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Home , Bullet Group, Density wave theory, Field galaxy, Hubble sequence

Ay21 Lectures 10 and 11: A Summary of the Key Ideas

Large Scale Velocity Field, Biasing, Galaxy Clusters

Galaxy Morphology, , Properties of the Spiral 1 Large-Scale Density Field Inevitably Generates a Peculiar Velocity Field

The PSCz survey local 3-D density field

A galaxy is accelerated towards the nearby large mass concentrations

Integrated over the Hubble time, this results in a peculiar velocity The pattern of peculiar velocities

should thus reflect the underlying mass density field 2 CMBR Dipole: The One Peculiar Velocity We Know Very Well

We are moving wrt. to the CMB at ~ 620 km/s towards b=27°, l=268° This gives us an idea of the probable magnitude of peculiar velocities in the local universe. Note that at the distance to (LSC), this corresponds to a ~ 50% error in Hubble velocity, and a

~ 10% error at the distance to . 3 How to Measure Peculiar Velocities? 1. Using distances and residuals from the Hubble flow:

Vtotal = VHubble + Vpec = H0 D + Vpec • So, if you know relative distances, e.g., from Tully-Fisher, or Dn-s relation, SBF, SNe, …you could derive peculiar velocities • A problem: distances are seldom known to better than ~10% (or even 20%), multiply that by VHubble to get the error of Vpec • Often done for clusters, to average out the errors • But there could be systematic errors - distance indicators may vary in different environments 2. Statistically from a survey • Model-dependent 4 Redshift Space vs. Real Space

“Fingers of God” Thin filaments Spatial Ÿ depth ˜ ˜ Real space distribution ˜ Ÿ The effect of cluster ˜ velocity dispersion The effect of infall Observed redshift Redshift space apparent distrib.

Position on the sky 5 Measuring Peculiar Velocity Field Using a Redshift Survey • Assume that galaxies are where their imply; this gives you a density field • You need a model on how the light traces the mass • Evaluate the accelerations for all galaxies, and their esimated peculiar velocities • Update the positions according to new Hubble velocities • Iterate until the convergence • You get a consistent density

and velocity field 6 The “Great Attractor” aka the -Centaurus

7 The Flow Continues? The Shapley Concentration of clusters at ~ 200 Mpc, beyond the Hydra-Centaurus may be responsible for at least some of the large-scale bulk flow

8 Peculiar Velocities: Summary • Measurements of peculiar velocities are very, very tricky – Use (relative) distances to galaxies + Hubble flow, to infer the peculiar velocities of individual galaxies. Systematic errors? – Use a redshift survey + numerical modeling to infer the mass density distribution and the consistent peculiar velocity field • Several general results: – We are falling towards Virgo with ~ 300 km/s, and will get there in about 10 - 15 Gyr – Our peculiar velocity dipole relative to CMB originates from within ~ 50 Mpc – The LSC is falling towards the Hydra-Centaurus Supercluster, with a speed of up to 500 km/s – The whole local ~ 100 Mpc volume may be falling towards a larger, more distant Shappley Concentration (of clusters) • The mass and the light seem to be distributed in the same way on

large scales (here and now) 9 Galaxy Biasing Suppose that the density fluctuations in mass and in light are not the same, but

(Dr/r)light = b (Dr/r)mass Or: 2 x(r)light = b x(r)mass Here b is the bias factor. If b = 1, light traces mass exactly (this is indeed the case at z ~ 0, at scales larger than the individual galaxy halos). If b > 1, light is a biased tracer of mass. One possible mechanism for this is if the galaxies form at the densest spots, i.e., the highest peaks of the density field. Then, density fluctuations containing galaxies would not be typical, but rather a biased representation of the underlying mass density field; if 1-s fluctuations are typical, 5-s ones certainly are not. 10 High Density Peaks as Biased Tracers Take a cut through a density field. Smaller fluctuations ride atop of the larger density waves, which lift them up in bunches; thus the highest peaks (densest fluctuations) are a priori clustered more strongly than the average ones:

Proto- cluster

Proto-

Thus, if the first galaxies form in the densest spots, they will be strongly clustered, but these will be very special regions. 11 An Example From a Numerical Simulation

All particles 1-s peaks 2-s peaks 3-s peaks

Gas/

Dark Matter

(From an N-body simulation by R. Carlberg) 12 Evolution of Clustering • Generally, density contrast grows in time, as fluctuations collapse under their own • Thus, one generically expects that clustering was weaker in the past (at higher redshifts), and for fainter galaxy samples • Deep redshift surveys indicate that the strength of the clustering decreases at higher redshifts, at least out to z ~ 1, as expected:

(Coil et al., DEEP survey team) 13 Evolution of Clustering • But at higher redshifts (and fainter/deeper galaxy samples), the trend reverses: stronger clustering at higher redshifts = earlier times! What is going on?

(Hubble Deep Field data)14 Biasing and Clustering Evolution Strength of clustering Higher density (= higher-s) fluctuations evolve faster 1-s flucs. 3-s flucs. 5-s flucs.

redshift At progressively higher redshifts, we see higher density fluctuations, which are intrinsically clustered more strongly … Thus the net strength of clustering seems to increase at higher z’s 15 Evolution of Clustering and Biasing • The strength of clustering (of mass) grows in time, as the gravitational infall and hierarchical assembly continue – However, the rate of growth and the strength of clustering at any given time depend on the mass and nature of objects studied – This is generally expressed as the evolution of the 2-point correlation function, x(r,z) = x(r,0) (1+z) -(3+e) – Clustering/LSS is observed out to the highest redshifts (z ~ 4 - 6) and it is surprisingly strong • What we really observe is light, which is not necessarily distributed in the same way as mass; this is quantified as bias: 2 (Dr/r)light = b (Dr/r)mass, x(r)light = b x(r)mass – Bias is a function of time and mass/size scale – Galaxies (especially at high redshifts) are biased tracers of LSS, as the first objects form at the highest peaks of the density field

– Today, b ~ 1 at scales > galaxies 16 Clusters of Galaxies: • Clusters are perhaps the most striking elements of the LSS • Typically a few Mpc across, contain ~ 100 - 1000 luminous 14 15 galaxies and many more dwarfs, masses ~ 10 - 10 M • Gravitationally bound, but may not be fully virialized • Filled with hot X-ray gas, mass of the gas may exceed the mass of stars in cluster galaxies • is the dominant mass component (~ 80 - 85%) • Only ~ 10 - 20% of galaxies live in clusters, but it is hard to draw the line between groups and clusters, and at least ~50% of all galaxies are in clusters or groups • Clusters have higher densities than groups, contain a majority of E’s and S0’s while groups are dominated by spirals • Interesting galaxy evolution processes happen in clusters 17 The Virgo Cluster The Coma Cluster • Irregular, relatively poor cluster • Nearest rich cluster, with • Distance ~ 16 Mpc, closest to us >10,000 galaxies • Diameter ~ 10° on the sky, 3 Mpc • Distance ~ 90 Mpc • ~ 2000 galaxies, mostly dwarfs • Diameter ~ 4-5° on the sky, 6-8 Mpc

X-ray / visible overlay 18 Hot X-ray Gas in Clusters • Virial equilibrium temperature T ~ 107 – 108 K, so emission is from free-free emission • Many distant clusters are now being discovered via x-ray surveys • Temperatures are not uniform, we see patches of “hot spots” which are not obviously associated with galaxies. May have been heated as smaller galaxies (or clumps of galaxies) fell into the cluster • In densest regions, gas may cool and sink toward the cluster center as a “cooling flow” • Unlikely that all of it has escaped from galaxies, some must be primordial, from the cluster formation process. It is heated via shocks as the gas falls into the cluster potential • But some metals, ~ 1/3 Solar, must be from stars in galaxies • X-ray correlates with cluster classification, regular clusters have high x-ray luminosity, irregular clusters have low x- ray luminosity 19 Virial Masses of Clusters: Virial Theorem for a test particle (a galaxy, or a proton), moving in a cluster potential well: 2 Ek = Ep / 2 ! mg s / 2 = G mg Mcl / (2 Rcl) where s is the velocity dispersion

2 Thus the cluster mass is: Mcl = s Rcl / G Typical values for clusters: s ~ 500 - 1500 km/s Rcl ~ 3 - 5 Mpc 14 15 Thus, typical cluster masses are Mcl ~ 10 - 10 M The typical cluster (~ 100 - 1000 galaxies) 12 are Lcl ~ 10 L, and thus (M/L) ~ 200 - 500 in solar units ! Lots of dark matter! 20 Dark Matter and X-Ray Gas in Cluster Mergers: The “” (1E 0657-56) The dark matter clouds largely pass through each other, whereas the gas clouds collide and get shocked, and lag behind

Blue: dark matter, as inferred from weak gravitational lensing

Pink: X-ray gas

(Bradac et al.) 21 Clusters as Cosmological Probes • Given the number density of nearby clusters, we can calculate how many distant clusters we expect to see • In a high density universe, clusters are just forming now, and we don’t expect to find any distant ones • In a low density universe, clusters began forming long ago, and we expect to find many distant ones • Evolution of cluster abundances: – Structures grow more slowly in a low density universe, so we

expect to see less evolution when we probe to large distances22 Hydrogen Gas Deficiency • As gas-rich galaxies (i.e., spirals) fall into clusters, their cold ISM is ram-pressure stripped by the cluster X-ray gas • Evidence for stripping of gas in cluster spirals has been found from HI measurements • Most deficient spirals are found in cluster cores, where the X-ray gas is densest All H I gone • HI deficiency correlates with X- ray luminosity (which correlates with cluster richness) • It is the outer disks of the spirals that are missing

• Thus, evolution of disk galaxies H I still there can be greatly affected by their large-scale environment Little X-ray gas Lots of X-ray gas 23 Intracluster Light • Zwicky in 1951 first noted “an extended mass of luminous intergalactic matter of very low surface brightness” in Coma cluster • Confirmed in 1998 by Gregg & West, features are extremely low surface brightness >27 mag per arcsec2 in R band • Also discoveries of intracluster red giant stars and intracluster planetary nebulae in Virgo & , up to ~ 10-30% of the total cluster light • Probably caused by galaxy- galaxy or galaxy-cluster potential tidal interactions, which do not result in outright mergers – This is called “galaxy harassment” – Another environment dependent process affecting

galaxy evolution 24 Clusters of Galaxies: Summary • Clusters are the largest bound (sometimes/partly virialized) elements of the LSS 2 3 14 15 – A few Mpc across, contain ~ 10 - 10 galaxies, Mcl ~ 10 - 10 M – Contain dark matter (~80%), hot X-ray gas (~10%), galaxies (~10%) – This maps into discovery methods for clusters: galaxy overdensities, X-ray sources (via emission of SZ effect), weak lensing, etc. • Clusters are still forming, via infall and merging – Studied using numerical simulations, with galaxies, gas, and DM • Galaxy populations and evolution in clusters differ from the general field – While only ~ 10 - 20% of galaxies are in clusters today, > 50% of all galaxies are in clusters or groups – Clusters have higher fractions of E’s and S0’s relative to spirals – Interesting galaxy evolution processes happen in clusters 25 Galaxies • The basic constituents of the universe at large scales – Distinct from the LSS as being too dense by a factor of ~103, indicative of an “extra collapse”, and a dissipative formation • Have a broad range of physical properties, which presumably reflects their evolutionary and formative histories, and gives rise to various morphological classification schemes (e.g., the Hubble type) • Understanding of galaxy formation and evolution is one of the main goals of modern cosmology • There are ~ 1011 galaxies within the observable universe 8 12 • Typical total masses ~ 10 - 10 M 7 11 • Typically contain ~ 10 - 10 stars 2 6 Elliptical Galaxies • About 20% of field galaxies are E’s, but most E’s are in clusters • There are subtypes: – E’s (normal ellipticals) – cD’s (massive bright ellipticals at the centers of galaxy clusters) – dE’s (dwarf ellipticals) Not really ellipticals, a – dSph’s (dwarf spheroidals) } different class of objects • Smooth and almost featureless: no spiral arms or dust lanes. Generally lacking in cool gas, and hence few young blue stars • Classified by the apparent ellipticity: b ε =1− a a b Elliptical galaxies are b n =1− denoted En, where: a 10 A round elliptical is E0, the most elongated ellipticals are E7 27 € € Lenticular (S0) Galaxies • Transition class between ellipticals and spirals are the S0 galaxies, also called lenticulars • S0 galaxies have a rotating disk in addition to a central elliptical bulge, but the disk lacks spiral arms or prominent dust lanes, i.e., no active formation • Lenticulars can also have a central bar, in which case they are labeled SB0 • May originate from spirals that have exhausted their gas, or that were stripped Spiral Galaxies Named for their bright spiral arms, which are prominent due either to bright O and B stars (evidence for recent ), and dust lanes. Define two parallel sequences of spirals: Sa Sb Sc Sd

Central bulge becomes less important Disk becomes more important Spiral arms become more open and ragged SBa SBb SBc SBd As above, except that these galaxies also have a central, linear bar, while the Sa, Sb… are unbarred 29 Dwarf Galaxies Starforming gas rich dwarf 6 10 • Low-luminosity: 10 – 10 L¤, low-mass: 7 10 10 – 10 M¤, small in size, ~ few kpc • Often low surface brightness, so they are hard to find! • More than one family of objects: – Gas-poor, passive (dE and dSph) – Gas-rich, star forming • Why are dwarf galaxies important? – Majority of galaxies are dwarfs! – Dwarf galaxies may be remnants of Sagittarius dSph galaxy formation process: “proto-dwarf” gas clouds came together to form larger galaxies (hierarchical formation) – Dwarf galaxies are currently being cannibalized by larger galaxies – Dwarf galaxies are relatively simple systems, not merger products: in some sense, “pristine” galaxies 30 Hubble’s Classification Scheme

Ellipticals classified by the apparent ellipticity

Spirals classified by the prominence of the spiral arms, and the presence of bars

Hubble thought (incorrectly) this was an evolutionary sequence, so ellipticals are called “early-type” and spirals “late-type” galaxies 3 1 Problems With Traditional Galaxy Classification Appearance of galaxies is strongly dependent on which wavelength the observations are made in.

e.g., the nearby galaxy M81:

X-ray UV Visible Near-IR Far-IR Note: large change in appearance between the UV and the near infrared images. Galaxies look “clumpier” in the UV, and increasingly smooth as we go to the visible and longer wavelengths. 32 Problems With Traditional Galaxy Classification Subjective - especially for spiral galaxies However, there are automated, objective schemes to classify galaxies, using measured image parameters.

Superficial - based on appearance, not physical properties Galaxy types or families can be defined in a parameter space of various measured/physical quantities. Different galaxy families follow different correlations.

Incomplete - misses the major dichotomy of dwarfs and giants (not separated in the traditional Hubble sequence) Dwarfs also exist in gas rich / gas poor, star forming or not, and perhaps other varieties 33 The Meaning of Galaxy Classification • Galaxy morphologies and other properties reflect different formative and evolutionary histories • Much can be explained by considering galaxies as composites made of two dominant visible components: 1. Old, pressure supported bulges, where most of the star formation occurred early on 2. Young(er), rotationally supported disks, where star formation happened gradually and is still going on • Note that we do not involve in this the dominant mass component - the dark matter … and that spiral arms may be mainly ornamental … • Nevertheless, there are some important and meaningful trends along the Hubble sequence 34 Galaxy Properties and the Hubble Sequence Hubble sequence turned out to be surprisingly robust: many, but not all, physical properties of galaxies correlate with the classification morphology: E S0 Sa Sb Sc Sdm/Irr

Pressure support ! Rotational support Passive ! Actively star forming Red colors ! Blue colors Hot gas ! Cold gas and dust Old ! Still forming High luminosity density ! Low lum. dens. … etc. But, for example, masses, luminosities, sizes, etc., do not correlate well with the Hubble type: at every type there is a large spread in these fundamental properties. 35 Interpreting the Trends Along the Hubble Sequence • Probably the best interpretation of many of these is a trend in star formation histories: – Ellipticals and early type spirals formed most of their stars early on (used up their gas, have older/redder stars) – Late type spirals have substantial on-going star-formation, didn’t form as many stars early-on (and thus lots of gas left)

– Spirals are forming stars at a few M¤ per year, and we 9 know that there is ~ a few x 10 M¤ of HI mass in a typical spiral • How long can spirals keep forming stars?? It seems that some gas infall/resupply is needed 36 Star Formation History in Galaxies

37 Stellar Populations • A key concept in our understanding of galaxies • In 1944, Walter Baade used the 100-inch Mt. Wilson telescope to resolve the stars in several nearby galaxies: M31, its companions M32 and NGC 205, as well as the elliptical galaxies NGC 147 and NGC 145 • Realized the stellar populations of spiral and elliptical galaxies were distinct: – Population I: objects closely associated with spiral arms – luminous, young hot stars (O and B), Cepheid variables, dust lanes, HII regions, open clusters, metal-rich – Population II: objects found in spheroidal components of galaxies (bulge of spiral galaxies, ellipticals) – older, redder stars (red giants), metal-poor 38 Stellar Populations and Dynamical Subsystems in Galaxies • The picture today is more complex: it is useful to thing about generalized stellar populations as subsystems within galaxies, characterized by the: – Location and morphology, density distribution – Dynamics (rotation, random motions, their distribution) – Star formation rate and mean age – The presence and nature of its etc., etc. • For example, in the , we can distinguish: – Young – Old – Metal-rich bulge (and bar?) – Metal-poor stellar halo Formation of Galaxy Spheroids and Dynamics of Stellar Populations

Stars “remember” the dynamics of their orbits at the time of formation, since dynamics of stellar systems is dissipationless. If stars form in dwarf protogalactic fragments which then merge, this will result in a pressure-supported system, i.e., a spheroid (bulge or halo, or an ). Their will reflect the abundances in their parent systems. 40 Formation of Galaxy Disks and Dynamics of Stellar Populations

Q Q

If protogalactic clouds merge dissipatively in a potential well of a dark halo, they will settle in a thin, rotating disk = the minimum energy configuration for a given angular momentum. If gas settles into a (dynamically cold) disk before stars form, then stars formed in that disk will inherit the motions of the gas (mainly an ordered rotation). 41 Chemical Self-Enrichment in Young Stellar Systems

In a low-mass system, supernova shocks and star winds from In a massive system, supernova massive young stars expell the ejecta are retained, and reused enriched gas and may supress any for subsequent generations of subsequent star formation. The stars, which achieve ever system retains its initial (low) higher metallicities. metallicity. 42 Quantifying Properties of Galaxies For galaxies of different types, we would like to quantify: • The distribution of light - need photometric measurements • The distribution of mass - need kinematical measurements • Relative distributions and interplay of various components, e.g., stars, gas, dark matter - need multiwavelength measurements, as different components tend to emit most energy in different wavebands, e.g., stars ® visible/near-IR, cold gas ® radio, dust ® far-IR, hot gas ® x-rays, etc. • Chemical composition, star formation rates - need spectroscopy All these measurements can then be analyzed using: • Dynamical models • Stellar population synthesis models Note: we tend to measure • Galaxy evolution models different observables for different galaxy types! 43 Global Properties of Spiral Galaxies Spirals are complex systems, generally more complex and diverse than ellipticals: • Wide range in morphological appearance • Fine scale details – bulge/disk ratios, structure of spiral arms, resolution into knots, HII regions, etc. • Wide range in stellar populations – old, intermediate, young, and currently forming • Wide range in stellar dynamics: – “cold” rotationally supported disk stars – “hot” mainly dispersion supported bulge & halo stars • Significant amounts of cold interstellar medium (ISM) Spirals tend to avoid high-density regions (e.g., clusters) as they are dynamically fragile, and can be merged and turned into E’s 44 Spiral Galaxies: Basic Components • Disks: generally metal rich stars and ISM, nearly circular orbits with little random motion, spiral patterns – Thin disks: younger, star forming, dynamically very cold – Thick disks: older, passive, slower rotation and more random motions • Bulge: metal poor to super-metal-rich stars, high stellar densities, mostly random motion – similar to ellipticals • Bar: present in ~ 50 % of disk galaxies, mostly older stars, some random motions and a ~ solid body rotation? • Nucleus: central (<10pc) region of very high mass density, massive black hole or starburst or nuclear star cluster • Stellar halo: very low density (few % of the total light), metal poor stars, globular clusters, low density hot gas, little or no rotation • Dark halo: dominates mass (and gravitational potential) outside a few kpc, probably triaxial ellipsoids, radial profile ~ singular isothermal sphere, DM nature unknown 45 Photometric Properties of Galaxies Empirically, the surface brightness declines with distance from the center of the galaxy in a characteristic way for spiral and elliptical galaxies For spiral galaxies, need first to correct for: • Inclination of the disk • Dust obscuration • Average over spiral arms to obtain a mean profile Corrected disk surface brightness drops off as: -R h I(R) = I(0) e R where I(0) is the central surface brightness of the disk, with a broad range of values, but typically ~ 21 - 22 mag/arcsec2, and hR is a characteristic scale length, with typical values: 1 kpc < h <10 kpc R 46 Bulge-Disk Decomposition In practice, surface brightness at the center of many spiral galaxies is dominated by stars in a central bulge. Central surface brightness of disk must be estimated by extrapolating inward from larger radii

Component profiles (µ is the logarithmic surface brightness in mags/arcsec2):

log surface brightness

radius 47 Spiral Galaxies: Gas Content • Gas in spirals: – Cool atomic HI gas

– Molecular hydrogen H2, CO, many other molecules – Need cold gas to form stars! Star formation associated with dense ISM – Can observe ionized hydrogen via optical emission-lines (Ha) – Observe HI via radio emission – 21 cm line due to hyperfine structure – a hydrogen atom that collides with another particle can undergo a spin-flip transition • Spirals show HI disks (amount of HI depends on Hubble type) • HI gas is optically thin, 21 cm line suffers little absorption, so we can measure gas mass directly from line intensity • HI is much more extended than optical light • Can use radial motion of 21 cm line to measure rotation in spiral galaxies 48 A Basic Tool: Spin-Flip (21 cm) Line of H I

In emission generally originates from warm (T ~ 100 - 6000 K) ISM, which accounts for ~ 30 - 65% of the total ISM volume in the Galactic disk. In absorption, it probes a cooler ISM (can be also self-absorbed).

Typical line profile !

A major advantage: it is not affected by the dust absorption! 49 Visible Light and Molecular Gas (CO)

50 Spiral Arms Defining feature of spiral galaxies - what causes them?

Observational clues:

Seen in disks that contain gas, but not in gas poor S0 galaxy disks.

Defined mainly by blue light from hot massive stars, thus lifetime is << galactic rotation period

When the sense of the galactic rotation is known, the spiral arms almost always trail the disk rotation 51 • Spiral arm patterns muct be persistent. Density wave theory provides an explanation: the arms are desity waves propagating in differentially rotating disks • Spiral arm pattern is amplified by resonances between the epicyclic frequencies of the stars (deviations from circular orbits) and the angular frequency of the spiral pattern – Spiral waves can only grow between the inner and outer Linblad resonances (Wp = W - k/m ; Wp = W + k/m ) where k is the epicyclic frequency and m is an integer (the # of spiral arms) – Stars outside this region find that the periodic pull of the spiral is faster than their epicyclic frequency, they don’t respond to the spiral and the wave dies out – Resonance can explain why 2 arm spirals are more prominent • We observe resonance patterns in spirals 52 Spiral Density Waves • The orbits in spiral galaxies are not quite circles – they are ellipses. These ellipses are slightly tilted with respect to each other.

• Thus there are regions of slightly higher density than their surroundings. The higher density means higher gravity.

• Objects (such as a gas cloud) will be attracted to these regions and will drift towards them. 53 Spiral Density Waves • When the gas cloud collides with other gas clouds, stars will be formed. (This is where most of the galaxy’s star formation takes place.)

• Many of the stars will be faint, red main sequence stars, but some will be bright blue OB stars. These stars will continue to drift through the region.

• The OB stars don’t go far before they explode. The brightest (and bluest) of a galaxy’s stars will never be far from the spiral arm where they were born. 54 The Density Wave Theory

M51 Density Wave Theory Summary • Spiral arms are waves of compression that move around the galaxy and trigger star formation • Star formation will occur where the gas clouds are compressed • Stars pass through the spiral arms unaffected • This theory is successful in explaining the properties of spiral galaxies • Two outstanding problems with it: 1. What stimulates the formation of the spiral pattern? Tidal interactions? 2. What accounts for the branches and spurs in the spiral arms? 56