dark matter velocity.

About the author, Nora Shipp:

I am a 2nd year grad student at the University of Chicago. I work on combining simulations and observations to learn about the Milky Way and dark matter.

dwarf galaxy

Editor’s note: Astrobites is a graduate-student-run organization that digests astrophysical literature for undergraduate students. As part of the partnership between the AAS and astrobites, we occasionally repost astrobites content here at AAS Nova. We hope you enjoy this post from astrobites; the original can be viewed at astrobites.org!

Title: Hunting Faint Dwarf Galaxies in the Field Using Integrated Light Surveys
Authors: S. Danieli, P. van Dokkum, C. Conroy
First Author’s Institution: Yale University
Status: Submitted to ApJ

One marvelous fact about our universe is that at the largest scales, it is fractal. Unlike true fractals, which exhibit exact self-similarity, the universe is only statistically self-similar. If you looked at the most massive objects in the universe, a record held by the gargantuan, invisible dark matter blobs holding clusters of galaxies together, which clock in with masses upwards of 1015 times the mass of the Sun, you’d find that they’re rife with smaller blobs, or “halos” of dark matter. Many of these smaller dark matter halos are inhabited by a galaxy, including giant bright elliptical galaxies, the smaller and fainter spiral galaxies, and hordes of yet smaller, fainter galaxies. Peering closer at, say, one of the dark matter halos of a Milky Way-like spiral galaxy, which clocks in at about 1012 times the mass of the Sun, you’d find that it in turn is surrounded by a similar but down-sized army of even smaller dark matter halos, which may contain even fainter “dwarf” galaxies. And the halos of each of these dwarf galaxies in turn can host their own army of even tinier dark matter halos. If you just looked at the dark matter of a galaxy cluster, a single spiral galaxy, or a dwarf galaxy, it would be hard to tell which was which — they would roughly look like scaled up (or down) versions of each other.

How far down does this fractal structure go? We can search part of the way down by searching for the smallest, faintest galaxies that live within them — which is an incredibly difficult task. To go further down to the smallest dark matter halos, which may be completely dark, and thus unobservable by usual means (i.e. by light), we’ll have to turn to more exotic methods. The faintest dwarf galaxies we’ve found thus far have been discovered by hunting for clusters, or “overdensities,” of stars. This technique can only uncover dwarfs in which we can observe individual stars, which we can distinguish only out to a pitiful distance — just to up to about 5 Mpc away, which is a little beyond the edge of the Local Group of neighborly galaxies.

What about the faint dwarfs that live even further away, far enough that they appear as fuzzy patches of light, and not as collections of stars? The authors of today’s paper discuss our prospects for finding these “integrated light” images of dwarfs (see Fig. 1) as far away as 10 Mpc. To do this, they set out to ask a simple question: how many galaxies could they find, given a telescope with a particular resolution and sensitivity?

Figure 1. Simulated observations of a faint dwarf galaxy at different distances from the Milky Way. If the dwarf is just outside the Milky Way, at about 500 kpc, we can see the individual stars within the galaxy. However, if such a galaxy is farther away, it becomes increasingly difficult to resolve individual stars. At 4 Mpc, it only appears as a fuzzy blob of light, and we can no longer see the individual stars in the galaxy. Detecting such faraway, faint dwarfs requires new search methods. [Danieli et al. 2017]

To carry out this calculation, the authors assume — for lack of data — that the faint galaxies as far as 10 Mpc look like the ones we’ve seen that are close by. Based on these nearby galaxies, they estimate how large and faint the distant dwarfs they’re searching for would be, then determine whether or not we could see them. They also estimate the mass in stars each of these faraway dwarfs have — a key quantity that allows them to guess the mass of the dark matter halos the dwarfs inhabit. It’s as yet unclear how much dark matter a galaxy with a given mass in stars has — a mapping we call the stellar mass-halo mass (SMHM) relation — so the authors adopt two different ones. With the dark matter mass of the dwarf galaxies, it’s simple enough to determine the number of such dwarfs that should exist out to 10 Mpc from simulations of the Milky Way and its surroundings.

Figure 2. The number of dwarfs between 3–10 Mpc that we could see with a telescope of a given resolution and sensitivity. The angular resolution is shown on the horizontal axis, and the sensitivity, here quantified as the surface brightness μ, is shown on the vertical axis. The colors and black contour lines denote the number of dwarfs you can see per square degree in the sky (an area equivalent to five times the Moon’s). The two panels show results from two different SMHM relations (see above paragraph for details). Brighter dwarf galaxies — those with smaller μ and thus at the bottom of the plots — can be seen no matter the resolution of your telescope. Fainter dwarfs — those with larger μ and higher up on the plot — are found in greater abundance, and we need good spatial resolution (fewer arcsecs) to detect them all. [Danieli et al. 2017]

The authors find (see Fig. 2) that, as you might expect, brighter dwarfs can be discovered no matter how good the resolution of your telescope is. However, when attempting to discover fainter dwarfs, the resolution really begins to matter. If, say, you had a telescope with a resolution of 9 arcsec versus one that was twice as good, you could detect up to six times as many of the faintest dwarfs. They also find that the SMHM they assume can affect the number of the faintest dwarfs they expect to find by as much as a factor of five.

The authors calculate that using this “integrated light” method to hunt for faint dwarfs using the Dragonfly Telescope Array, a telescope that was designed for the task, we could find a similar number of galaxies — if not more — as with surveys that rely on the traditional method of finding clusters of individually resolved stars. This is an exciting result. The few smallest, faintest galaxies we’ve found so far currently puzzle astronomers: how many of them are there? Why are they so faint? How did their stars form? We could begin to unravel these mysteries once we find more of these tiny galaxies.

About the author, Stacy Kim:

I am a fourth-year graduate student in The Ohio State University’s Department of Astronomy. On a day-to-day basis, you can typically find me attempting to smash clusters of galaxies together inside big supercomputers with Dr. Annika Peter to see if cluster mergers are good testbeds for dark matter collisionality. As an undergraduate at Caltech, I spent a few years chasing photons where planets are thought to form (or, as they say, performing Monte Carlo radiative transfer calculations of protoplanetary disks) with Dr. Neal Turner of the Jet Propulsion Laboratory. When I’m not sitting in front of a computer trying to translate cosmic thoughts into pithy lines of code, you can find me in the kitchen or on the walls of a climbing gym.

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