Lecture Three: Observed Properties of Galaxies, contd.! Hubble Sequence. Environment! Globular Clusters in Milky Way. kpc

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1 Hubble Sequence Lecture Three: Fundamental difference between Elliptical galaxies and galaxies with disks, and variations of disk type & importance of bulges Observed Properties of Galaxies, contd.! Monday 15th Feb Longair, chapter 3 + literature Early type Late type Hubble 1936, the Realm of Nebulae Environment! Globular Clusters in Milky Way Mateo 2008, Garching workshop kpc ~140 globular clusters, 65% <8kpc from centre

2 Globular Clusters & galaxy formation and evolution vrot = 193 +/- 29km/s!los = 59 +/- 14km/s" vrot = 43 +/- 29km/s!los = 116 +/- 11km/s" Metallicity dispersion is large; mean metallicity decreases with increasing distance from galactic centre metal rich Armandroff 1989 AJ metal rich (disk) Globular Clusters & galaxy formation and evolution metal poor (halo) milky way disk metal poor flattened from rotation dominated by random motions Zinn 1985 ApJ, 293, 424 Zinn 1985 ApJ, 293, 424 Formation of Halo? Outer Halo: dsph kpc Mateo 2008, Garching workshop Bullock & Johnston 2005 ApJ

3 The Local Group Nearest Cluster! hello Near centres of mass: gas-less pressure supported dsphs Anomalies: more distant dsph Outer regions: dominated by gas rich quiescently evolving dwarf irregulars kpc Mateo 2008, Garching workshop Local Super-Cluster! What is a galaxy cluster? Half the galaxies in the Universe are found in clusters or groups, systems of galaxies that are a few Mpc across. Within the central Mpc, clusters typically contain luminous galaxies (L> L * ~ 2 x L! ). Most famous catalogues: Abell 1958 and it s 1989 supplement, with 4073 rich clusters, having at least 30 giant members within a radius of ~1.5h -1 Mpc. Galaxies in clusters are bound together by their mutual gravitational attraction: the cluster is generally filled with hot interstellar gas, also retained by gravity. Clusters differ from groups by having higher densities. Cluster galaxies live in such proximity that they significantly affect each others development.

4 What causes diversity of galaxy types? Virgo Cluster Velocity dispersion 715km/s Virial radius 730kpc ~17Mpc distance There are a number of ways of reducing the variables in a study of galaxy properties and one is to look at a group or cluster of galaxies. You remove uncertainties due to different distances of your sample of galaxies as well as different environments. Your sample completeness is easily defined. HOWEVER it is not clear that you obtain a complete sample of all types of galaxies, and it may not even be a good average sample. closest rich cluster of galaxies, centred on giant elliptical galaxy M87 Virgo Cluster! Brighter galaxies are redder Global Properties! Elliptical galaxies in Virgo (open symbols) & Coma (closed symbols) Coma galaxies are shown 3.6 mag brighter as they would be at distance of Virgo This trend could be explained if small elliptical galaxies were either younger or more metal poor than large bright ones (or both). ~17Mpc distance, ~2000 member galaxies Bower et al. 1992

5 What do colours mean?! Spectrum of an Elliptical galaxy! U B V R I What does it mean?! U B V R Star Formation History >10Gyr ~8Gyr ~1.5Gyr ~5Myr Stellar spectra Elliptical Galaxy O Connell 1986 PASP, 98, 163

6 Velocity dispersion 148km/s Virial radius 880 kpc Ursa Major Group! Galaxies get bluer and fainter Global Properties! On average S0 galaxies are luminous and red Sd, Sm systems are fainter and bluer Only late type galaxies with no particular concentration towards any centre Verheijen & Sancisi 2001 Ursa Major Group M. Verheijen 1997 Fainter galaxies have proportionately more HI Global Properties! Virgo Cluster! M HI /L K (solar units) Open circles, low SB galaxies (I K (0) > 19.5), the least luminous and richest in HI; not efficient at turning HI into stars Disk has lower central surface brightness Ursa Major Group M. Verheijen 1997 VIVA: VLA Imaging of Virgo in Atomic gas

7 Inter-cluster Medium! Virgo Cluster ngc4402 falling into centre of Virgo ngc4522 in Virgo Ursa Major Group vs. Virgo! Fraction of E & Sp HI properties of late-type galaxies HI / optical diameter HI content of galaxies in centre of Virgo LESS than in the outskirts or in lower density systems (Ursa Major) outskirts HI deficiency Gas disks are SMALLER in centre of Virgo than in the outskirts or in lower density systems (Ursa Major) centre Distance to cluster centre (degrees) Verheijen 2004 Oemler 1974

8 Morphology-Density Relation First large (55 clusters, 6000 galaxies) study of morphological segregation (Dressler 1980). The frequency of different galaxy types was found to vary as a function of the number density of galaxies in which they are found. Is this related to R? Or N? Difficult to ascertain: N # R -1. galaxy type appears to be dictated by LOCAL DENSITY of galaxies, although presumably galaxies move through a range of densities, thus there must be coherent sub-structure. study of poor groups (Postman & Geller 1984) - the centres of which have similar densities to outer regions of clusters follow same relations as clusters. galaxies with a nearby companion are more likely to be Es (Whitmore, Gilmore & Jones 1993), so morphology-density a local phenomenon. there is a single universal morphology -density relation over 6 orders of magnitude in density. What causes diversity? Galaxies in clusters more likely to be Es or S0s than those in the field Environment plays a role Not all clusters are the same - large E fraction correlates to regular symmetric clusters; low values to ratty ones Oemler (1974) Also E/Sp varies with position in a cluster -> depends on density. Fraction of spirals increases out from centre; essentially no spirals in cluster cores MORPHOLOGY-RADIUS RELATION Spirals closer to the centre have less gas than those further away Why? Spatial segregation should give rise to kinematic differences - ie., spirals follow more energetic orbits - ie., spirals at a given distance from centre of cluster should have larger random velocities than E Fraction Sp/E goes up moving out from cluster centre Dressler 1980 fraction of Sp goes down with size of cluster. Effect of sub-structure? Simulating Interacting Systems Lack of spirals compared to ellipticals in dense environments has lead people to consider that merging spirals result in an elliptical galaxy... Galaxy luminosity function Just as the distribution of stellar luminosities reflects the physics of star formation and stellar structure, we might hope to learn about galactic evolutionary processes by studying the distribution of galaxy luminosities. The galaxy luminosity fn. $(M), $(M)dM is proportional to the number of galaxies that have absolute magnitudes in the range (M, M+dM): Where % is the total number of galaxies per unit volume The field galaxy luminosity function, in its simplest form, involves measuring the apparent magnitudes of all the galaxies in some representative sample. The individual brightnesses are converted to absolute magnitudes by estimating the galaxies distances usually by applying the Hubble law to their observed redshifts. Josh Barnes 1998 Toomre & Toomre 1972, ApJ

9 Short comings Malmquist bias - magnitude limited surveys - luminosity function distorted if function has a finite spread in luminosity. Even if all galaxies have intrinsically identical luminosities, but a range of estimated absolute magnitudes due to errors in their adopted distances. Estimating distances using hubble law intrinsically approximate process. Particular problem for nearby galaxies - local motions dominate over hubble flow. Particular problem for low luminosity galaxies - which can only be observed nearby. So faint end of luminosity fns remains rather poorly defined. Spatial structure - incomplete sampling of variations in galaxy distribution (filaments vs. Voids). Luminosity Functions of galaxies In an attempt to find a general analytic fit to galactic luminosity functions, Schechter (ApJ 203, p297, 1976) proposed the functional form: Which can also be written (in terms of astronomical magnitudes): In both forms & (the slope of the power-law at low luminosities) and L * (the break luminosity) are free parameters that are used to obtain the best fit to the available data. Local: &= -1.0 and M * B = -21 Virgo: &= and M * B = -21 ± 0.7 A power law with a high luminosity exponential cut-off i.e., this is NOT a universal luminosity function. It seems to depend upon environment. Press-Schecter Thus & sets the slope of the luminosity fn at the faint end L * or M * gives the characteristic luminosity above which the number of galaxies falls sharply and $ * sets the overall normalisation of the galaxy density. This formula was initially motivated by a simple model of galaxy formation (Press & Schecter 1974 ApJ ), but has proved to have a wider range of application than originally envisaged. Integration over previous eqn has limitation that it effectively predicts infinite number of small faint galaxies (alpha lies close to ~-1) However we know universe is finite (dark sky) However only place we can easily detect low lum galaxies. They exist in large numbers Galaxy Counts Schechter function Number of galaxies $(M) per 10Mpc cube between absolute magnitude M R and M R + 1; vertical bars are errors

10 Galaxy cluster LFs Different environment. Easier to obtain LF members lie in small region of sky. So photometry can be obtained efficiently, and all members at same distance. Reducing distance errors. Only problem is rich clusters are rare. So typically at large distances. Making it hard to detect fainter members. faint bright Jerjen & Tammann 1997 Faint end slope in cluster significantly steeper than in the field: encounters don t end in mergers as often as in field relative velocities are higher bright faint $* larger in clusters Galaxy luminosity function in the Virgo cluster (Sandage, Binggeli & Tammann 1985, AJ 90, 1759) Relative numbers of different types Largest fraction in either environment of all galaxies are dwarfs (de and Irr). Even though S and E the most prominent in terms of mass and luminosity. More E in Virgo... see: Thomas: ESO Astrophysics Symposia (1999) The total luminosity function in either environment is the sum of the individual luminosity functions of each Hubble type. Binggeli, Sandage, Tammann ARAA (1988) ARAA,26, 509 This is usually represented by the luminosity function '(M)dM. Defined to be the number of galaxies in a particular sample that have absolute magnitudes between M and M+dM. All dis and des

11 Masses of galaxies Spiral Galaxies: application of Gauss s theorem to Newton s law of gravity: Rotation curves allow mass determination. The constant rotational velocities in the outer regions - suggests that mass increases linearly with distance from the centre. In stark contrast to the light distribution, which decreases exponentially over the same distance. Meaning a rapidly increasing mass-to-light ration (M/L) and a hidden dark matter halo in spiral galaxies (Bosma 1981). Elliptical Galaxies: application of virial theorem, assuming isotropic stellar distribution This kind of analysis has led to the prediction of large dark matter haloes around elliptical galaxies (e.g., Côte et al. 2001, 2003), using globular clusters as tracers. The line-of-sight velocity dispersion remains remarkably constant out to the limits of observation. This has the same explanation as flat rotation curves in HI. To bind globular clusters with large velocity dispersions at large radii means that the mass within R must increase proportional to R. Rotation Curves of Galaxies Surface Brightness profiles sample the distribution of luminous matter in a galaxy. This does not necessarily tell us about the mass of the galaxy - about the presence and amount of DARK MATTER. The most direct way to do this is via the rotation curve of the HI. When rotation curves are compared with either luminosity or Hubble type a number of correlations are found:! for increasing L B rotation curves tend to rise more rapidly with distance from centre and peak at higher maximum velocity (V max ). Tully-Fisher! for equal L B spirals of earlier type have larger V max.! within a given Hubble type more luminous galaxies have larger V max.! for a given value of V max the rotation curves tend to rise slightly more rapidly with radius for earlier type galaxies. The fact that galaxies of different Hubble types, and therefore different bulge-to-disk luminosity ratios, exhibit rotation curves that are very similar in form if not in amplitude suggests that the shapes of the gravitational potential do not necessarily follow the distribution of luminous matter. V max is significantly lower in Irrs (50-70km/s). This suggests that this is the minimum rotation speed required for the development of a well ordered spiral pattern Bosma 1981 Internal dynamics of Ellipticals Source of galaxies shape? It might be thought that the internal dynamics of elliptical galaxies would be relatively simple - the surface brightness distributions appear to be ellipsoidal, with a range of flattenings, which it might be thought could be attributed to rotation. This can be tested by measuring the mean velocities and velocity dispersions of the stars through out the body of a galaxy. These measurements can be compared with the rotation and internal velocity dispersions expected if the flattening can be attributed to rotation. Solid line: amount of rotation necessary to account for observed ellipticity of galaxy relative to! of stars. Ellipticals rotate too slowly for centrifugal forces to be the causes of their observed flattening. This means that the assumptions of asymmetric spatial distribution and/or an isotropic velocity distribution of stars at all points within galaxy must be wrong. Converting luminosity to mass IMF (initial mass function) ((m, t), number of stars formed per unit volume at t=0 LF (luminosity function) currently observed number of stars observed per unit luminosity per unit volume PDMF (present day mass function) number of stars observed today per unit mass per unit volume. This needs to be corrected for the time evolution of the IMF up to the present day, low mass, long lived stars often approximated as a power law: ((m) dm = ( 0 m -& IMF TRIAXIAL SYSTEMS this means a system with 3 unequal axes and consequently anisotropic stellar velocity distributions PDMF high mass, short lived stars Star Formation History From Davies et al Kroupa, Tout & Gilmore 1993 MNRAS, 262, 545