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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: LOFAR Discovery of a 23.5 s Radio Pulsar
Authors: C. M. Tan et al.
First Author’s Institution: Jodrell Bank Centre for Astrophysics, University of Manchester, UK
Status: Published in ApJ

Pulsar Rotation Rates

Neutron stars are formed from massive stars that undergo violent supernova explosions after they run out of nuclear fuel and collapse under their own gravity. Radio pulsars are highly magnetized, rotating neutron stars that emit beams of radiation from their magnetic poles. When these beams of radio emission sweep across our line of sight, they generate radio pulses that can be detected with radio telescopes on Earth. The surface magnetic field strength, age, and internal structure of these objects can be studied through measurements of their rotational rates. Astronomers have now discovered more than 2,700 pulsars in our galaxy, and they’re constantly on the lookout for rare breeds. In today’s astrobite, we cover the discovery of the slowest known spinning radio pulsar, PSR J0250+5854, which has a rotational period of 23.5 s. This exciting finding demonstrates that radio pulsars can rotate much slower than expected and still produce radio pulsations.

LOFAR Superterp

Figure 1: An aerial view of the LOFAR Superterp, part of the core of the extended telescope located in the Netherlands. [LOFAR / ASTRON]

PSR J0250+5854: A Record-Setting Slow-Spinning Radio Pulsar

The authors discovered PSR J0250+5854 on 2017 July 30 using the LOw Frequency ARray (LOFAR) radio telescope (see Figure 1) as part of the LOFAR Tied-Array All-Sky Survey (LOTAAS). Additional follow-up radio observations were performed using the Green Bank, Lovell, and Nançay radio telescopes. Pulsations were detected between 120 and 168 MHz with LOFAR and at 350 MHz using the Green Bank Telescope (GBT), but no pulsed emission was detected at ~1.5 GHz using the Lovell and Nançay telescopes. The pulsar’s radio spectrum (spectral index of α = -2.6 ± 0.5, assuming its flux density follows a power-law as a function of frequency) is remarkably steep compared to the average pulsar population (<α> ≈ -1.8). This suggests that its radio emission is significantly brighter at lower frequencies (see Figure 2).

Radio spectrum of PSR J0250+5854

Figure 2: Radio spectrum of PSR J0250+5854 using LOFAR and GBT observations. The black line shows the fitted spectral index, with 1-σ uncertainties indicated by the shaded gray region. The circle corresponds to the measured flux density from LOFAR Two-meter Sky Survey imaging observations, and the triangles correspond to upper limits on the flux densities from LOFAR Low Band Antenna, Nançay, and Lovell radio telescope observations, respectively. [Tan et al. 2018]

Based on measurements of the pulsar’s rotation spanning more than 2 years, PSR J0250+5854 has an inferred surface dipole magnetic field strength of 26 trillion Gauss, characteristic age of 13.7 million years, and a spin-down luminosity of 8.2 x 1028 erg s-1, assuming a dipolar magnetic field configuration. PSR J0250+5854’s radio beam is very narrow according to the measured width of its pulse profile (the pulse duty cycle is < ~1% below 350 MHz, see Figure 3). Individual single pulses were routinely detected from the pulsar at low radio frequencies, except during brief periods of “pulse nulling” when the pulsar stopped emitting radio pulses. This occurred 27% of the time on average. The pulsar’s slow rotation period of 23.5 s is similar to other classes of pulsars. In particular, magnetars have high magnetic fields, spin periods ranging between roughly 2 and 12 s, and often produce X-ray emission, and X-ray Dim Isolated Neutron Stars (XDINs) have spin periods ranging between 3.4 and 11.3 s. However, no X-ray emission was detected from PSR J0250+5854 during follow-up observations with the Neil Gehrels Swift Observatory X-ray Telelescope.

pulse profiles of PSR J0250+5854

Figure 3: Integrated pulse profiles of PSR J0250+5854 at observing frequencies of 350 MHz (GBT), 168 MHz (LOFAR), and 129 MHz (LOFAR). Here, only 5% of the rotational phase is shown. The inset shows the pulse profile across the whole LOFAR HBA band over a full rotation period. [Tan et al. 2018]

A Needle in a Haystack or a Haystack Full of Needles?

The P–Ṗ diagram is a key diagnostic tool for characterizing how pulsars evolve in time. Using pre-discovery LOTAAS data of PSR J0250+5854 from 2015, the authors measured a spin period derivative of Ṗ = 2.7 x 10-14 s s-1. The pulsar’s rotational parameters place it in the right region of the P–Ṗ diagram (see Figure 4) — an area where few pulsars have been found to reside. In particular, PSR J0250+5854 falls near/below many of the so-called “pulsar death lines,” beyond which pulsars are not expected to emit coherent radio emission. These models are based on assumptions about the conditions in the pulsar’s magnetosphere, such as pair production, which is thought to be essential for the generation of radio emission. Since the radio-emission mechanism in pulsars is not fully understood, searching for additional pulsars near these death regions will help to inform us about how pulsars produce radiation.

P–Ṗ diagram of pulsars

Figure 4: P–Ṗ diagram of pulsars derived from their measured rotational periods and rotational-period derivatives. The positive sloped gray lines indicate characteristic ages of 1 kyr, 100 kyr, 10 Myr, and 1 Gyr. The negative sloped gray lines correspond to inferred surface magnetic-field strengths of 10 GG, 100 GG, 10 TG, and 100 TG. Magnetars (green), XDINSs (orange), RRATs (yellow), and the 8.5-s radio pulsar PSR J2144–3933 are indicated on the plot. The colored lines show the various death-line models, where pulsars below these lines are not expected to produce radio emission. [Tan et al. 2018]

The discovery of PSR J0250+5854 begs the question: Is this a special kind of pulsar, or are there more to be found? The authors argue that more of these slow-rotating pulsars may be lurking around our galaxy, but we simply haven’t been sensitive to detecting them because commonly used Fast Fourier Transform (FFT)-based periodicity search algorithms are not well-suited to detecting slow pulsars with small duty cycles. The authors also point out that the radio emission observed from PSR J0250+5854 was much more erratic at higher frequencies. Therefore, if other slow rotating pulsars are similar to PSR J0250+5854, then this suggests that low-frequency radio telescopes, like LOFAR, may prove to be excellent observatories for searching for these slow rotators.

About the author, Aaron Pearlman:

I am a Ph.D. candidate in Physics at Caltech. My research focuses on searching for new pulsars near the center of the Galaxy using JPL’s Deep Space Network radio dishes in the southern hemisphere. I am also interested in studies of magnetars, fast radio bursts, gravitational-wave searches, and high-energy observations of compact objects. When I’m not hunting for pulsars, I can usually be found hanging out with my dogs or trying the latest vegetarian cuisine Los Angeles has to offer!

planet formation

Editor’s note: AAS Nova is on vacation this week. Normal posting will resume next week; in the meantime, we hope you enjoy this post from Astrobites, a graduate-student-run organization that digests astrophysical literature for undergraduate students. The original can be viewed at astrobites.org.

Title: On the Terminal Rotation Rates of Giant Planets
Authors: Konstantin Batygin
First Author’s Institution: California Institute of Technology
Status: Published in AJ

The rotation periods of Jupiter and Saturn are 9.93 hours and 10.7 hours, respectively. Now, compared to our tiny Earth that lazes around on a 24-hour rotational period, you might think, “wow, those are some zoomy-bois.” However, our best theories of planet formation tell us that, based on how massive they were when they formed, they should really be doin’ a faster spin.

Fun fact alert: While the sun holds most of the solar system’s mass, Jupiter and Saturn hold the majority of our solar system’s angular momentum.

How Do You Form a Jupiter?

So let’s say you want to make a Jupiter, just for the heck of it. If you follow the rules of our understanding of general planet formation, there are three main steps. You start inside of a protoplanetary disk. There is some sort of gravitational instability where heavy metals can gravitationally collapse and start to form a metallic core. This core acquires a gaseous envelope which can then feed the newly forming planet. Once that gaseous envelope is about the size of the initial core, the planet enters a stage called runaway accretion. That just means that material around the planet falls quickly and efficiently, adding mass rapidly. And BAM — you have a Jupiter-like planet (technically called a Jovian planet). But, following this simple model, once the runaway process begins, the planet is accreting so much mass that our new Jupiter spins faster and faster and has no way to let go of any of its angular momentum. In this simple model, the surface of the planet can reach speeds that equal the escape velocity, which means that the planet breaks apart. That’s not great for planet formation. Plus, when we observe Jupiter and Saturn (and now that we’re gathering more and more data of Jovian planets outside of our solar system), we continue to see rotational velocities well below the planets’ escape velocities. So how in the heck do we slow down a young energetic planet? We turn to the answer that all astronomers look to in times of need: magnetic fields.

Setting Up the Problem of Slowing Down a Chonker Like Jupiter

Today’s paper attempts to lay the groundwork for solving this angular-momentum problem in Jovian-planet formation using magnetohydrodynamics. Big (scary) word, yes, but put more simply, this paper creates a semi-analytic model of a newly forming Jovian planet with a strong magnetic field, and it then explores how the field might slow the planet down. The model breaks the problem into two parts: the circumplanetary disk and the planet itself. Each part of this problem has equations that describe key parameters, such as the temperature, density, and abundance of metals in the surrounding envelope. For the planet part, the author calculates a magnetic field strength based on a typical luminosity of young exo-Jovian planets and uses these properties to calculate the electric conductivity and magnetic induction of the system, which would produce the forces that affect the speed of planet rotation. “Running” this model consists of calculating each of these equations over a series of time steps so that one can further understand how each of these factors change and affect each other as the planet forms.

How the Giant Chonker Was Slowed Down

The results of the model are illustrated in Figure 1 below. The finding of this paper is that if we consider the Jovian protoplanet to have a significant magnetic field, that field will invoke a force in the opposite direction of the rotation of both the circumplanetary disk and the planet itself. Basically, the magnetic field couples to the surrounding disk. Since there is now a force in the opposite direction of the original motion, the planet slows its spin. Angular momentum leaves the system as material feeding the planet gets kicked out of the system and back into the surrounding protoplanetary-disk environment.

planet formation

Figure 1: An illustrative view of planet formation and the effect of magnetic fields (red lines). We are taking a look inside a protoplanetary disk, with the host star to the left, zooming in on a Jupiter-like planet being formed. The planet has its own circle of influence, the edges of which correspond to the purple regions. We can see that material flows onto the planet from above, and that material can only fall onto the planet if it is very nearly falling directly down. Material that falls just off to the side gets added to the de-cretion disk and thus shucked off into the gaseous nebula. The planet slows down via magnetic-field induction that invokes a force in the opposite direction of the original Keplerian rotation, which is the same direction as the planet is rotating. [Batygin 2019]

This paper set up a simple semi-analytic model that did the job of adding magnetic fields to our picture of planet formation. And this model has shown that with a strong magnetic field, is it not only possible to slow the spin of the planet down to speeds that we observe, but also to slow it down quickly. Of course, there are more details that could be added to this model, and there are assumptions made that are difficult to back up with observations. Planet formation, in general, seems to happen relatively fast and early in a solar system’s lifetime, so it is hidden from view and hard to catch. But even with this “simple” model, we can see that magnetic fields are certainly a key factor in the mystery of the slowly spinning Jovian planets.

About the author, Jenny Calahan:

Hi! I am a second-year graduate student at the University of Michigan. I study protoplanetary disk environments and astrochemistry, which set the stage for planet formation. Outside of astronomy, I love to sing (I’m a soprano I), I enjoy crafting, and I love to travel and explore new places. Check out my website: https://sites.google.com/umich.edu/jcalahan

soho image of sun

Editor’s note: AAS Nova is on vacation this week. Normal posting will resume next week; in the meantime, we hope you enjoy this post from Astrobites, a graduate-student-run organization that digests astrophysical literature for undergraduate students. The original can be viewed at astrobites.org.

Title: Was the Sun a Slow Rotator? — Sodium and Potassium Constraints from the Lunar Regolith
Authors: Prabal Saxena, Rosemary M. Killen, et al.
First Author’s Institution: NASA Goddard Space Flight Center
Status: Published in ApJL

Journey to the Sun’s Past

Throughout the solar system’s history, the frequency of flares and eruptive events from the Sun have had a strong effect on the development of the inner planets, from the top of their atmospheres right down to their surfaces. The number of eruptive events the Sun produces just so happens to be closely related to the rate at which it rotates. Therefore, to understand how planets like the Earth came to be, it is incredibly important to understand the Sun’s rotation during the early stages of our solar system’s development. Previous studies have attempted to do so by considering other Sun-like stars. However, today’s authors have found answers by looking much closer to home.

The Moon, Earth’s only natural satellite, is a surprisingly ideal place to look for clues about the history of solar activity. The lack of a thick atmosphere causes solar eruptions that reach the Moon to strip material from its surface, leaving behind an imprint that can be used to understand the Sun’s tumultuous past (but more on that later).

A Model for the Early Sun

Before the authors of today’s paper dove into how the Moon is an ideal place to look for evidence of past solar activity, their first challenge was to model the Sun’s activity over its entire lifetime. As mentioned previously, the primary cause for depletion of material from the lunar surface is from space weather events — most notably coronal mass ejections (CMEs), which occur when large volumes of material are ejected from the Sun during a solar eruption. A cartoon depicting this is shown in Figure 1.

rotation rates

Figure 1: Cartoon showing the relationship of solar rotation rate with the amount of material lost from the surface of the Moon. [Saxena et al. 2019]

The authors considered three rotation classes for their model of the early Sun: slow rotators, medium rotators, and fast rotators, which correspond to rotation rates observed for Sun-like stars in a previous study. To construct a flare/CME relation, they looked at data from both the Kepler Space Telescope and Earth’s geological record. The Kepler Space Telescope observed several Sun-like stars in a single patch of sky over the span of four years and was able to characterize the activity of Sun-like stars with respect to their rotation. An approximately linear flare–rotation relation was found, with faster rotators being more active than slow rotators. CMEs are also always associated with a flare, however not every flare produces a CME — especially at low energies. The authors consider only the most energetic flares, and therefore assume a 100% CME–flare association.

The authors’ fully constructed CME frequency for the entire Sun’s lifetime is shown in Figure 2. Regardless of initial rotation rate, all scenarios converge to the same CME/flare rate, as shown in the figure.

CME frequency vs. time

Figure 2: Plot showing CME frequency versus time for three initial solar rotation classes. Different colored lines indicate different flare energies over time from different data sources (Kepler and Earth’s geological record), while solid, dashed, and dotted lines indicate the fast, medium, and slow rotator model, respectively. [Saxena et al. 2019]

Moon Rocks Rock

With their model CME frequency history in hand, the authors could now dig deeper into how solar activity could have affected the present-day composition of the Moon. The Moon is a surprisingly ideal place study past solar activity due to the cataclysmic event that formed the Earth–Moon system.  The most widely accepted theory for our Moon’s formation involves a large Mars-sized object, Theia, which crashed into the primordial Earth some four billion years ago. At the time of the Moon’s formation, the Earth and Moon had the same surface composition, having been formed from the same mass of rock. Thanks to our thick atmosphere and magnetic field, Earth has been able to hold on to a lot of the material on its surface since its formation. The Moon, however, is far too small to have an atmosphere. Therefore, material on the surface is constantly being stripped off due to many factors, the most effective being space-weather events. To be able to accurately assess the difference between material on the Earth and Moon, the authors focused on two volatile elements on the lunar surface, sodium and potassium. These elements have moderately lower abundances on the Moon than on Earth, but they are abundant enough on the Moon to be accurately measured.

Based on the amount of sodium and potassium currently present on the Moon, the authors then determined how fast the Sun had to be rotating in order to account for the difference relative to the abundance on Earth. They found that a fast-rotating Sun would have depleted the Moon’s sodium and potassium far more than the present-day values suggest. For the medium rotator case, the authors found that this model would account for the present-day values of sodium, but the potassium values would not have lined up. That leaves the slow rotator model. An initially slowly rotating Sun would account for the difference of both sodium and potassium between the Moon and the Earth.

A Solar-System Scale Puzzle

Understanding the initial rotation rate of the Sun is necessary for understanding the evolution of planets in the inner solar system. Many other mechanisms were used to try to explain the degree of sodium and potassium depletion, namely the amount of exposed vs. buried material in the Moon’s surface, meteorite impacts, volcanism, and magnetism (both from the Earth and the Moon itself). However, none of them would fully account for the amount of depletion observed. In today’s work, the authors found that a slow rotator model for the Sun best explains the current amounts of volatile elements present on the surface of the Moon and puts yet another piece into the billions-year long puzzle of our solar system’s history.

About the author, Ellis Avallone:

I am a first-year graduate student at the University of Hawaii at Manoa Institute for Astronomy, where I study the Sun. My current research focuses on how the solar magnetic field triggers eruptions that can affect us here on Earth. In my free time I enjoy rock climbing, painting, and eating copious amounts of mac and cheese.

M80

Editor’s note: AAS Nova is on vacation this week. Normal posting will resume next week; in the meantime, we hope you enjoy this post from Astrobites, a graduate-student-run organization that digests astrophysical literature for undergraduate students. The original can be viewed at astrobites.org.

Title: New s-process Site in Rapidly-Rotating Massive Pop II Stars
Authors: Projjwal Banerjee, Alexander Heger, Yong-Zhong
First Author’s Institution: Indian Institute of Technology Palakkad, India; Shanghai Jiao Tong University, China
Status: Submitted to ApJ

One of the main goals of nuclear astrophysics is to understand and explain the composition of the universe. Starting out with hydrogen, helium, and a teeny bit of lithium, the universe evolved to have an entire periodic table worth of chemical elements. How did this happen? We know that stars and supernovae played a key role in this chemical enrichment. However, exactly how and where different elements are produced, and in what quantities, remains a topic of vigorous ongoing research. Heavy element nucleosynthesis is a key piece of this puzzle and has received a lot of attention recently in the context of neutron star mergers.

Heavy elements — i.e., elements heavier than iron — are mainly formed through the s-process and the r-process. The names describe the speed of the process relative to the decay time of the isotopes involved: s-process involves slow neutron capture and the r-process involves rapid neutron capture.

Although both these processes make roughly equal overall contributions to the total abundance of heavy elements, the s-process does not get the same hype as its cool twin, the r-process. This is partly due to the fact that the site of the r-process was a huge mystery until very recently, when neutron star mergers were identified as one of the sites where the process occurs. It still remains to be seen whether they are the only site of the r-process.

The s-process, on the other hand, is known to occur inside asymptotic giant branch (AGB) stars. This is the final stage of evolution for long-lived, low-mass stars between 1–3 solar masses. The AGB stars can produce elements up to 209Bi, forming what’s called the “main” component of the s-process. However, a weaker version of the s-process also happens in massive stars (>10 solar masses), producing elements up to atomic mass number A ~ 90. Today’s paper shakes things up by introducing a new site for the main s-process: in rotating metal-poor massive stars!

The authors find that above a critical rotation speed, massive metal-poor stars evolve in a quasi-chemically-homogeneous (QCH) manner. This means the stars are spinning so fast that the mixing caused by rotation becomes very efficient. At the end of core hydrogen-burning, which produces helium, these stars look like helium stars. We can see this in panel (a) of Figure 2. Now starts the story of the s-process, which goes like this:

  1. The star in the QCH state starts burning helium in the center, producing 12C. Its core becomes convective.
  2. The convective core grows in size. Some of the 12C is mixed outward into the radiative region of the star. In this region, there are still some protons present, along with the right thermodynamic conditions to allow 12C to react with protons. This produces 13C.
  3. Some 13C is mixed back into the convective He-burning core. In regions that are hot enough, 13C reacts efficiently with the 4He, producing oxygen, as well as neutrons. This leads to high neutron densities that allow the s-process to happen!

We can see this play out in panels (b) through (d) of Figure 1.

Figure 1. Isotopic composition of the star during different stages of its evolution. The y-axis gives the mass fraction of different isotopes and the x-axis indicates the mass coordinate, i.e., how much mass of the star we’re looking at as we move away from the center of the star. We can see the mass fraction of 13C growing as 12C gets mixed into and reacts with protons in the outer layers. This 13C gets mixed inward and reacts with 4He, producing both 16O and neutrons. [Banerjee et al. 2019]

In Figure 2, we can see the amount of heavy elements present during different stages of the star’s evolution. There is a clear and strong s-process pattern leading all the way up to the element lead!

Figure 2. Time evolution of the s-process yield pattern for a 25-solar-mass star. Number yields of heavy isotopes inside the star are plotted as functions of mass number. The different lines correspond to different stages in the star’s evolution. [Banerjee et al. 2019]

Apart from being interesting in its own right, this new site may also be important for explaining the abundance patterns seen in very metal-poor (VMP) stars. These are old stars thought to reflect the composition of the interstellar medium at  ~1 Gyr after the Big Bang. This is much too early for low-mass stars to contribute any s-process elements (they evolve slowly), implying that the heavy-element abundances seen in VMP stars must come from the r-process. However, if a strong s-process is possible in massive stars as well, we cannot be so sure! Massive stars live fast and die young, which would allow for an earlier onset of s-process enrichment. Figure 3 shows s-process yield patterns from the authors’ calculations, compared to abundances observed in two VMP stars. The fits are quite good, presenting a potential explanation for VMP observations. 

Figure 3. Heavy-element abundances from this work compared to VMP stars. The abundances are from the wind ejecta (left) and the wind ejecta combined with outer stellar ejecta (right) for a 25-solar-mass star. The stars shown here are CEMP stars. The one on the left is a CEMP-s star while the one on the right is a CEMP-r/s star. [Adapted from Banerjee et al. 2019]

Of course, we should remember that this is an initial exploration of a new idea, though it seems promising. Stay tuned to see where the research leads!

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.

M101

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: Untangling the Galaxy I: Local Structure and Star Formation History of the Milky Way
Authors: Marina Kounkel, Kevin Covey
First Author’s Institution: Western Washington University
Status: Accepted to AJ

Despite our home planet being embedded in it, the Milky Way and its immediate environment remain an enigma to astronomy. Once thought to have few satellite neighbors, The Milky Way has been found to have many dwarf galaxies orbiting it. New stellar streams are being uncovered as well, likely remnants of past gravitational interactions with dwarf galaxies, in which the Milky Way pulled rivers of stars from its now-dissipated partners. This burst of discoveries of new nearby and entangled structures are thanks to advancements in technology allowing astronomers to observe dimmer objects and to track stars with greater precision.

Today’s paper utilizes one of these advancements, the much lauded Gaia mission, in tandem with machine learning methods to identify new substructures within the Milky Way and, in so doing, learn about its murky past.

Re-Clustering the Star Clusters

To begin, the authors are presented with the challenge of identifying stellar structures within the enormous Gaia dataset. In order to group stars together the authors use a clustering algorithm, which is a series of steps designed to isolate populations of objects with similar characteristics; the characteristics in question here are the stars’ coordinates within the Milky Way, their parallaxes, and their proper motions. A data sample of over 19 million stars are selected from the Gaia catalog, chosen to isolate stars for which the above characteristics are measured with high certainty. After much testing of the algorithmic parameters, the model returns over 1,900 star clusters, many of which have been independently identified in other studies. However, they also identify new structures that appear to have eluded other investigations (Figure 1).

star clusters on sky map

Figure 1: Map projection of the portion of the sky considered in today’s paper, with algorithm-identified star clusters marked in blue. Yellow markings indicate star clusters previously identified using different methods. Galactic coordinates are indicated with b and l. [Kounkel & Covey 2019]

In order to learn about our galaxy’s past, the authors must gain more information about these clusters to construct a star formation history. The star formation history of a galaxy is exactly what it sounds like: a combination of all star-forming events in a galaxy’s past that contribute to the current picture seen by astronomers. However, one can’t fully understand the history by only knowing the what and the where of star formation; also important is the when.

The authors determine the ages of their identified clusters testing two separate methods: analysis by a convolutional neural network (CNN) and isochrone fitting. Training the CNN using both known real clusters and a multitude of artificial ones, they only reproduce the accepted ages of clusters in 44% of cases. Similarly, using isochrones alone is only successful in a minority of cases. Using the CNN age estimate as an input to their isochrone model, however, increases the success rate to 77%, so this methodology is used to obtain ages for the remainder of the work.

Finding Loose Strings

While investigating the distribution of their identified star clusters, the authors noticed that they tended to be distributed in long, narrow structures. These strings, as the authors call them, are about 200 parsecs in length and lie parallel to the plane of the Milky Way. They appear similar to stellar streams, but are these simply new streams, or something new entirely? The answer lies in a peculiar trend noticed by the authors: although these strings act very similarly to normal clusters in terms of their motion, they are markedly younger than the population of clusters as a whole (Figure 2).

star cluster age distribution

Figure 2: Histogram of the age distribution of the star clusters (called “groups” here) compared to the strings. Notice how the distribution of string ages appears to have lower ages. [Kounkel & Covey 2019]

Now, one might intuitively think that the strings were formed by tidal stretching, i.e., that the stars formed in a roughly spherical cloud that was then stretched out by tidal interactions with other structures. However, many of the strings don’t show any evidence of a residual core of stars, leading the authors to conclude that they just formed this way. This interpretation is supported by previous observations of molecular filaments within the Milky Way, long string-like structures of the dense, molecular gas that is so crucial to forming stars. The authors suggest that the strings formed from these very same molecular filaments.

string subsample

Figure 3: 3D plot of a subsample of the strings, where the thick lines represent the “spine” of the string and the thin lines perpendicular to the spine indicate the velocities of the stellar components of the string. Color indicates age, and a redder string is a younger string. Check out an interactive version of this plot on the Dr. Kounkel’s website. [Kounkel & Covey 2019]

Further, analysis of the global distribution of strings (Figure 3) indicates that strings of different ages seem to lie close together, coagulating into four coherent streams of structure. Due to a correlation between the position of the youngest stream and the Local Arm of the Milky Way, the authors contend that these collections of strings may correspond to past star formation in old spiral arms within the Milky Way that have become less visible after losing their star-forming gas.

If so, deeper analysis of these strings might provide a way of studying the past structure and star formation history of our home galaxy.

About the author, Caitlin Doughty:

I am a fourth-year graduate student at New Mexico State University. I use cosmological simulations to study galaxy evolution during the epoch of reionization, with a focus on metal absorption in the circumgalactic medium.

Centaurus A

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: Positive and Negative Feedback of AGN Outflows in NGC 5728
Authors: Jaejin Shin, Jong-Hak Woo, Aeree Chung, Junhyun Baek, Kyuhyoun Cho, Daeun Kang, Hyun-Jin Bae
First Author’s Institution: Seoul National University, Republic of Korea
Status: Accepted to ApJ

One of the many mysteries of galaxy evolution is how the formation of stars is affected by a process called feedback. Unlike comments coming from a teacher on an essay, in the galactic context, feedback is coming from powerful sources of energy such as active galactic nuclei (AGN). Star formation in galaxies requires a lot of dense gas (also called the interstellar medium, or ISM), so any feedback processes that disrupt the presence or the denseness of said gas can affect the ability of a galaxy to form stars. Simulations have shown that AGN are theoretically capable of providing negative feedback by heating up the ISM or blowing it away. However, they might also provide positive feedback by compressing the ISM with their winds, making it denser and triggering bursts of star formation.

Each of these options have theoretical merit and are observed in simulations, but it can be hard to observe the effects in the wild. Today’s paper takes advantage of a particularly well-situated Seyfert 2 galaxy, NGC 5728, to enhance our understanding of AGN feedback processes. The Seyfert 2 designation is used to describe galaxies containing AGN that are similar to quasars, but that have visible host galaxies while most quasars do not.

Positive Feedback from Outflows?

The authors use observations of NGC 5728 in the optical wavelength range, the range visible to humans, to target light from stars and several emission lines generated by ionized (i.e. heated) gas. Targeting emission from a molecular transition (the carbon monoxide CO (J=2–1) transition, to be specific) also helped them trace out the distribution of molecular gas within NGC 5728, which is useful for examining how much material is available to form stars.

NGC 5728 Halpha

Figure 1: Emission map of NGC 5728 in hydrogen alpha, where the color indicates the amount of flux in a given pixel. The position of a star-forming ring and spiral arms are noted with grey dashed lines, while a biconical outflow is traced in white dashed lines. Note the black square, region A, that indicates an intersection between the AGN outflow and the star-forming ring. [Shin et al. 2019]

From these observations, the authors noted a few prominent structures. First, there are two spiral arms (only faintly visible in Figure 1). Second, there is a ring of star formation about 1 kiloparsec from the center. Lastly, there are prominent biconical outflows made up of ionized gas and full of high-energy radiation, like X-rays. Most importantly, there is an apparent intersection (labeled “A” in the figure) between the star forming ring and the northwest (i.e. the upper right in Figure 1) cone of the AGN outflow. Luckily, this intersection provides an ideal scenario for testing whether the AGN is helping or hindering star formation.

The authors define three other regions in the star forming ring (B, C, and D in Figure 2) that are located well away from the northwest biconical outflow, and can therefore serve as controls when looking for peculiarities in the star-forming characteristics of region A. Using a BPT diagram, the authors were able to calculate the percentage of flux contributed by stars in each pixel of the image and from this calculate the star formation rate in their selected regions.

AGN fraction

Figure 2: Fraction of emission in each pixel from AGN contributions (take 1 minus the AGN fraction to find the stellar contribution). A very blue color corresponds to a low AGN fraction, and thus a pixel containing gas whose hydrogen-alpha emission is dominated by stellar light. [Shin et al. 2019]

While region A has more solar masses of stars formed per year than the combined average of regions A–D, it is not particularly unusual in this respect. However, the brightness of the emission from molecular gas at region A is quite low compared to the other regions, meaning that it has an unusually high star formation efficiency (Figure 3). Indeed, it is a factor of 3–5 higher in star formation efficiency than the control regions — a significant difference! This result seems to indicate that the presence of the AGN feedback had a positive influence on the star formation, boosting it significantly and consuming a large amount of molecular gas.

star formation efficiency

Figure 3: The star formation efficiency of the galaxy, which is the ratio between the star formation rate and the available mass of molecular gas to form stars. [Adapted from Shin et al. 2019]

This seems like a pretty good indicator that an AGN should always significantly alter a galaxy’s star forming ability, right? Not quite! While this factor of 3–5 appears pretty large, when comparing to the calculated star formation rate in the entire galaxy, region A only accounts for a miniscule <10% of the total.

The jury is still out, then, on whether AGN will invariably cause galaxies to have significantly different star formation rates. This scenario is even further complicated by the fact that the biconical outflows may be simultaneously expelling accreted molecular gas from the spiral arms of the galaxy, preventing star formation from occurring there. Regardless, NGC 5728 is providing a rich test case to examine theories of AGN feedback and star formation.

About the author, Caitlin Doughty:

I am a fourth-year graduate student at New Mexico State University. I use cosmological simulations to study galaxy evolution during the epoch of reionization, with a focus on metal absorption in the circumgalactic medium.

gas-giant transit

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 Intrinsic Temperature and Radiative-Convective Boundary Depth in the Atmospheres of Hot Jupiters
Authors: Daniel P. Thorngren, Peter Gao, Jonathan J. Fortney
First Author’s Institution: University of California, Santa Cruz
Status: Submitted to ApJL

HAT-P-7b

Artist’s impression of HAT-P-7b, an inflated hot Jupiter. [NASA, ESA, and G. Bacon (STScI)]

Jupiter-sized gas-giant exoplanets in close orbits around their stars, commonly referred to as hot Jupiters, have been the prime targets for probing planetary atmospheres beyond our solar system. One of the many mystifying features of hot Jupiters — which, ironically, also makes them easier to detect and characterize — is their inflated radii. A good fraction of known hot Jupiters have sizes larger than those predicted by evolutionary models that take into account the properties of the system like temperature, age, and metallicity of the system. What could be causing these hot Jupiters to puff up?

A proposed mechanism to explain hot-Jupiter inflation is deposition of energy from stellar irradiation deep into the interiors of the planet. However, in addition to inflating the planet, energy from stellar flux heating up the planetary interiors can also radically alter the thermal structure (temperature variation with altitude) of its atmosphere which has direct consequences on its inferred atmospheric properties. Today’s paper attempts to draw a connection between the stellar irradiation of hot Jupiters and their intrinsic temperature, and how that ultimately affects the observations and our understanding of the atmospheres of these gas giants.

Structuring the Atmosphere of a Gas Giant

The vertical thermal structure — also referred to as the pressure–temperature profile — of a planetary atmosphere is directly related to change in the mode of heat transport (radiation or convection) within the atmosphere at different heights. You can think of this in the context of the Earth’s atmosphere: closer to the surface heat exchange occurs through convection, with hot parcels of air rising up and adiabatically expanding and cooling. This causes the temperature to steadily decrease as you go up until a certain altitude called the tropopause; above this you hit the stratosphere, where the air absorbs most of the heat from ultraviolet radiation from the Sun, causing the temperature to now increase with altitude. Even before this happens convection begins to weaken considerably and radiation takes over as the dominant mode of heat exchange. The altitude or the pressure level at which this happens is called the radiative-convective boundary (RCB; see Figure 1 for example). Such stratification of atmospheres is very commonly seen in planetary atmospheres in the solar system and has been studied extensively from measurements by probes like Galileo and Cassini-Huygens.

hot Jupiter profiles

Figure 1: Pressure-temperature profiles for hot Jupiters at different distances from a Sun-like star, and hence different equilibrium temperatures (Teq). Note that on y-axis, the pressure decreases as you go up, corresponding to going higher up in the atmosphere. The thick parts of the profiles mark the regions of the atmosphere that are convective, and you can see how the radiative–convective equilibrium boundary moves to lower pressures for hotter planets. [Thorngren et al. 2019]

Determining the height of the RCB for gas-giant atmospheres requires an understanding of the heat flux from the planetary interiors, which can be described by the planet’s effective intrinsic temperature (Tint). To give you an idea of the numbers, Jupiter with Tint ~ 100 K has an RCB around the height corresponding to the pressure of 0.2 bars (1 bar = pressure at sea level on Earth). In the case of hot Jupiters, on the other hand — which, given their proximity to the star, receive radiation of thousands of times that received by Jupiter — the atmosphere remains radiative to a much greater depth. Here the RCB can be expected to lie much deeper, at pressures of around 1 kilobar (remember pressure increases with depth). However, this is a good estimate only if you assume Tint ~ 100 K for hot Jupiters as well. As mentioned before, observed radii inflation of hot Jupiters points toward possible heating of their interiors by stellar irradiation (the strength of which is reflected by the equilibrium temperature of the planet Teq). This implies that hot Jupiters can have much higher Tint, which would push the region of convection and hence the RCB to larger altitudes (lower pressures). Since Tint and Teq both affect the height of the RCB, and Tint also depends on Teq, at what height should we expect the RCB for a hot Jupiter with a given Teq?

To answer this question, the authors calculate temperature–pressure profiles from thermal equilibrium atmospheric models of archetypal hot Jupiters with a range of Teq, and they then investigate how the height of the RCB changes with respect to different levels of stellar irradiation (see Figure 1 and 2).

Marking the Boundary

As is evident from Figure 1, the RCB moves to lower pressures (larger altitudes) with higher Teq, similar to how Tint increases with Teq. The surface gravity and metallicity of the planet also affect the RCB height, as seen in Figure 2.

RCB pressure vs Teq

Figure 2: The RCB pressure level with respect to the Teq of the planet, as calculated for different surface gravities and metallicities of the planet. Note that the RCB ends up at higher pressures for higher surface gravity and lower pressures for higher metallicity. [Thorngren et al. 2019]

The variation of RCB height and Tint with respect to Teq of the planet has several significant implications for models and observations of hot Jupiters. When the RCB lies at lower pressures (larger altitudes), this implies that more heat can now be deposited into the convective region of the atmosphere from stellar irradiation, allowing the mechanism of Ohmic dissipation to be even more efficient at inflating the planet. A higher Tint (of the order of few 100 K) would also affect predictions of day–night energy transport and atmospheric circulation predicted by global circulation models. It would also mean that phase curve observations of some hot Jupiters might be able to probe flux from this intrinsic heat of the planet. Moreover, higher Tint means that cloud condensation will occur much higher up in the atmosphere, affecting the observed emission from the day side of the planet.

With more exoplanet discoveries from TESS and exoplanet characterization opportunities from JWST on the horizon, we can hope to obtain a stronger constraint on atmospheric boundary conditions such as these, which would be important for accurate interpretations of exoplanet atmosphere observations.

About the author, Vatsal Panwar:

I am a PhD student at the Anton Pannekoek Institute for Astronomy, University of Amsterdam. I work on characterization of exoplanet atmospheres to understand the diversity and origins of planetary systems. I also enjoy yoga, Lindyhop, and pushing my culinary boundaries every weekend.

galaxy and CGM simulation

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 Impact of Enhanced Halo Resolution on the Simulated Circumgalactic Medium
Authors: Cameron B. Hummels, Britton D. Smith, Philip F. Hopkins, et al.
First Author’s Institution: TAPIR, California Institute of Technology
Status: Submitted to ApJ

It can be easy to think of galaxies as islands in the universe, floating around in isolation. However, a galaxy is actually surrounded by a huge sea of low-density gas that extends out to its virial radius and beyond. This gas is known as the circumgalactic medium (CGM), and more and more research is showing that the CGM has a crucial role to play in galaxy evolution. Observing the CGM has proven difficult due to its extremely low density, though, so simulations have played a large role in understanding the physics of this region. In today’s paper, the authors detail the effects of running a CGM simulation with significantly increased resolution, capable of resolving cool gas that precipitates in the CGM and rains down on the galaxy.

What Do We Know About the CGM?

Residing just outside of the galaxy, the CGM is home to large-scale flows of gas that drive galaxy evolution. These gas flows provide fuel for star formation, regulate the interactions between dark matter halos and the intergalactic medium, and contain the energy, mass, and metals of large outflows from a galaxy. In fact, the CGM is predicted to hold at least as many baryons and heavy elements as galaxies themselves, and most of the metals in the universe are found in the CGM. These metals (meaning anything heavier than hydrogen or helium in astronomy terms), deposited by galactic outflows, serve as the dominant coolant for the CGM. They are capable of radiating energy away more easily than elements like hydrogen, so an increased abundance of metals can lead to cooler gas. Consequently, this influx of metals helps to create two phases of gas: “cool” (10,000 Kelvin) gas composed of neutral hydrogen and other elements in low-energy ionization states, and “hot” (300,000–1,000,000 Kelvin) gas that contains oxygen, nitrogen, and neon in high-energy ionization states.

Unfortunately, computational work has chronically underproduced the observed abundances of these ions across redshifts by orders of magnitude. Recent work has shown that AGN feedback can increase the abundances of oxygen and other ions in the hot gas, but the discrepancy remains for hydrogen and other ions in the cool gas. In today’s paper, the authors discuss the effect of increased simulation resolution on these discrepancies.

Resolving the Resolution Issue

Perhaps one reason that simulations struggle to reproduce observations of the CGM lies in their resolution limits. Similar to how using more pixels in a television or computer screen gives a better image, increasing the resolution in a simulation means using more cells or particles to obtain a better physical picture of what is going on. However, each increase in resolution increases the computational cost of the simulation. This means your simulation that took a few days to run could instead take a few months.

Consequently, most simulations of galaxies apply their highest resolution to regions of high density where most of the matter is. This is great for figuring out what happens in the dense disk of a galaxy, but not ideal for studying the low-density CGM. Today’s paper runs simulations that force high resolution upon the CGM, reaching resolutions that are comparable to those normally obtained in the disk of the galaxy. This technique is appropriately named Enhanced Halo Resolution (EHR). Figure 1 shows the resolutions obtained by both a normal cosmological simulation and an EHR one for a region encompassing a galaxy and its surrounding filaments.

resolution plots

Figure 1: Plots of resolution for a traditional (AMR — adaptive mesh refinement) and EHR simulation. Each of these grids is made up of many cells, and spatial resolution refers to the physical length (in kiloparsecs) of the smallest cell that is present in a region. In the left panel, many galaxies are present and a particularly massive galaxy lies at the center. Its virial radius is shown by the dotted white line. Resolution in the CGM is roughly 16 times worse than in the disk of the galaxy. On the right, the EHR simulation enforces high resolution approximately to the virial radius, ensuring that interactions within the CGM are given much more computational attention. [Hummels et al. 2019]

What Does this Computational Cost Buy You?

By better resolving the gas in the CGM, the authors note that a number of physical effects present themselves. Firstly, the balance of cool and hot gas is shifted, leaving more cool gas and less hot gas than in simulations with lower resolution. The clouds of cool gas that form are also greater in number and smaller in size. Finally, the amount of neutral hydrogen and other low-energy ions found in the cool gas increases, while the abundances of oxygen, nitrogen, and neon in high-energy ionization states fall due to the decrease in hot gas. Coupled with the aforementioned work on AGN feedback, this can bring simulations closer to the observed abundances for these ions.

simulated galaxy and CGM

Figure 2: A galaxy and the CGM in an AMR simulation and an EHR one. A significant increase in HI (neutral hydrogen) can be seen in the EHR simulation. Recall that neutral hydrogen tracks the cool gas, which condenses into many clumps on the right that weren’t resolved in a traditional AMR simulation. Many of these clumps fall back into the galaxy because they no longer have enough thermal energy to resist the gravitational pull of the galaxy. [Hummels et al. 2019]

In other words, EHR causes more gas in the CGM to cool, condense into clouds, and potentially fall back into the galaxy. This is completely analogous to water vapor in our own atmosphere, which often cools, forms clouds, and rains back down to Earth. In this way, the CGM can be conceptualized as the atmosphere of a galaxy. Figure 2 shows cool gas condensing into these clouds, some of which fall into the galaxy.

Why does an increase in resolution result in more cool gas? The answer lies in how gas mixes in simulations. With lower resolution, clouds of cool gas are typically resolved only by a few cells, inducing artificial mixing between the hot and cool gas. The authors perform a test simulation demonstrating this, shown in Figure 3.

cloud test problem

Figure 3: In this test problem, a 4-kiloparsec-wide cloud of cool gas sits in a flow of hot gas for 260 million years. In the low-resolution test, the boundary of the cloud is only resolved by a few cells. This artificially thick boundary means that much of the cool gas quickly mixes with the hot gas and eliminates the HI (neutral hydrogen). In the high-resolution case, the boundary becomes much thinner, allowing the interior cool gas to survive much longer. [Hummels et al. 2019]

Resolution clearly makes a big difference in understanding the physics of the CGM and galaxies. For example, just like plants on Earth sprout after a rain, cool gas that condenses in the CGM and falls into a galaxy can trigger star formation. Understanding the ecology and geology of Earth requires a detailed picture of the atmosphere, and perhaps unlocking the mysteries of galaxy evolution may depend just as strongly on our understanding of the CGM. 

About the author, Michael Foley:

I’m a graduate student studying Astrophysics at Harvard University. My research focuses on using simulations and observations to study stellar feedback — the effects of the light and matter ejected by stars into their surroundings. I’m interested in learning how these effects can influence further star and galaxy formation and evolution. Outside of research, I’m really passionate about education, music, and free food.

cosmic distance ladder

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 Carnegie-Chicago Hubble Program. VIII. An Independent Determination of the Hubble Constant Based on the Tip of the Red Giant Branch
Authors: Wendy L. Freedman, Barry F. Madore, Dylan Hatt, Taylor J. Hoyt, In Sung Jang, et al.
First Author’s Institution: University of Chicago
Status: Accepted to ApJ

Author’s note: Credit for “H0tTake” goes to the conference “Tensions between the Early and the Late Universe” hosted at the Kavli Institute for Theoretical Physics!

The value of the Hubble Constant (H0) is a beast to pin down. However, it’s integral to our understanding of how the universe evolved and will continue to evolve. H0 relates the speed with which distant objects are moving away from us — due to the universe’s expansion — to how far away they are (see this Astrobite for a detailed explanation of how H0 assumed its place of importance). Measurements of H0 can be made using the early universe, from the cosmic microwave background (CMB), and the late universe, from distance measurements for stars, galaxies and other objects.

Under our current understanding of the universe, these two sorts of measurements ought to yield similar values of H0. Instead, we’ve witnessed a growing divergence between them that’s only gotten worse (or more interesting?) with time (see Figure 4, though it does come with a spoiler). Currently, early universe measurements of H0 rely on CMB observations made by the Planck satellite, while late universe measurements rely on Cepheid variables and Type Ia supernovae (Sne Ia). The discrepancy between these early and late measurements of H0 could be chalked up to new physics in the early universe that is outside our current models. But before claiming that, we’d want to rule out any hidden issues in how these measurements are being made.

On the side of the late universe, this requires using other astronomical objects to make measurements of H0  and to calibrate the distances to standard candles (objects whose brightness we understand very well), like Sne Ia. Very recently, a new measurement of H0 was announced, which used strong gravitational lens systems for distance calibration (see this Astrobite for a good summary). The paper being discussed in today’s Astrobite comes out of the Carnegie-Chicago Hubble Program, which was established to calibrate Sne Ia through alternate methods. Here, the authors use something called the Tip of the Red Giant Branch (TRGB).

The TLDR on the TRGB

color-magnitude diagram

Figure 1. A color-magnitude diagram of globular cluster Messier 55 (M55). The TRGB can be seen at the upper-right. [B.J. Mochejska, J. Kaluzny (CAMK), 1m Swope Telescope]

The TRGB consists of stars that are at a pivot point in their evolution. Red Giant Branch (RGB) stars are stars that have nearly exhausted the hydrogen in their cores. The next stage of their life is triggered when they start fusing helium in their cores instead. TRGB stars have just begun this stage of helium burning, and they can be distinguished by their characteristic redness and brightness (see Figure 1). These standard features of the TRGB make it highly suitable for measuring distances, since we know how bright it ought to appear at a certain distance.

The authors use the TRGB in lieu of Cepheids to calibrate the distances to galaxies that have hosted Sne Ia. TRGB stars have some advantages over Cepheids: they are much more common and can be found in uncrowded regions of their host galaxies, making them easier to identify. They also don’t need multiple observations to be recognized. Another useful quirk of TRGB stars is that their brightness in the I-band does not vary greatly with metallicity (the composition of the star), so the TRGBs in different galaxies shouldn’t look terribly different.

Could it (TRG)be?

In the near future, parallax measurements of Milky Way TRGB stars taken by the Gaia satellite will be available to anchor TRGB calibrations. For now, the authors use I-band observations of the Large Magellanic Cloud’s TRGB as well as parallax measurements for their analysis. The authors analyze the TRGB of 18 Sne Ia hosts, ranging from 7 to nearly 20 Mpc away, to calibrate the distances to those galaxies (see Figure 2). Their sample consists of galaxies that were not obscured by dust and had observations of their halos, where the TRGB could be cleanly measured. The TRGB calibrations were then used with a larger sample of Sne Ia to measure the distances to those Sne.

Sne Ia host galaxies

Figure 2. Nine of the eighteen Sne Ia host galaxies whose TRGB were studied. The squares represent the areas of the halo that were targeted. The hatched areas show the regions that were analyzed. [Freedman et al. 2019]

Finally, *drumroll* the authors present their measurement of the Hubble constant — 69.9 ± 0.8 ± 1.7 km s-1 Mpc-1 (the two errors are statistical and systematic respectively). This new result is shown clearly in a Hubble diagram showing their 18 TRGB calibrators and 99 Sne Ia from the Carnegie Supernova Project (see Figure 3). A Hubble diagram is a plot of distance versus speed, and the slope of the plot gives us a value of H0.

Hubble diagram

Figure 3. The Hubble diagram produced from the TRGB calibrators and Sne Ia from the Carnegie Supernova Project. The slope of the line is where the measurement of H0 comes from. The y-axis of the upper plot is the distance modulus (a measurement of distance using the relation between the absolute magnitude and apparent magnitude for an object). The y-axis of the lower plot is the difference between the points and the fit to the data. The x-axis of both plots is a quantity relating the distance modulus with redshift (see Section 7.1 of the paper). [Freedman et al. 2019]

This number falls squarely between the CMB and Cepheid-Sne Ia measurements of H0 (see Figure 4). The authors are careful to note that their result does not resolve the discrepancy in H0 values, but reiterate that additional, independent late universe measurements of H0 could change that. And the future is teeming with possibilities: aside from Gaia, the James Webb Space Telescope and LIGO and Virgo offer other avenues for measuring distances across large swaths of space, not to mention better measurements of strong lensing systems and tried-and-tested Cepheids. All in all, this is a very exciting time for cosmology!

Hubble constant over time

Figure 4. Measured values of H0 over time, showing where the TRGB measurement lands relative to the CMB and Cepheid measurements. The red star is the measurement from the paper being discussed. [Freedman et al. 2019]

About the author, Tarini Konchady:

I’m a graduate student at Texas A&M University. Currently I’m looking for Mira variables to better calibrate the distance ladder. I’m also looking for somewhere to hide my excess yarn (I’m told I may have a problem).

binary supermassive 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: Discovery of a close-separation binary quasar at the heart of a z ∼ 0.2 merging galaxy and its implications for low-frequency gravitational waves
Authors: Andy D. Goulding, Kris Pardo, et al.
First Author’s Institution: Princeton University
Status: Published in ApJL

With the announcement of the experimental confirmation of gravitational waves by LIGO in 2016 in tandem with additional electromagnetic follow-up of a neutron-star merger, astronomy was quickly ushered into an era of truly multi-messenger science. Although the number of gravitational-wave events observed by LIGO since is already substantial, the sheer number of black holes (and neutron stars) predicted to exist within our universe vastly outweighs this cumulative yield. One reason why LIGO is not constantly finding strong gravitational waves from all of these black holes (the gravitational wave background or GWB) is that not every black hole exists in a pair, which is a necessary condition to spiral inwards, merge, and set off a gravitational-wave event. Predictions suggest that the timescales required for some of these events to occur are a sizable fraction of the age of our universe! 

LIGO has, however, shown direct evidence for the merging of solar-mass black holes with atypically large masses between 10–40 M☉, and although easier to observe due to their strong gravitational signal, their existence has continued to challenge theoretical explanations. Scaling up to even greater masses does not reduce the pressure on theory either, as the predicted dominant mass contribution to the yet undetected GWB is from supermassive black holes (SMBHs) on the order of 108–109 M☉.

SDSS J1010+1413

Figure 1: Hubble Space Telescope images of SDSS J1010+1413, showing wide-field galaxy morphology and a zoom-in view of the central SMBH pair with the F621M, F689M, [OIII], and F160W bands. The galaxy shows evidence of a past merger: a disturbed shape and stripped gas streams. The [OIII] extent is seen to be coincident with the continuum F689M light. [Goulding et al. 2019]

It has long been hypothesized that SMBHs inhabit the central-most regions of almost all galaxies, and they accumulate mass through the slow accretion of gas and stellar material. When galaxies undergo the often violent processes of a wholesale merger, these SMBHs are predicted to collect within the central region of a galaxy and become gravitationally bound on short Myr timescales, accelerated by dynamical friction. However, the black-hole pair can only bleed off so much energy through interacting with nearby material, and at some point within the final parsec, the merger is predicted to stall out. This so-called “final parsec problem” has yet to be resolved. For the sample of intermediate-mass-black-hole mergers suggested by the LIGO observations, it appears that nature has ways around this problem. Whether this is true for SMBH mergers has yet to be seen.

Putting the final parsec problem aside, the authors of today’s astrobite provide definitive evidence for a precursor system that may one day produce a low-frequency gravitational-wave event consistent with a strong contribution to the GWB.

As shown in Figure 1, observations with the WFC3 instrument aboard the Hubble Space Telescope revealed a pair of tightly bound SMBH candidates in a highly luminous post-merger galaxy, poetically named “SDSS J1010+1413”. Accounting for the cosmological distance, the separation between the SMBHs is found to be ~430 pc (1,400 light-years). Previous studies examining the velocities and dynamics of the galaxy confirm the outward telltale signs of a trainwreck galaxy and provide the context for the apparently resolved pair of SMBHs in its core. Special imaging with an appropriate [OIII] narrowband filter corroborates this picture by defining the extent of the extremely luminous [OIII] emission — characteristic ionized gas associated with powerfully accreting SMBHs (AKA quasars). However, coincident X-ray observations with the Chandra Space Telescope showed very little X-ray light, a fact that, when compared to infrared estimates, suggests an obscuring cloud of thick gas along the line of sight.

Despite the awesome resolution of the Hubble Space Telescope (~0.04”), such an observation of a supposed SMBH pair may be ambiguous. To increase confidence in this interpretation, the authors modeled each of the sources with a coincident point-like model and an extended Gaussian model. Even so, a single SMBH with extended [OIII] and bisecting obscuration due to a dust lane could mimic this scenario, albeit with a significantly worse fit. Given the former scenario, each SMBH is estimated to have a minimum mass of 4 x 108 M☉ based on the Eddington luminosity limit, putting them in the sweet spot of GWB contribution.

merger stages

Figure 2: Merger stages with timescales shown. From the right, dynamical friction accelerates the first stage of coalescence, followed by stellar hardening. Gas infall may help prevent stall-out and overcome the “last parsec problem”. Then, in the final stages of coalescence at < 1 pc, Pulsar Timing Arrays should be sensitive to the predicted nanohertz gravitational-wave signal prior to the merger. [Goulding et al. 2019]

Lastly, the authors make tentative predictions for a future low-frequency high-mass gravitational-wave event, as shown in Figure 2. By considering carefully motivated arguments for dynamical friction, they surmise a coalescence timescale of 0.1–2 Gyr, where the lower limit is argued from the seemingly large gas reservoir near the SMBHs, which may help dissipate energy and rapidly close their orbital separation. Ignoring the “final parsec problem”, they argue that once the system reaches < 0.1 pc separation, the gravitational wave emission will enable the pair to finally merge within ~700 Myr. Given that the lookback time to this galaxy at z ~ 0.2 is on the order of the predicted merger timescale, this discovery provides strong evidence that such galaxies could be contributing to the GWB right now.

The low-frequency nanohertz GWB signal will not be detectable by current observatories such as LIGO. However, these increasingly powerful facilities will soon be complemented by hyper-sensitive Pulsar Timing Arrays which should be able to detect a nanohertz GWB signal, as may be produced by such a pair of quasars in this and other trainwreck galaxies.

About the author, John Weaver:

I am a PhD student at the Cosmic Dawn Center at the University of Copenhagen, where I study the formation and evolution of galaxies across cosmic time with incredibly deep observations in the optical and infrared. I got my start at a little planetarium, and I’ve been doing lots of public outreach and citizen science ever since.

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