Davor Krajnović's Recent Research Highlights

A list of recent highlights include:

Two channels of supermassive black hole growth

(appeared on the main page in January 2018.)

Suppermassive black holes are large dark objects residing in centres of galaxies. They are likely related to stellar black holes, which are the end products of the evolution of very massive stars, but the formation and evolution of supermassive black holes is still not well understood. They are called "supermassive" as their masses are well above million solar masses, up to tens of billions solar masses. There are also "supermassive" black holes with somewhat lower masses, and there is an expectation that also "intermediate" mass black holes exists (their masses would be of the order of tens of thousands solar masses), but these are not yet securely detected (a nicely controversial field to work on!). The supermassive black holes (I'll refer to them as SMBH from now on), next to their sole existence, have another puzzling property: their masses seem to be in proportion to the masses, and a number of other properties, of galaxies in which they live. The correlation between the SMBH masses and host galaxies are often called black hole scaling relations. These are very exciting as they connect objects which have very different sizes. SMBH, although very big, are comparable with the Solar System in size, perhaps 10 times as big. Galaxies, however, are very large objets. Their size is measured in 1000 parsecs (where 1pc ~ 3e13 km). So, there is something like 6-7 orders of magnitude in the size difference between SMBH and the host galaxy. Nevertheless, they seem to know of each other: large, bright and more massive galaxies have more massive SMBH (see Figure 1).

Figure 1. The tightest black hole - host galaxy scaling relation, relating the mass of SMBH and the velocity dispersion within an effective radius for a sample of nearby galaxies presented in van den Bosch (2016). Yellow symbols are early-type galaxies, while green symbols are spirals. Squares are galaxies with new SMBH mass estimates from a paper still to be published (Krajnović et al. 2018). The solid line is the best fit relation (from van den Bosch 2016), dashed and dotted lines show the scatter around the mean relation.

A simple way to explain this connection is that an SMBH and its host galaxy influence each other during their evolution. In other words, each "know" how the other grows. They could grow synchronous, or SMBH could grow faster, with galaxy catching up, or the other way around. This is one of the fundamental questions in modern astrophysics, especially, as the SMBH, in the form or Active Galactic Nuclei (AGN), can radiate huge amounts of energy, creating jets of plasma that pass through galaxies and illuminate the intergalactic space. This makes them as potential sources that regulate the growth of galaxies. But, what is meant by the "growth"?

Like humans, SMBH grow by eating. Black holes accrete matter via their strong gravitational fields. As the matter falls in the SMBH, it also forms an accretion disc around it, and the friction and complex interplay of friction, plasma physics and magnetic fields create the AGN phenomena. The more material falls in the SMBH, the more energy will be released, up to a point where the energy that is begin released starts blowing away the material (this is called Eddington luminosity, after Arthur Eddington who derived it to explain the physics of stars). So, SMBH can grow by "feeding" of the ambient material (gas and dust), but only until they blow it out.

Galaxies grow by increasing the number of their stars; they need to convert gas into stars. Similarly to SMBHs, they need to "capture" gas, which can come from another galaxy, or from what is called a "cold stream", a stream of gas from intergalactic space. Once captured, the gas has to cool and compress into giant dense clouds where the conditions are right for star formation. The more gas, the more stars can be made, but once stars are ignited, their radiation can heat the rest of the gas, evaporate the dust and stop further star formation.

This is why people making cosmological simulations like SMBHs and AGNs: they provide a lot of energy which can be used to stop galaxies forming stars, and evolve galaxies "passively" (without making new stars) the way observations are suggesting. The "tightest" black hole scaling relation (meaning the data points scatter the least from the mean relation) is between the SMBH mass and the "velocity dispersion" of the host galaxies (Figure 1). This velocity dispersion is a measure of the total binding energy of the galaxy. The way it is measured (by taking spectra within a certain large area of the galaxy) it actually includes both the ordered and random motion of stars and in this way it traces the gravitational potential as well as gives an estimate of the kinetic energy of the system. This is useful as it can be related to the Eddington luminosity and the energy feedback from the active SMBH which regulates the rate at which a galaxy can form stars and grow. So far so good, but the devil is in the details.

First problem is that very massive SMBHs (10 billion solar masses!) are found to exist when the Universe was something like 800 Myr old. These high redshift quasars (very luminous AGNs) pose problems how such big black holes are made in such a short time since the Big Bang. This is a fascinating problem (have a look at this semi-popular
article by Smith, Bromm and Loeb for possible solutions), but I'll not go into this further.

The other problem is that it is actually very difficult to find direct observational evidence for the connection between SMBH activity (AGN) and the suppression of star formation. And finally, the black hole scaling relations do not seem to be exactly universal, in the sense that all types of galaxies follow the same scaling relation. There are indication that different types of galaxies follow different type of scalings, that some galaxies are more strongly correlated than the others, suggesting SMBH influence different galaxies in different ways. (And not all astronomers agree on how one should interpret these varied scaling relations. For example, here and here are two recent summaries of the SMBH scaling relations, with not exactly the same conclusions.)

And then there is another channel of growth for galaxies: they can collide and merge. When this happens in galaxies without any gas, new stars can't be born. If the initial galaxies have similar masses, the final galaxy will then be twice as massive as either of the progenitors. And if they had SMBHs (as it seems is the case for every massive galaxy), their SMBHs will also merger, and again just double the mass. The interesting thing is, however, that when such a mergers between galaxies happens, because stars actually don't collide, the velocity dispersion of the final galaxy will remain the same.

Figure 2. Mass-size diagram, a non-orthogonal projection of the virial mass plane, for galaxies with measured SMBH masses. The mass of the galaxy is calculated from the total K-band luminosity of the galaxy, while the size is also measured in the K-band. Both quantities come from the 2MASS Observations, but the final numbers are obtained using various relations and values from van den Bosch et al. (2014), van den Bosch (2016) and Cappellari (2016). The diagonal lines are lines of constant velocity dispersion, and the red line is the so-called "Zone-of-Exclusion", beyond which galaxies are not normally found. Colour scale shows the mass of the SMBH. These values are adaptively smoothed using LOESS method (Cleveland & Devlin 1988). Note how along a dashed line the colour of the symbols is constant. This seems to be breaking once the mass of the galaxy M* > 2×1011 MSUN, when along a velocity dispersion line there seem to be a change of colour from orange to red, or red to white. Investigating this deviation, subtle as it is, is the topic of this "astro highlight" and this paper.

So, both the galaxy and SMBH mass doubles, but the velocity dispersion does not change. This implies that at some point, the relation between SMBH mass and the galaxy velocity dispersion, the tightest known black hole scaling relation, should break down. What are the evidence for this. Well, they are not rally clear, but its has been known for some time now that very massive galaxies, those that could expect to go through such mergers, actually have black holes with masses which often deviate from the predicted scaling relation. The problem is that there are not so many such galaxy-SMBH systems: they are simply rare in the nearby Universe, where we can measure their black hole masses. Nevertheless, we went to investigate this a bit further and checked where do galaxies with measured SMBH masses fall in the mass - size diagram (Figure 2).

In this diagram the lines of constant velocity distribution are sort of diagonal (as mass and the galaxy size are linked via the velocity dispersion through the Virial Theorem, and galaxies seem to obey it). The colours indicate the measured black hole masses, smoothed using a statistical technique called LOESS (for some practical details have a look here). And indeed, if you are a galaxy with mass less then about 2e11 Msun, your black hole will sit where is it expected based on your velocity dispersion. But once you are a more massive galaxies, the black holes don't seem to follow this trend.

This can be explained using a simple toy model (Figure 3) where the growth of black holes has two channels. For lower mass galaxies, SMBHs grow through accretion of gas, and they regulate the growth of the host by energy feedback. This is why SMBH masses correlate with the velocity dispersion of galaxies. For higher mass galaxies, which do not have any gas, and their SMBH can't grow by accretion, nor galaxies can grow by star-formation, the channel of growth is different: they grow by merging. This, however, changes the dependence of the scaling relations, and SMBH mass does not correlate anymore with the galaxy velocity dispersion.

Figure 3. A toy model simulation of SMBH mass on the mass-size plane (left-hand panel), and the same model with over-plotted LOESS-smoothed black hole masses (right-hand panel). The toy model SMBH is based on a simple prescription in which SMBH masses are calculated based on the black hole mass - velocity dispersion relation for galaxies less massive than 2 × 1011 MSUN, and on a black hole mass - velocity dispersion relation modulated by the galaxy mass for higher mass galaxies. The colour scale of the model is limited to the same range as the LOESS-smoothed black hole masses of our galaxies, as shown on the colour bars. The red solid lines show the Zone-Of-Exclusion. Diagonal dashed lines are lines of constant velocity dispersion.

The model works well qualitativey, as can be seen from the right-hand panel of Figure 3, but the sample is still small. This is of course good, as to confirm/disprove the proposed conjecture, we need more data and more telescope time! Crucially, these usually also generate more ideas. For the moment, however, this model is also consistent with our understanding of the evolution of massive galaxies. The most massive galaxies, those that reside in the centres of groups of clusters of galaxies, are devoid of gas and can't grow via star formation. They however, capture smaller satellite galaxies, and by merging with similar size galaxies (this happens rarely, but it seems that is does happen often enough to be a factor). Such mergers of massive galaxies don't only change the black hole scaling relations, but also the internal structure of galaxies and the way stars orbit in them. So, for example, lower mass galaxies have ordered rotation, similar to those of spiral (disk) galaxies, but the most massive galaxies are all rotating very slow, and show very irregualr motions of stars (have a look at a few past "highlights"). Or, the most massive galaxies typically have something astronomers call a "core", a nuclear region where there is less light (less stars) than expected. The only model we have at the moment that can explain such cores are actually mergers of SMBHs, that kick the surrounding stars in the process of spiralling in. We haven't seen such a merger yet, but gravitational waves should find it, when the instrument become sensitive to it.

Mass assembly of NGC3115

(appeared on the main page in January 2017.)

One of the standard galaxies to target with new instruments is
NGC3115, and inevitably we observed it with MUSE during the commissioning. It is a large galaxy on the sky, with a half light radius of about 44 arcsec, and as such it was used as a test of making a large mosaic with MUSE without using an optimal strategy for the sky fields (these are necessary to take out the contribution of the Sun spectrum from the object) and general calibrations. It was envisaged as a quick and dirty try on a target with lots of available imaging to test MUSE image quality and mosaicing capabilities. Seeing the data in 2014, we decided that they are worthy of more then just a small performance test for MUSE. The results is a wonderful paper led by Adrien Guérou with a title: "Exploring the mass assembly of the early-type disc galaxy NGC3115 with MUSE".

Figure 1. Image of NGC3115 and the 4 MUSE pointings mapping the disc.

NGC3115 was observed with 4 adjacent MUSE pointings tracing the extent of its disc (Fig. 1). To maximise the coverage, MUSE was oriented such that the disc was along the diagonal of the square(ish) field-of-view, hence the toothy pattern of the data. The quality of the data are best seen in Fig. 2, which shows the stellar kinematics, from top to bottom: the mean velocity, the velocity dispersion, and the higher order moments of the line-of-sight velocity distribution. The stars in this galaxy rotate fast and in a regular manner. There is a rather thin disc (also visible on the images), but the rotation is present above the disc as well. It seems the whole galaxy is spinning rapidly.

Figure 2. Maps of the first four moments of the line-of-sight velocity distribution: (a) the mean velocity, (b) the velocity dispersion, (c) the third order Gauss-Hermite moment, and (d) the fourth order Gauss-Hermite moment. The colour bar and the numbers around it give the range of the values on each map. Note the thin and fast rotating structure in the mid-plane of the galaxy, corresponding to the optical thin disk. The velocity dispersion is small in the disc and rises only in the very central region.

MUSE can give us more than just the motions of stars. We can also investigate the chemical properties of the stars, and in comparison with models of stellar populations (and stellar evolution) can give us a decent handle on how old are the stars. This is called the star formation history and it is shown in Fig. 3., where Adrien separated stars by age in six panels. On each you can see the distribution of stars of a certain age (as written in upper left corner). Each coloured dot on these panels represent a region in NGC3115 where the total mass of stars (in solar masses) of a particular age is as given by a value that can be read of the colour bar on top (and the numbers in the lower left corner). There are three important conclusiones here: 1) most of the stars are very old, born within the first few giga years of the age of the Universe. But, 2) stars have been forming in this galaxy ever since, not much, but continuously, and 3) they were born in the disc like structure. The distribution of stars of different ages in different structures is a very interesting results. It shows that NGC3115 is made of multiple structural components (which is also suggested by the imaging and kinematics), and it seems these different components can be age tagged, enabling to date their formation and therefore follow the (structural) evolution of this galaxy.

Figure 3. Maps of the stellar mass distribution in NGC3115. These maps show the present day stellar mass distribution as a function of the (mass-weighted) age. Each pannel shows the mass in stars of a certain age (as indicated in the upper left corner), thile the total stellar mass within a given age range is given in the lower right corner. The difference in the distribution of the stellar mass between panels (e) and (f) suggest a different type of mass assembly: for age>12 Gyr the distribution is more centrally concentrated and therefore it can be linked with a dissipative process (such as a collapse of a gas cloud). Panles (a) to (c) suggest an extended star formation within the disc of NGC3115.

The last figure in this highlight shows what can be found out when you combine both ages of stars and their chemistry (or as astronomers like to call it: metallicity, which measures the fraction of say iron with respect to hydrogen). The metallicity is an important tool as metals (all elements except hydrogen and helium) are only made in stars or supernova explosions, and, hence, more metal rich the stars were born from material that was already pre-processed in stars that are no more. Fig. 4 shows that young and metal rich stars are found in the thinnest structure (how much is something thin or thick depends how far it extends above the midplane of the galaxy), the thin disc, which is also most rapidly rotating (Fig. 2). Young but metal poorer stars are born in a thicker component. This components is similar to the thick disc of the Milky Way, a still rapidly rotating structure, in which stars have drifted away from their original birth place within the thin mid-plane of the galaxy due to mutual gravitational interaction. Finally, there are also old stars of different metallicity. Those that are poor are in an extended spheroidal structure that surrounds NGC3115, while those that are metal rich are in the very centre of the galaxy.

Figure 4. Stellar mass fraction of NGC3115 in four bins of age and metallicity pairs: young (age<12 Gyr), old (age>12Gyr), metal-rich ([Z/H]=0,2) and metal-poor ([Z/H]<0.2). The old and metal-poor stars in NGC3115 are mostly located in the spheroid, while the young and metal-rich stars are mostly in the disc.

These fascinating MUSE data can be used to piece a picture of the formation and subsequent evolution of NGC3115. The galaxy was likely created some 1-2 Gyr after the Big Bang (redshift of ~3), from a collapse of a gas cloud. Most of stars were born in that event. Then over next 1-2 Gyr (to redshift of about 2) this proto galaxy went through a few mergers with other similar systems, containing already chemically enriched gas. Finally, over the last 10 Gyr, NGC3115 was evolving passively, meaning without any large galactic collision with a similar size galaxy (although accretion of smaller galaxies were possible at larger radii beyond the MUSE data). This evolution was dominated by the so-called secular processes: the stars age and die and new stars are born from the stellar material (and not some new gas from outside of the galaxy), while structure of the galaxy is shaped by processes associated with disc instabilities such as bars and spiral arms (you'll have to look at Adrien's paper to check out those).

Are kinematically distinct cores really decoupled components: the case of NGC5813

(appeared on the main page in June 2015.)

One of the first targets with MUSE, still during the early stages of commissioning of the instrument on Paranal, was an elliptical galaxy
NGC5813. Why? Because it intrigued me for a long time... In 2003, we published a paper, led by Eric Emsellem and the SAURON team, showing integral-field kinematics of a representative sample of early-type galaxies. There were many famous object in that sample, and this one had another distinction: NGC5813 was the first known object to have a "kinematically distinct core", discovered by Geroge Efstathiou, Richard Ellis and David Carter in 1980. The SAURON data were very good, but they didn't reveal all secrets of NGC5813. The structure of the NGC5813 remained a mystery for good 35 years. The new MUSE data, however, allow us to start understanding the hidden nature of NGC5813 and provide a good topic for An Astro Highlight based on this paper.

Kinematically distinct cores, or kinematically decoupled cores, or kinematically peculiar cores, or most commonly only referred to as KDCs, are exciting features on the mean velocity maps of galaxies. As you can see from the first panel on the plot below, which shows the map of the mean velocities of stars in NGC5813 (essentially a map of ordered motions of stars), the centre of the galaxy, its core, seems to rotate, while the rest of the galaxy's body does not seem to have any ordered rotation. Why is that? How do you make structures like that? Are they stable and long lived?

The mean velocity and the velocity dispersion maps of NGC5813 as seen by MUSE, showing the famous KDC, but also a very specific structure on the velocity dispersion map (two peaks along the major axis). This figure is based on Fig 2. from a recently published paper in MNRAS (Krajnović et al. 2015). NGC5813 is the first know KDC, and the MUSE data, constraining dynamical models, allowed us to understand its internal structure. It is not a decoupled component, separated from the rest of the galaxy's body. Instead, it is made of two components of stars, each following a family of orbits with opposite net angular momentum: they are counter-rotating. This counter-rotation together with a specific distribution of stars within the components defines the structure we see on the kinematics maps. North is up, East is left. A specialist talk about KDCs in general, which I gave at the Macquarie University, Sydney, in February 2015, is available here (27Mb).

These are some of the question astronomers are trying to answer since the discovery of KDCs. There are many types of KDCs (velocity maps of a number of KDC you can see on the figure at the bottom of the page): some have similar velocity maps as NGC5813, where the core seems to rotate and the rest of the body doesn't, but there are even more spectacular cases, such as NGC4365, where the rotation axes of the core and the body are misaligned by about 90 degrees, or NGC5322 where the core is counter-rotating with respect to the rest of the body (180 degrees misalignment). Sometimes, but not very often, the misalignments is even different from the 90 and 180 degrees (can you spot such galaxy on the figure below?). When astronomers started discovering more and more KDCs in 1990s, they became the topic of the day as they look like kinematic proofs that galaxies indeed grow by interactions. KDCs are the smoking guns of mergers, but what kind of mergers and how are they actually created is still a bit of an uncharted territory (although there are plenty of ideas out there).

KDC are often found in large and massive elliptical galaxies. These objects are essentially balls, or even rugby balls, of stars embedded in dark matter halos and hot gas. Some have active galactic nuclei and large scale jets, which seem to stir the halos of hot gas, rising bubbles of gas around the central galaxy (all this is true for NGC5813). KDC are also found in smaller galaxies, but these seem to be different from the KDCs in big galaxies: they are made of young stars and they are likely to disappear as their stellar population ages, as it was elegantly shown by Richard McDermind and collaborators in this paper.

On the second panel of the figure above you can see what makes NGC5813 very special indeed. It shows the map of the velocity dispersion, a measure of random motions, and it is very unusual. Typically the velocity dispersion maps of early-type galaxies (including galaxies with KDCs) have a peak (highest values) in the centre of the galaxy, there the gravitational potential is the deepest. From that point outwards, velocity dispersion decreases pretty much steadily, following the distribution of the light. But the velocity dispersion map in NGC5813 is very different: it has a central peak, but this is then followed by a drop of some 100 km/s, followed by a rise of about 40-50 km/s and a subsequent drop. While general drops and rises in velocity dispersion are known to occur, typically in galaxies with small central stellar discs or bars, the structure of the velocity dispersion map in NGC5813 is very specific. The rise happens only on two sides of the nucleus (and not everywhere around it), along the major axis, and it coincides with the end of the KDC (of the region that exhibits clear ordered rotation). What could be the cause for that?

As I mentioned earlier, NGC5813 intrigued us for quite some time. The drop in the velocity dispersion was known since the discovery of the KDC, and the SAURON observations revealed part of its structure. But unfortunately the field-of-view of SAURON was too small, and the final mosaic of three SAURON pointings didn't cover the full structure (we didn't really know what to expect). The MUSE, however, did it in one shot. It revealed this special structure, which has been actually seen before, but in very different types of galaxies.

A small percentage of early-type galaxies, something like 7%, typically less massive and luminous and quite flat, disc-dominated systems, show what we call a double-sigma profile on their velocity dispersion maps: drop in the centre, followed by two peaks in the velocity dispersion along the major axis. Their velocity maps can be diverse: from very nice counter-rotating KDCs (180 degrees difference in angle between the central and the outer part of the body) to very messy velocity maps, with no ordered rotation. The most famous of these galaxies is NGC4550, and there is a very reasonable explanation for its structure (see this as well): it is made of two discs of stars, which rotate in opposite directions. The two peaks in the velocity dispersion are then the consequence of the counter-rotation. Imagine that in the same spectrum you are detecting one population of stars that go in one direction (quite fast) and another population that goes in the opposite direction (as fast). This will result in the mean observed velocity close to zero (the relative velocities cancels out), but the spread in velocities (the velocity dispersion) will be large. Depending on the type of these two counter-rotating discs, how many stars they have and how are they distributed, one can get a plethora of structures on the velocity and the velocity dispersion maps.

Once we saw the shape of the velocity dispersion maps of NGC5813, we had a good hunch that that NGC5813 could also be built of two populations of stars, which rotate in oposite directions. But NGC5813 is not a flat, disc-like galaxy, and the situation is a bit more complicated. Still with the help of dynamical models, we found that indeed, the galaxy can be described as having two components of stars. They couter-rotate, but not very fast; neither of them is a disc (as in a typical double-sigma galaxy). The distribution of stars is such that one component is dominant in the centre: within the KDC one component consists of 70% of the stars - that is why we see the rotation. Outside of the KDC, the components become more similar, contributing about 50-50% in mass. This 50-50 ratio also explains why there is no visible rotation: stars in both components steadily rotate, but the net motion gets canceled and we don't see it.

What does that mean for formation of the KDCs? Well, it seems KDC in NGC5813 should not be considered as a separate entity, divorced from the rest of the galaxy. It is not a ball of stars that has different kinematics from the rest of stars. It is not "decoupled". It is a result of mixing of different orbital families of stars. There are still many ways you can form this structures, such as accretion of gas from cosmological cold flows, turbulent dics of gas rich galaxies in the early universe (stars in NGC5813 are almost as old as the Universe), mergers of similar-in-mass galaxies.... Plenty of possibilities, and only by understanding the structure of other KDCs in a similar way as here, we can hope to understand their formation.

NGC5813 is a wonderful system. It has a complex kinematics which can be explained with a relatively simple, but very specific dynamical model. It also has very complex multi-phase gas structure, linked to the active galactic nucleus and its jet. But that might be another highlight, in the future.

The mean velocity maps of galaxies with KDCs from ATLAS3D Paper II (Krajnović et al. 2011). All velocity maps were obtained with the SAURON instrument. One before the last in the 2nd row is the SAURON velocity map of NGC5813, also published in 2004. North is up, East is left.

Kinematically distinct core in M87

(appeared on the main page on September 2014.)

Here is a highlight of my recent work based on the MUSE Science Verification data obtained at the end of June 2014 (past highlights can be found
here). It comes from a paper by Eric Emsellem (ESO), Marc Sarzi (University of Hertfordshire) and myself (Emsellem, Krajnović & Sarzi, 2014, MNRAS, accepted). It spawned press releases in Germany and the UK, and it featured as a Picture of the Week by ESO.

Multi Unit Spectroscopic Explorer, or MUSE for short, is a next generation instrument on ESO's Very Large Telescope situated on Cerro Paranal in Chile. It was installed at the beginning of 2014 on UT4 and we were commissioning it throughout the spring. The first light and some of the commissioning data can be seen on this ESO press release. The standard procedure with the new instruments, before they are offered to the astronomical community, is to demonstrate their scientific capabilities. This is done during the Science Verification (SV) observations. The idea is that astronomers devise programmes which will test the instrument and all the tools that are required to use it, and bring some exciting science. Once MUSE was successfully commissioned, ESO announced a call for the SV proposals, selected a bunch of sound looking ideas and started observing in June (the SV run was about 2 weeks long, split between late June and mid August). The catch with the SV data is that they become immediately public and anyone in the world can use them.

We observed one of the biggest nearest galaxies. Its name is Messier 87 (M87 for short), or NGC4486 (if you prefer the New General Catalogue of Nebulae and Clusters of Stars of John Dreyer), and it is in the centre of the Virgo Cluster of galaxies, some 54 million light-years away. Galaxies like M87 are fascinating from many aspects. They are the largest galaxies, reside in the centres of galaxy clusters (bottoms of their gravitational wells), consists of old stars, have very massive black holes in their centres, which are typically very active, shooting out plasma jets and steering the inter-cluster medium very far beyond the extent of their host galaxies. Galaxies like M87 are often considered the end products of galaxy evolution, but the question still is: how have they become what they are now?

This is a copy of Fig. 2 from Emsellem, Krajnovic & Sarzi (2014). The first panel shows the SAURON velocity map, similar to the one presented in Emsellem et al. (2004), but now binned to a much higher signal-to-noise ratio (of 300, instead of 60 in that paper), which suggested that there is indeed some rotation in the central parts. The second panel presents the MUSE velocity map, and the third panel is a kinemetry reconstruction of the features on the middle panel, highlighting the KDC and the prolate-like rotation of the outer body. The color-bar on the right describes the rotation pattern: red are those regions that are receding from and blue are those that are approaching the observer. Green are the regions with no-net motion. The recession velocity of M87 was subtracted. A dashed magenta line marks the photometric major-axis of M87 (at large radii). The central 2" were masked as these kinematics are significantly influenced by the AGN (having very broad and strong emission-lines). North is up, East is left.

The first time I saw a map of stellar velocities of M87, the one coming from SAURON integral-field spectrograph observations, but not as good as the one shown on the left most panel on the figure, I was puzzled: why are stars not moving? (Have a look at the collage at the bottom of the page of stellar velocity maps typical of the galaxies in the nearby universe). The answer is that they are moving, actually very fast (often faster than a million km/h), but they do not move in a coherent way. Basically every star is moving in a direction different from that of the other stars and the mean motion observed at each point (what is plotted in the figure above) is very low. There was, however, something odd in the SAURON velocity map of M87 and we decided to have a second look at the centre of this galaxy, with an instrument of better characteristics. MUSE has a larger field-of-view than SAURON, covering about 1 arcmin x 1 arcmin on the sky. It also has a beter resolution, both in terms of probing smaller spatial scales (one square pixel is 0.2 arcseconds on a side), as well as in the resolving power of the spectrograph. Basically, with MUSE one can distinguish better than with SAURON between the various velocities of the stars. The improvement can be seen on the middle panel of the figure below.

The MUSE velocity map of M87 (middle panel above) is somewhat larger than the previous SAURON map. Notably it adds information in the top-left and bottom-right corners, or, in astro terms, in the North-East and South-West regions. This part, which was not covered by SAURON, is actually very important, as it shows that stars in M87 do coherently rotate about the center of the galaxy (evidence of this could also be seen in this paper). Another unexpected finding brought up by MUSE is that the stars in the very centre also rotate coherently, but their rotation is distinct from the large scale rotation (seen in the corners). It is distinct in two senses: the two rotations are not spatially connected and are around different axes, with some 140 degrees difference between them. The axis of the outer rotation is actually close to be aligned with the major axis of the galaxy (dashed magenta line above), the main symmetry axis which traces the orientation of the galaxy. There are very few galaxies that show this kind of rotation, often described as being "prolate-like", as a body with a prolate symmetry is expected to exhibit such a rotation.

Similar structures are know to exist, typically in massive elliptical galaxies. They are called "kinematically distinct cores" (KDCs). What makes this kinematic configuration even more surprising is that neither of the axis of rotation coincides with the minor axis of the galaxy. That kind of configuration is rarely found, and it very strongly suggest that the body of M87 is of a triaxial shape. As a matter of fact, now that we have detected the coherent rotation, albeit low, we can have some hope in determining the actual shape of M87. Before, with essentially zero net rotation, many models of very different shapes could be made to describe M87 (its light distribution and no rotation), but to hardly any one could attach a decent significance. Now, we have a possibility to change this.

To summarise, MUSE observations of M87 revelad that this central galaxy in the Virgo cluster has a very complex kinematic structure, comprising of an outer prolate-like rotation and a central KDC. The two rotation patterns seem distinct and with a mutual orientation of about 140 degrees. How does one build a galaxy like this? M87 is a product of a long history of galaxy mergers, some of which were between galaxies of smilar masses and most of which were between a big galaxy (M87) and many small satellites. By knowing the actual shape and kinematics, by knowing the chemistry and age of the stars, but also the properties of globular clusters and surrounding satellite galaxies, we may start deciphering the colourful history of M87.

And MUSE is a really impressive instrument!

ATLAS3D Paper 23. Angular momentum and nuclear surface brightness profiles

(appeared on the main page on 01.09.2013.)

Here is a highlight of my recent work based on ATLAS3D data. These two diagrams are from
Krajnović et al. 2013 (MNRAS, 433, 2812K) and with colours and symbols show the nuclear shape of the radial surface brightness profiles in the λR - ε space of nearby early-type galaxies. (λR and ε respectively measure the specific angular momentum of stars and the apparent shape of galaxies within the half-light radius). The shape of the light profiles can be related to the formation of galaxies. They either continue increasing at the last resolution point (typically about 10 pc for galaxies in this study) or they turn downwards at some point and stay flat until the last observable point (so, it is critical to use Hubble Space Telesope observations). The nuclear regions in galaxies where the light profile turns down and stays flat are called "cores". The other type of light profiles, those that keep increasing at the smallest spatial scales we can probe, are sometimes called "core-less" (or power-laws). The idea is that cores are made by massive black holes which are bound in a bianry system and kick far out all stars that cross their paths. These black holes are brought together in a merger of two galaxies of simular sizes. However, if there is a lot of gas present during the merger (i.e. brought by one of the galaxies), even if cores are created by the massive black hole binary, new stars will be made from the gas and the light profile will be core-less.

There are several ways to determine which galaxy has a core or a core-less profiles. One can, for example, fit an outer piece of the light profile with a function (such as the Sersic 1968 profile) and then estimate how much does the nuclear light profile deviate from this outer fit (e.g. Kormendy et al. 2009), in other words shows an "excess" (e.g. core-less) or a "deficit" (core). Anotehr way is to fit a more specific function, such as the so-called core-Sersic function (Graham et al. 2003). In this work, we used the so-called "Nuker-law" (Lauer et al. 1995), which is a double power-law function meant to be applied only to the inner regions of the light profiles (not on the full galaxy). Each method gives somewhat different results (as it measures different things), but, as we show in the paper, they are not very different.

The two diagrams above show the angular momentum versus the ellipticity for 260 ATLAS3D galaxies. On both panels, small open symbols are galaxies with no available HST observations (we do not have data to investigate their nuclear light profiles) and filled small symbols are galaxies for which the classification was not possible (it is uncertain due to significant presence of nuclear dust which masks the central light). Also on both panels, colours of symbols indicate the class of the nuclear profiles: red -- core, blue -- core-less. The green solid line separates fast from slow rotators (divsion of early-type galaxies into slow and fast rotators is described in Emsellem et al. (2007) and Emsellem et al. (2011) papers). Left: Core galaxies are shown with squares and core-less galaxies with circles. The grid of dashed and dotted lines show the region where one can expect to find galaxies that are oblate and axisymmetric, and of certain anisotropy in the distribution of velocities (more details can be found in the paper by Cappellari et al. 2007). Right: Shapes of symbols indicate the kinematic group (Krajnović et al. 2011): group a -- non-rotating galaxies, group b -- featureless non-regularly rotating galaxies, group c -- kinematically distinct cores, group d -- 2σ galaxies made of two counter-rotating discs, and group e - regularly rotating (disc-like) galaxies. Kinematic classification is not provided for galaxies with no HST data. The contours show the distribution of a family of oblate objects with an intrinsic shape of εintr= 0.7 ± 0.2 (from Emsellem et al. (2011)).

Conclusion: fast rotators are typically core-less while slow rotators are cores, but there are significant exceptions to this rule, with an implication that either the accepted formation of cores may not be the only path for their formation, or that there are some additonal core perservations mechanisms at play. Perhaps the most interesting finding is that there are some core-less slow rotator galaxies (blue symbols below the green line), pointing to more complex formation of these galaxies (i.e. formation in gas rich mergers). Unfortunately, a significant number of slow rotators which could be in this regime has not HST imaging. This will be changed soon, as we were awarded HST/WFC3 observations during the Cycle 21 of 12 galaxies in this region, which will allow us to determine the formation paths for these galaxies.

ATLAS3D Paper 17, Fig.5

(appeared on the main page on 17.09.2012.)

Bellow is a highlight of my recent work based on ATLAS3D data. The two diagrams are from
Krajnović et al. 2012) and show the relation between λR, a proxy for specific angular momentum of stars within the effective radius (enclosing half of the total galaxy light), and ε, the measure of the apparent shape, for a complete and volume limited sample of nearby early-type galaxies. Galaxies that are above the green line are fast rotators (with high angular momentum) and galaxies below the line are slow rotators (with low angular momentum). In the traditional classification of early-type galaxies, fast rotators are basically S0s, while slow rotators are true ellipticals (for details see Emsellem et al. 2011, MNRAS,414, 888).

Symbols on the left panel are related to the Sérsic index (a measure of the curvature of the light distribution) of the bulge component (galaxies were decomposed into "bulges" and "discs", but this didn't work for all: red symbols are single components fits), while the colours relate to the fraction of disc light in the these galaxies. The right panel is similar, but the symbols are now related to the type of motions of stars found in these galaxies (and explained in Krajnović et al. 2011, MNRAS, 414,2923). Conclusion: a) most of early-type galaxies have a significant fraction of stars in discs, b) Sérsic indices of bulges of early-type galaxies are often small (nb<3), c) fast (ordered) rotation is related to existence of discs, d) slow rotators (ellipticals) actually have large Sérsic indices and often can not be decomposed into two components.

Last modified: 12 Sep 2014