Galaxies Astro 430/530 Spring 2018 Prof. Jeff Kenney

Transcription

1 Galaxies Astro 430/530 Spring 2018 Prof. Jeff Kenney CLASS 3 January 22, 2018 Non- parametric QuanItaIve Morphology & VerIcal distribuions of starlight 1

2 FuncIons fit to Galaxy Radial light profiles ExponenIal disk: I(r) = I(0) exp (- r/r d ) DeVaucouleurs r 1/4 bulge law: I(r) = I(r eff ) exp {- 7.67[( r/r eff ) 1/4-1]} Sersic law: I(r) = I(r eff ) exp {- b n [( r/r eff ) 1/n - 1]} n = Sersic index n = 1-4 typically If n=1 exponenial (all disk) [b n chosen to make r eff the effecive radius] If n=4 devaucouleurs r/4 law (all bulge) If n<2 small bulge- disk raio If n>2 large bulge- disk raio Advantage of Sersic law: can describe en5re profile shape with just one number n

3 Color- magnitude diagram of 183,000 nearby SDSS galaxies g- r All environments Red sequence galaxies Have liile or no star formaion Galaxies end up here Blanton+03 Green valley galaxies TransiIoning from blue cloud to red sequence What processes drive their evoluion? Blue cloud galaxies AcIvely forming stars Galaxies start here M I Color- magnitude diagram: disinct red sequence and blue cloud Most luminous (and massive) galaxies are red For red galaxies, color- mag relaion (slope of line) due to more metals in more massive gals Galaxy evoluion galaxies start blue (star- forming) eventually become red (less star- forming)

4 DistribuIon of opical color, central surface brightness, Sersic index, & absolute magnitude, for ~10 6 SDSS galaxies Red sequence Blanton etal 2005, ApJ, 629, 143 Blue cloud Red sequence Blue cloud Color- magnitude diagram: disinct red sequence and blue cloud Color- sersic index diagram: most red galaxies have n>2 i.e., bulge- dominated most blue galaxies have n<2 i.e., disk- dominated Color- surface brightness diagram: most red galaxies have higher central SB i.e., bulge- dominated most blue galaxies have low central SB i.e., disk- dominated Most luminous (and massive) galaxies are red 4

5 g- r Color- magnitude (~ssfr- M * ) vs. environment ALL low ρ high ρ v.high ρ M I SDSS Hogg+04 Dense environments like clusters have: Fewer blue cloud galaxies Less disinct green valley Stop star formaion more effecively

6 Color- magnitude (~ssfr- M * ) vs. environment ALL low ρ high ρ v.high ρ g- r SDSS Hogg+04 in low density environments, very few red disky (n<2) galaxies in high density environments, many red disky (n<2) galaxies

7 Not all galaxies obey Sersic law! (or have radial light distribuions that are well fit by any simple equaion with a small number of parameters) à Measure properies of light distribuion which doesn t require fivng the data to some equaion: Non- parametric analysis Merger remnant NGC 6240 InteracIng dwarf NGC

8 Galaxy light distribuions: parametric vs. non- parametric descrip5ons Parametric fit a funcion to the distribuion (e.g. Sersic index), the distribuion can be characterized by a parameter of that funcion Non- parametric don t fit a funcion! Do something else. 8

9 Non- parametric descripions of light distribuions: one example: CAS: ConcentraIon, Asymmetry, Smoothness Conselice

10 CAS: ConcentraIon What fracion of the light is in the central regions? Concentra3on index (C): central concentraion of light; raio of radii that contain 20% and 80% of the total light good measure of B/D but also shows galaxies concentrated without tradiional bulge (e.g. concentrated due to circumnuclear disk (pseudobulge) or SF due to interacion). 10

11 bulge to disk raio (B/D) L disk = I disk (r) da L bulge = I bulge (r) da B/D = L bulge /L disk 11

12 CAS: Asymmetry What frac5on of the light is in non- symmetric components? Original image image rotated 180 o difference Asymmetry index (A): measure of how asymmetric a galaxy is azer rotaing the image 180 o along the line of nodes. NOTES: 2- arm spiral structure is symmetric in this index but 1- arm or 3- arm spiral structure is asymmetric. Other kinds of asymmetry indices may be beier for showing paricular types of asymmetries. 12

13 CAS: Smoothness (Clumpiness) What fracion of the light is in small scale features? Original image blurred image difference Smoothness clumpiness index (S): Subtract smoothed version of Image from unsmoothed version. To show small scale substructure like spiral arms and star forming complexes. (aka unsharp masking ). 13

14 Non- parametric photometric quaniies: CAS+Gini Lotz+2004 ConcentraIon Asymmetry Smoothness Gini coeff Red/Early type galaxies are highly concentrated, symmetric, have liile substructure Blue/late type galaxies are less concentrated, asymmetric, have lots of substructure Bar+bulge Light distributed Odd- m Spiral arms Structure due to & Star formaion Star forming regions among more pixels (true for disks); corr w/b/d but works even 14 if light conc are not in center

15 Gini coefficient: income distribuion among people country where all people have same income (gini = 0) country where 1 person has all the income (gini = 1) actual country (gini = 0-1) The Gini coefficient (also known as the Gini index or Gini ra;o) is a measure of staisical dispersion intended to represent the income distribuion of a naion's residents, and is the most commonly used measure of inequality. The Gini coefficient measures the inequality among values of a frequency distribuion. A Gini coefficient of zero expresses perfect equality, where all values are the same. A Gini coefficient of one (or 100%) expresses maximal inequality among values. 15

16 Gini coefficient: light distribuion among pixels percentage of total light galaxy where all pixels have same light (gini = 0) percentage of pixels galaxy where 1 pixel has all the light (gini = 1) actual galaxy (gini = 0-1) 16

17 CAS parameters for nearby galaxies Conselice

18 CAS parameters for nearby galaxies Conselice

19 Non- parametric photometric quaniies: concentraion & asymmetry Early types Late types Concentra;on vs. Mean Surface Brightness: Galaxies with high surface brightnesses are more concentrated, on average (trend is obvious but note lots of scaier) Late types Early types high SB low SB blue red Asymmetry vs. Color: Blue galaxies are more asymmetric, on average (more star formaion) 19 Yagi+2006 MNRAS 368, 211

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21 Disks of highly- inclined spiral galaxies Disk of S0 galaxy What do these images tell us about stellar disks? 21

22 good view of edge- on stellar disk in S0 galaxy NGC 4452 HST stellar disk is relaively thin disk thickness is small compared to disk radius 22

23 NGC 4565 Sb galaxy dust (& cold gas) are disk component not bulge component 23

24 NGC 891, a nearby galaxy similar to Milky Way, viewed edge- on compare verical distribuions of stars and dust stellar disk thicker than dust (& gas) disk 24

25 Stellar disks oden but not always associated with dust/gas disks. Stellar bulges not associated with dust/gas. Stellar disks generally thicker than gas disks. Gas disks in some galaxies appear disturbed. 25

26 Why does disk form? 26

27 disk must form from gas, since gas is collisional disk cannot form from stars, since stars are collisionless 27

28 How does stellar disk form? Gas, which is collisional and dissipates energy through collisions, seiles to a rotaing thin gas disk. Stars form in giant molecular clouds (GMCs) of dense gas, which are embedded within a thin disk of gas. The youngest stars are therefore in a disk with the same thickness as the layer of star- forming dense gas. 28

29 Why rotaing gas cloud forms thin rotaing disk Gas blob (nearly) all random mo5ons result in inelasic collisions, which dissipate energy KE random à other forms of energy Gas blob Gas blob Gas blob ordered motions (net rotation) do not result in collisions (orbits are parallel & non-intersecting) KE ordered à remains ~constant Only motions supporting cloud in directions other than rotation direction are random so if particles are collisional the cloud collapses in all directions 29 other than rotation direction

30 Why rotaing gas cloud forms thin rotaing disk Gas blob (nearly) all random mo5ons result in inelasic collisions, which dissipate energy KE random à other forms of energy 30

31 Why rotaing gas cloud forms thin rotaing disk Gas blob (nearly) all random mo5ons result in inelasic collisions, which dissipate energy KE random à other forms of energy this component results in collision so mo;on dissipates this component doesn t experience collision so mo;on remains 31

32 examine nearest disk: Milky Way Galaxy The structure of stellar disks reveals important informaion about their formaion and evoluion Both internal and external processes can make them thicker over Ime 32

33 Density of stars (of paricular type S) in Milky Way disk Use cylindrical coordinates (R,φ,z) and integrate over azimuthal angle φ h R = (radial) scale length of disk h z = (vertical) scale height of disk n(0,0,s) = density of stars of type S at galaxy center 33

34 Density of stars (of paricular type S) in Milky Way disk Radial distribution: Exponential is excellent fit to light profiles of many disk galaxies. Some galaxies deviate from this. Scale length differences for different types of stars in some but not all disks (generally not as important as type differences in vertical distributions). Hard to study in Milky Way 34

35 Density of stars (of paricular type S) in Milky Way disk Vertical distribution: Nothing special about exponential. It gives OK rough fit for some types of stars. Very different scale heights for stars of different type or age. (age is main parameter, but harder to determine than type.) Possible to study in detail in Milky Way! 35

36 Density of stars (of paricular type S) in Milky Way disk Use cylindrical coordinates (R,φ,z) and integrate over azimuthal angle φ h R = (radial) scale length of disk h z = (vertical) scale height of disk n(0,0,s) = density of stars of type S at galaxy center Radial distribution: Exponential is excellent fit to light profiles of many disk galaxies. Some galaxies have deviations from this. Scale length differences for different types of stars in some but not all disks (generally not as important as type differences in vertical distributions) Hard to study in Milky Way Vertical distribution: Nothing special about exponential, gives OK rough fit for some types of stars Very different scale heights for stars of different type or age (age is main parameter, but age is harder to determine than type) 36 Possible to study in detail in Milky Way

37 we can examine verical structure of Milky Way disk (if we can measure distances to the stars) Study stars in direcions ~perpendicular to disk plane 37

38 VerIcal distribuion of stars near Sun in Milky Way 38

39 VerIcal distribuion of stars near Sun in Milky Way sum thin disk thick disk halo Note that the A stars have a very small scale height (MS A stars have age<100 Myr) Why do different types of stars have different verical distribuions? G,K stars have large age range but many are old > 3 Gyr 39

40 What relaion do we expect between verical scale height and verical velocity dispersion? Poisson s equa;on (relaion between gravitaional force & density) Louisville s equa;on (one of Jeans equaions, which comes from Collisionless Boltzmann EquaIon)(equilibrium state for system of many idenical collisionless paricles) SoluIon to this pair of equaions is: ρ(z) = ρ o sech 2 (z/z e ) <v z2 > = 2πG ρ o z e 2 For a given mass surface density s, the verical velocity determines the scale height (In simple terms: for a given mass surface density, if the velocity dispersion is higher, then the stars will reach a greater height above the plane 40

41 What relaion do we expect between verical scale height z e and verical velocity dispersion <v 2 > 1/2? σ z2 =<v z2 > = A π G Σ z e a. A = 2 if ρ = ρ o sech 2 (z/z e ) b. A = 1.7 if ρ = ρ o sech (z/z e ) c. A = 1.5 if ρ = ρ o exp (z/z e ) sech 2 has ~constant density near disk midplane ( core ) exp is peaky ( cuspy ) near disk midplane sech 2 theoreical expectaion from (too) simple model exp or sech closer to observaions A good way to es5mate the mass surface density Σ of disks: measure <v 2 > and z e 41

42 τ<100 Myr (O,B stars) ~100 pc

43 Structure & kinemaics of Milky Way disk: key points from table 2.1 Increase in velocity dispersion corresponds to increase in scale height and decrease in mean rotaional velocity Velocity dispersion of disk stars increases with age of stars There are different velocity dispersions in different direcions (R, φ, z) why? Halo is not part of disk, and its origin is physically disinct. But there are halo stars located in the solar neighborhood and within the disk. 43

44 VerIcal velocity vs. age Nearby main sequence F and G stars O = low metallicity stars (Z < 0.25 Z sun ) 44

45 Evolu;on of stellar disks GravitaIonal interacions between stars and either GMCs or spiral arms transfer energy to the stars, heaing them up dynamically, thereby increasing their verical moions and their average height above the disk midplane internal, con5nuous process origin of gradual trend of increasing velocity dispersion with age of stars 45

46 Evolu;on of stellar disks GravitaIonal interacions between stars and either GMCs or spiral arms transfer energy to the stars, heaing them up dynamically, thereby increasing their verical moions and their average height above the disk midplane internal, con5nuous process origin of gradual trend of increasing velocity dispersion with age of stars Mergers of (small) galaxies with the Milky Way galaxy gravitaionally disturb the stars in the disk, also heaing them up dynamically external, discrete random events origin of thick disk 46

47 Not all spiral galaxies have thick disks! Superthin Galaxy UGC7321 Maihews+ 1999; Maihews 2000 WIYN R- band LSB galaxy, i=88 deg, V max ~100 km/s VerIcal scale height in H- band Single component fit at r=0 (minor axis) exponenial fit 2.9 =140 pc (~smallest value known) h r /h z = 14 (~largest value known) Possible 2- component fit to disk at r= +/- 60 Sech 2 fits, scale heights 3.8,8.7 It is possible that even this superthin disk has complex structure and internal dynamical heaing but probably no merger so no thick disk & not much of a stellar halo Simplest & least evolved stellar disks? 47

48 Thick disks in spiral galaxies 34 late- type, edge- on, undisturbed, disk galaxies spanning a wide range of mass Milky Way RaIo of verical scale heights for thick and thin components has range ~1.5-5 (factor of ~3) Yoachim & Dalcanton

49 Thick disks in spiral galaxies 34 late- type, edge- on, undisturbed, disk galaxies spanning a wide range of mass Disk thickness appears to scale with circular velocity (~galaxy mass) for both thin and thick disk components not well understood thin thick superthin UGC7321 Galaxy mass - > Galaxy mass - > Yoachim & Dalcanton

50 Which galaxy is making a thick disk? 50

51 Which galaxy is making a thick disk? Minor mergers can make thick stellar disks in 2 ways: 1. Pre- exising (old) stellar disk is dynamically heated by gravitaional interacion 2. Gas from either galaxy iniially has disturbed configuraion, stars form in it 51

52 Stellar disks oden but not always associated with dust/gas disks. you need to first form a gas disk to make a stellar disk, but azer that the gas in the disk can go away. Stellar bulges not associated with dust/gas. bulges formed by process(es) not involving gas dissipaion Stellar disks generally thicker than gas disks. stars form in thin disk but later things can happen to make stellar disk thicker (OR some stars formed at earlier Imes in thicker gas disk) Gas disks in some galaxies appear disturbed. mergers can make irregular gas disks which eventually dissipate energy and become thin 52

53 VerIcal disk structure - - caveats SIll debated whether Milky Way (& other galaxies) has disinct thick disk component, or gradual increase in z e with age (e.g. Bovy etal 2012) Dust exincion near disk midplane makes it difficult to accurately measure verical distribuions of stars in most galaxies (NIR beier than opical) Stellar populaion differences with z make it difficult to accurately interpret verical distribuions of stars would like to know distribuion of each stellar populaion separately AND distribuion of stellar mass (these are hard but worthwhile!) 53