Gaseous Phases in Galaxies

Cetraro, Giugno 3 - 7, 20021

Gaseous Phases in Galaxies Gaseous Phases in Galaxies

1. Introduction

4. Dust

2. Atomic gas

3. Molecular gas

5. Hot gas

6. Heating and cooling

8. Gas mass and b

7. Galactic Winds

Uli Klein

Univ. Bonn


Guess what and where!


... in the LMC

Gas phases ...


1. Introduction

BBNS: 76% H, 24%He, 10-4 3He, 10-4 2H, 10-9 - 10-10 Li, Be

today : 66% H, 32% He, 2% ‘metals‘

Interstellar cycle:


Phase n [cm-3] T [K] fV [%] fm [%] h [pc] Tracer

molecular 102 ··· 105 10 ··· 50 < 1 ~20 ~70 CO

cold neutral 40 ··· 80 50 ··· 200 2 ··· 4 ~40 ~140 HI (absorpt.)

warm neutral 0.1 ··· 0.6 5500 ··· 8500 ~30 ~30 ~400 HI (emission)

warm ionized ~0.2 ~8000 ~20 ~10 ~900 Hradio

hot ionized 10-3 ··· 10-2 105 ··· 107 ~50 ~1 1000 [OVI]X-rays

Gas phases:


ISM consists of different gas phases, i.e. components with different temperatures and pressures.

Most of them are in mutual pressure equilibrium:

P = n k T [P] = dyn cm-2

or

P/k = n T [P/k] = K cm-3

Molecular clouds and dust do not participate in pressure equilibrium.

Molecular clouds are self-gravitating behave like stars

Dust has fm 1% and T 20 ···30 K

Relativistic component: cosmic rays, in pressure equilibrium with the gas, coupled to it via magnetic fields; CRs have fm 10-6

Average energy density of each component participating in pressure equilibrium:

‹u› = 1.6 10-12 erg cm-3 = 1 eV cm-3


Number density of CRs much lower:

assume en energy equipartition between particles and magnetic field; then

28

2cemrelnnE

Brel

2

2

1

1

c

cm

Be

ec

2

4

3

obtain Lorentz factor from critical frequency (observing frequency), estimate B-field from synchrotron intensity

e.g. B = 5 G, vc = 5 GHz nrel 1.6 ·10-10 cm-3

See also lectures by

L. Gregorini


The meaning of pressure equilibrium:

More details: see Appendix...

Appendix

Assumption is justified because the sound crossing time is much larger than the

- mean time between SN shocks

- recombination time

- cooling time


2. Atomic gas

Neutral hydrogen

Formation: in the 1st 3 minutes ... (BBNS)

21 cm line radiation of neutral hydrogen at frequency

10 = 1.42040575178(6) GHz

hyperfine transition, interaction of electron and nuclear spin magnetic dipole radiation with

A10 = 2.86888 ·10-15 s-1

spTk

h

eg

g

n

n

0

1

0

1

Since h · « k ·Tsp relative population always 3:1, dominated by collisions

Level population given by spin temperature Tsp:


Measure Tsp against con-tinuum sources involving on-off and frequency switching technique

)1(, eTeTT spconb

Towards continuum source:

Frequency switching:

)1()(, eTTT csponb

off position:

)1(, eTT spoffb

so that

offb

onbcs

T

TTT

,

,1

1

Note: =(), T=Tb()


Total HI mass

oMskm

d

Jy

SM

MpcD

HI

1

25 )(

1033.2

total dynamical mass:

oMskm

RM

kpcR

tot

2

15 )(

1031.2

HI mostly optically thin

210

1810823.1

cmatomsskm

d

K

TN b

HI

Kilborn et al. (1999)


A massive galaxy:


A dwarf galaxy (1/100 M31):

Not necessarily the true gas distribution

... molecular gas !?


LB ~ 0.5 LMW

LB ~ 0.06 LMW

LB ~ 0.005 LMW


Distrubed galaxies: how much mass?

e.g. NGC 4449

- Mtot ~ 2 ·1010 M (?)

- MHI ~ 2 ·109 M

- heavily disturbed by 109 M companion (DDO 125)

- irregular velocity field in centre

IZw 18 HI

Hunter et al. (1998)

Bomans et al. (1997)


Ionized hydrogen

Massive young stars emit photons with 912 Å ionize surrounding gas; these HII regions emit

• recombination lines (Ly, Ly, ... H, H, ... etc.)

• free-free radiation

radiative transfer (Rayleigh-Jeans approximation for h· « k ·Te):

2

22)(

)1()(

c

TkTB

eTBI

ee

e

ff

e gK

T

GHzcmpc

EM

5.1

4

2

662

10101001.3

hence: I 2 for » 1

I -0.1 for « 1

171010

ln2.1315.1

4

GHzK

Tg e

ff

gff = Gaunt factor for free-free emission Te = electron temperature EM = emission measure

0

0

2s

e dsnEM



M82 408 MHz Wills et al. (1997)

A prime example of free-free absorption: M82


)(/)(1

1

),(

1033.215.11.24

HNHeNK

T

GHzTfkHzT

T e

ec

l

Radio recombination lines:

from recombination rates obtain line temperature Tl:


Diffuse ionized gas (DIG)

DIG found out to large heights above galaxy planes - ‘DIG’ = ‘WIM’: ne ~ 0.02 cm-3 T ~ 8000 K traced by H - large scale height (Reynolds 1989):

390070 025.0015.0 cmeezn

zz

e

where z is measured in pc.

Observed H intensity along the l.o.s.:

)(10

)(36.0

9.0

40

3

2

RRayleighsdsK

T

cm

snI e

Le

H

1 R = 106/4 photons cm-2 s-1 sr-1 = 2.41·10-7 erg cm-2 s-1

by the way....: ne also from - rotation measures RM- pulsar dispersion measures DM

dsnDMdsBnRML

e

L

e 00


Haffner, Reynolds & Tufte


Dettmar et al. N

GC

3109

NGC4700


Problem of ionization: two possible (and necessary) mechanisms!

- photo-ionization - shock ionization - CR ionization

line ratios [X/H] (X = NII, SII, OI, OIII) indicate mixture of processes UV photons from HII regions travel large distances out of the plane

without being absorbed (and re-emitted) by dust

Association with star-forming activity is obvious fountains & winds

total energy of > 1042 erg s-1 required exceeds power of all SNe!

correspondence with radio continuum & polarization: magnetic fields ‘guide’ ionizing CRs into the halo

Appendix


3. Molecular gas

Molecular hydrogen

Since about 20 years it is known that hydrogen in the ISM consists at least as much of H2 as of HI!

maps of neutral hydrogen at = 21 cm yield an incomplete picture!

However: direct measurement of H2 difficult; symmetric molecule, lacks permanent dipole moment

Ground state 1+: both electrons in lowest orbital

Energy spectrum given by

- vibration Ev = e (v + ½)

- rotation Ev = Bv J(J + 1) - D J 2(J+1)2 Bv = h/(8·)

e vibration frequency

D stretching constant

Bv rotation constant

moment of inertia


Selection rules radiative and collisional transitions between

even (para H2 , J = 0, 2, 4, ...) and odd (ortho H2 , J = 1, 3, 5, ...)

levels strictly forbidden.

Transitions within each species allowed, in particular electric quadrupole transitions within v = 0, obeying J = 2

= 28 m, T = 512 K emission only in hot regions (shocks, stellar vicinity) otherwise absorption against (few) bright sources

molecular clouds have T 10 ···50 K (cold), 80 ···100 K (warm)

IR emission of H2 not representative for general ISM?


Significance of H2:

• H2 is the most abundant molecule in the universe

• a significant fraction of non-stellar baryonic matter in spiral galaxies is in H2

• H2 is an important coolant of diffuse gas from T 104 K down to T 100 K

• H2 cooling influenced structure formation in the early universe

• H2 infrared emission traces warm gas, collisionally and/or radiatively excited

• H2 promotes all interstellar chemistry

H2 in galaxies:

pervasive Tk ~ 10 ··· 30 K nH2 1000 cm-3

GMCs Tk ~ 20 K nH2 ~ 10 2 cm-3

dark clouds Tk ~ 10 K nH2 ~ 10 3 ···10 4 cm-3

cores Tk 40 K nH2 10 4 cm-3


oMcm

n

K

TM J

2/1

3

2/35

100102

At recombination, i.e. z 1100, MJ 103 ... 106 M (~ globular clusters)

H2 controls early structure formation in bottom-up scenario (Tegmark et al. 1997)

Molecular hydrogen is an indispensable ingredient to star formation, hence for the overall fate of the universe (as we witness it now)!

Requirement for structure formation early on:

cooling time << Hubble time cooling rate >> expansion rate

i.e., mean free path for interaction of particles and photons must be small enough

H(t) = dR(t)/dt / R(t) << cool

Tegmark et al. (1997)


Formation of molecular hydrogen:

Simply by ‘gluing together’ 2 hydrogen atoms? Basically yes, as

coll 1010 · (n/cm-3) s

clouds with n 10 ···100 cm-3 this would imply coll 103 yr

However, simple 2-atom collisions cannot form H2, since the formation energy (4.5 eV in the ground state) must be expelled.

Emission of photon not possible, since the only repulsive state with energy close to zero, the 3+ state, is not radiatively connected to the 1+ state; this would require a change of electronic spin!

How, then, dows it work?


Dust as a catalyst!

Reaction rate, i.e. rate to hit a dust grain

coll = (vH · ng · <d>)-1

plugging in typical values, one arrives at

coll 2 ·1012 · (n/cm-3)-1 s

vH = velocity of H atoms relative to (much more massive) dust grains

ng = number density of hydrogen atoms d = geometric cross section of dust

grains

in clouds with n 105 cm-3, a few H atoms will hit a dust grain per year (!); enough to convert all of the H atoms into H2 in a 103 few years!

once there is dust, H2 forms fast (dust has to have Td < 20 K); becomes efficient at nH2 ~ 105 cm-3

Another process: H + H- H2 + e-

about 103 less efficient, however important in early universe


Destruction of molecular hydrogen:

Ionization potential of H2 is 15.4 eV (larger than HI) destruction mostly via photo-dissociation.

Selection rules require two-step process for photo-dissociation:

(i) upward transition from 1+ ground state to higher bound electronic state, followed by

(ii) radiative de-excitation to vibrationally excited state leading to dissociation.

H2 simple process can be calculated; narrow lines related to bound states imply self-shielding of H2 against UV radiation

lifetime of H2 in standard interstellar radiation field

H2 103 yr, pd = 5·10-11 s-1 unshielded

H2 106 yr, pd = 5·10-14 s-1 for columns of 3 ·1020 mol./cm-2

Van Dishoek & Black (1988)


Carbon monoxide

Molecular hydrogen most important, but most measurements not representative.

Second-most abundant molecule: CO, with [CO/H2] 10-4

higher inertia lower transition frequency

(J = 1 0) = 115.27 GHz ( = 2.6 mm) 5.3 K above ground

(J = 2 1) = 230.54 GHz ( = 1.3 mm) etc.

Isotopomeres: 12C16O 13C16O 12C18O 13C18O 12C17O Abundances: 1 1:60 1:240 1:15000 ?

Formation and destruction of carbon monoxide:

CO mostly from OH + C+ CO + H+ chemically very stable, large ionization potential (14 eV) destruction by photo-dissociation, Ediss. = 11.1 eV photons with < 1120 Å, which implies 912 < < 1120 Å

self-shielding much less efficient than in case of H2 ; becomes efficient at NH2 > 1021 mol./cm-2 ; beyond that the main isotope is optically thick


Measuring H2 via CO

Underlying mechanism: excitation of CO by collisions with H2

For optically thin radiation, e.g. 13CO column density from measured brightness temperature Tb:

23.5

114 .

1

106.213

cmmol

e

skm

d

K

T

N

exT

b

CO

with 12CO, optically thick, determine Tex:

1

1

)1()(

'

''

Tk

h

cexb

ek

hT

eTTT

Tb : brightness temperature

Tex : excitation temperature

Tc : continuum background temperature

Measure Tex with 12CO ( » 1)

if we know [13CO/CO] in low-density regions total column density


That’s still not NH2 ...!

First determinations of NH2 using the virial theorem; stable molecular clouds:

potkin EE 2

22

2

1

2

1 irii

ikin MmE v = line width

ji

ir

ij

jipot r

MG

r

mmGE

,

2

So, for a homogeneous cloud:

G

rM ir

2

r = radius of cloud


For density distribution (r) ~ r-

oMskmpc

rM ir

2

133.02

40.02210

Measure total CO luminosity of molecular cloud at distance D:

s

pcskmKdTdDL bCO212

Once this has been established

• measure ICO NH2 or

• measure LCO Mvir MH2 (don’t forget to add HI and to correct for helium!)

Define

or dTI bCO

Milky Way: XCO = 1.5 ·1020 mol. cm-2 (K km s-1) -1


What does this mean?

• ICO measures (‘counts’) the number of individual clouds within the telescope beam, weighted by their

temperatures

• Mvir (the total cloud mass) equals the sum of the atomic and molecular gas mass

ICO is a good measure for the H2 column density (or LCO is a good measure for the H2 mass)

Tests: measure

• LCO, v, r correlation Mvir r ·v2 ?

• check extinction vs. measured gas column density:

N(HI+2H2) / Av = 1.8 ·1021 cm-2 mag-1

Solomon et al. (1987)

Guelin & Cernicharo (1987)


Other methods/checks:

Other methods:

• FIR & submm emission (Thronson 1986)

S ~ NHI + 2 · NH2

• -rays: interaction of CRs with hydrogen nuclei, subsequent 0 decay (Bloemen et al. 1986)

I ~ NHI + 2 · NH2 ~ NHI + 2 · XCO · ICO

inelastic collision of CR protons with hydrogen, roughly 1/3 of resulting pions are neutral, decaying into two -rays with mean energy of 180 MeV

nH 1 cm-3 predicts L 1039 erg s-1, close to what is measured!


Other methods:

• X-ray absorption: measure NHI and analyse spectrum of soft X-ray emission 2 · NH2

Exercise: decide whether we view NGC253 from ‘above’ or ‘below’....!

... from below!


Caveat: XCO depends on

• metallicity (C & O abundance, e.g. Wilson 1995)

• radiation fields (dissociation)

• density (shielding)

• angular, hence linear resolution (XCO depends on r and v)

• CR heating (Glasgold & Langer 1973)

heating by

- energetic particles (1 ··· 100 MeV CRs)

- hard X-rays ( 0.25 keV)

process: H2 + CR H2+ + e-(~35 eV) + CR

primary electrons heat gas by (ionizing or non-ionizing) energy transfer

heating rate (Cravens & Dalgarno 1978; van Dishoek & Black 1986):


Klein (1999)

bottom line: detailed case studies indispensable!

circumstantial evidence:

but: CR flux at E 100 MeV not known in galaxies ....

In any case:

• high densities, strong excitation, high metallicities : small XCO

(e.g. M82, ULIRGS & mergers)

• low densities, weak excitation, low metallicities : large XCO

(e.g. dwarf galaxies, halo gas)


a normal galaxy ...

a dwarf galaxy ...

Large Magellanic Cloud!

M51

Examples


NGC 4214 D = 4.1 Mpc

3 molecular complexes in distinct evolutionary stages

• NW : no massive SF yet; excitation process?

• Centre : evolved starburst; ISM affected

• SE : SF commenced recently; ICO as in NW canonical threshold column density for SF: NHI ~ 1021 cm-2

comparison with HI above 1021 cm-2 primarily molecular

H2 : self-shielding because of high density

dissociation of CO in photon-dominated regions (PDRs) atomic carbon [CI], [CII]

[CI] and [CII] are important coolants of the ISM

radiative decay of excited states


• WLM D = 0.9 Mpc: - little SF, weak radiation field & CR flux- XCO 30 XGal (Taylor & Klein 2001)

- below 12 + log(O/H) = 7.9 no CO detections of galaxies (Taylor et al. 1998)

Two contrasting examples:


• M 82 D = 3.6 Mpc: - intense SF, strong radiation field and CR flux high

gas density, large amount of dust- XCO ~ 0.3 XGal in central region (Weiß 2000) from

radiative transfer models; requires many transitions,including isotopomers true gas distribution

- strong spatial variation of XCO - blind use of XCO leads to false results ....


Ultra-luminous Infrared Galaxies (ULIRGS):

gas densities comparable to stellar mass densities in the centres of elliptical galaxies (Solomon et al. 1995)!!

tracers: molecules with high critical densities

(HCN, CS, etc.)


‘Measuring’ temperatures and densities

Local thermodynamic equilibrium (LTE) and Large Velocity Gradient (LVG)

LTE assumes Tkin = Tex = T, i.e. the same temperature everywhere and for all components

everything is ‘thermalized’

remember:

23.5

114 .

1

106.213

cmmol

e

skm

d

K

T

N

exT

b

CO

gu, gl statistical weights

column density of optically thin CO then


LVG approach: different molecular species may have different excitation temperatures

- assumes that optical thinness is provided by turbulence

- rotating clouds, spherical symmetry velocity is a function of distance from centre of a cloud, i.e. V = V0 · r/r0

- this avoids ‘line trapping’, i.e. photons emitted by certain molecular species in certain transition gets absorbed by the same species

- assuming the turbulence v » natural line width, then the photons emitted somewhere in the cloud can only interact with nearby molecules, reducing the global problem of photon transport to a local one


LVG requires many transitions of a molecule (J = 1 0, 2 1, 3 2, etc.) and its isotopomeres (12C16O 13C16O 12C18O 13C18O)

LVG code calculates for given (fixed) input parameters (abundances, velocity gradient, radiation field, beam filling factor) line ratios in the Tkin - nH2 plane

least-squares procedure finds the most likely Tkin and nH2


Weiß et al. (1999)

Distribution of molecular gas in M82


An effective path length in LVG:

L=|dv/dr|-1· v,

where v is the observed line width Velocity gradient and CO abundance are input parameters; then

Weiß et al. (1999)


Direct measurements of H2

Direct observation rendered difficult, owing to lack of dipole moment

Measurements with ISO SWS

e.g. NGC891 (Valentijn & van der Werf 1999):

S(0): J = 2 0 28.2 m S(1): J = 3 1 17.0 m

rotational lines, quadrupole transition, 512 K above ground

warm component : 150 - 230 K cooler component : 80 - 90 K

could amount to 5 - 15 times the HI mass significant fraction of DM!


Clouds near SF regions: PDRs

Photon dominated regions

precise structure depends on

- metallicity

- photon field

PDR models describe individual molecular clouds

not appropriate to obtain quantitative results for whole galaxies

galaxy would have to be synthesized from suitable ensemble of clouds

Van Dishoeck & Black (1988)


[CI], [CII] lines are tracers of dissociation of CO

[CII] line very important to study distant galaxies:

- high radiation fields - quasi independent of

metallicity

Transition [m ] [GHz] Tcool [K] ncrit [cm-3] Layer

[CI] 3P1 3P0 610 492 24 3P2 3P1 371 809 39

[CII] 2P1/2 2P3/2 158 1899 92

[OI] 3P2 3P1 185 1620 98 3P2 3P1 63 4757 228

CO J= 1 0 2600 115 > 5.3

3·103 surface

3·103 surface

> 105 intermed.

3·103 core


Structure of molecular clouds

HI : thick disk, FWHM 260 ··· 440 pc

CO : thin disk, FWHM 150 pc

Clouds have fractal structure:

M r3- = 0.3 ··· 1.3

mass spectrum:

dN/dM M- Heithausen et al. (1998)

= 1.5 ··· 2.0


4. Dust

Dust has cardinal importance for the evolution of the ISM

• catalyst in formation of H2

• fate of molecular gas in star-forming regionsregulates strength of radiation field

influences star formation

Formation: gas shed by red giant stars; gas cools and forms mostly oxygen-rich molecules seed molecules coalesce to dust particles; at high gas densities, material condenses out onto dust dust particle growth

Composition:

Graphites : at pressures of 102 ··· 103 dyn cm-2 free carbon (i.e. C, C2, C3 etc.) condenses and grows to unisotropic graphite particles, i.e. Cn, n » 1.

Silicates : heat-resistent silicates condense at temperatures below 1600 K, e.g. Ca2SiO4, Al2SiO4, Mg2SiO4


Signatures:

• extinction

• polarization of starlight

• (sub)mm/FIR emissionnote: FIR not a tracer of dust mass, but rather of ‘re-processed’ starlight


Böttner et al. (2001)

)(

2

dd TB

DSM

Dust mass best determined in the mm/submm regime; measure total continuum flux density:

2

)(

D

TBMS dd

00

1.5

dust absorption coefficient

Total dust mass then:

Fit parameters: - Td ( 2 components)- dust composition


NGC 4449 (center):

Böttner et al. (2002) fit 3 dust temperatures: 138 10, 39 3, 16 2 K

MHI ~ 1.5 ·108 M

MH2 ~ 4.4 ·108 M

Md ~ 1.8 ·106 M

Mg/Md ~ 330 (accounting for He)

10 + log(O/H) = 8.2

Note: Galactic Mg/Md ~ 150

XCO ~ 13 XGal Lisenfeld et al. (2002)NGC 1569:

Lisenfeld et al. (2002) propose lack of PAHs, owing to strong radiation field

XCO ~ (25 - 30) XGal

Md ~ 3.2 ·104 M

Mg/Md ~ 1500 - 2900!

10 + log(O/H) = 8.2


Polyaromatic hydro carbons (PAHs):


Dust in a normal, massive galaxy: NGC891

CO and dust at 850 m and 450 m (Israel et al. 1999; Alton et al. 1998) exhibit similar distributions

Alton et al. (1998)

Israel et al. (1999)


5. Hot gas

Existence

Hot phase of ISM postulated by Spitzer (1956) to explain existence of neutral gas clouds outside (above/below) the MW disk; if stable, they need to be held together by pressure of a hot surrounding gas:

n1 · k · T1 = n2 · k·T2 n = 10-2 · · · 10-4 cm-3

T = 105 · · · 107 K

Simple hydrostatic model T 106 K at 10 kpc from the plane.

Existence of this component meanwhile confirmed by numerous observations:

• interstellar absorption lines of highly ionized elements

• X-ray emission: thermal bremsstrahlung and emission lines


Cooling the gas

line radiation and thermal bremsstrahlung

shortest cooling time below Te = 105 K, depending on metallicity of plasma

above Te = 105 K, bremsstrahlung dominates:

Heating the gas

Heating sources?

• Stars : hottest have T 106 K

• PNe : up to 1.5 ·105 K

• Shocks : strong shocks (M >> 1) shock waves provided by stellar winds and SNe; e.g. SNe II; ESNII = 1051 erg, SN rate e.g. 1 every 100 years LSNII = 3 ·1041 erg s-1; 1/3 radiation, 2/3 mechani-cal energy

• magnetic reconnection: short time interval strong particle acceleration (shocks & magn. reconnection relevant in the solar corona)

yrK

T

cm

n eecool

2

1

8

1

3310

1010105.8

typical values for clusters of galaxies


Ion obs [Å] Eion. [eV] [X/H] T [K] (a) (b)

C+3 47.9 - 64.5 -3.44 1.0 ·105

N+4 77.5 - 97.9 -3.95 1.8 ·105

O+5 113.9 - 138.1 -3.07 2.9 ·105

(a): solar abundances

(b): temperature of gas in thermodyn. equilibrium at which ion has max. relative abundance

1548.195 1550.770

1238.821 1242.804

1031.926 1037.617

Observing the hot gas

UV absorption lines (IUE, FUSE) and soft X-ray emission (ROSAT, CHANDRA, XMM)

UV absorption lines:


X-ray emission:

• emission lines with H- or He-type spectra

dominates below T 5 ·106 Ktransitions down to n = 1 K-series (, , , ...) n = 2 L-series

(, , , ...) etc.• thermal bremsstrahlung

dominates above T 5 ·106 Kcontinuum spectrum similar to radio free-free emission, with exponential tail at high energies

emission coefficient:

),(3

2

3

323

62

effTk

h

eee

ie TgemTkcm

nneZe

dsI

h

Tkg e

ff 4

9ln

3

X-ray intensity:


X-ray model spectra with all ingredients (metals)

• optically thin line features

• exponential cut-off at high energies

absorption by hydrogen omitted here


Examples:

NGC 1569

M 82

NGC4631


Hot gas in clusters of galaxies

X-rays and optical

• ejected by early (dwarf?) galaxies

cool » H0-1

• heated by - cluster merging - galactic wakes?


6. Heating and cooling

ISM expected in stable disk configuration if all phases cool as quickly as they are being heated

if heating rate » cooling rate expansion of gas, blow-out and winds; worth reading: Wolfire et al. (1995)

Heating processes: deposit kinetic (thermal) energy in the gas

• photo-electric heating of small grains & PAHs; main heating process of diffuse ISM, UV photons absorbed by dust grains dislodge electrons, some make it to the surface and can be ejected;small grains and PAHs have low binding energies; without grains, there would be solely C, N, O, with high binding energies (difficult to heat)

• photo-ionization via species with ionization potentials below 13.6 eV (mainly CI)

h · + X X+ + e- + Ekin

• CR heating

H + CR(1 - 100 MeV) H+ + e-(Ekin 35 eV) + CR

• X-rays (similar to CRs)

Here: the neutral atomic phase in galaxies


Cooling processes: convert kinetic (thermal) energy of the gas into photons

• collisional excitation of fine structure lines; CII and OI the main coolants; dominates at T 8000 Kcollisional impacts of CII and OI with H, HII, H2, e-

additional cooling via CI, SiI, SiII, SI, FeI, FeII

• electron recombination onto positively charged grains; this occurs at moderately high (T 104 K) temperatures

• collisional excitation of H Ly and of low-lying metastable transitions of CI, CII, OI, OII, SiI, SiII, FeI, FeII; this occurs at the highest temperatures (T 104 K)

First models were guided by observe-ational evidence: early measurements of HI emission and absorption led to two-phase model. Field, Goldsmith & Habing (1969):

- cold nc 100 cm-3, Tc 30 K - warm nc 0.4 cm-3,

Tc 8000 K

in pressure equilibrium


Two-phase diagram:

detailed calculations of heating-cooling balance (Wolfire et al. (1995) show why a two-phase ISM basically always builds up

for an equilibrium pressure of P/k = 3000 cm-3 K, the following parameters are found:

Phase n [cm-3] ne/n T[K]

CNM 4.2 - 80 (13 - 3.2) ·10-4 210 - 41

WNM 0.1 - 0.59 (4.6 -1.3) ·10-2 8700 - 5500

Nota bene: in the unstable regime,

d(log P)/d(log n) < 0 !!

once a volume of gas undergoes a slight density decrease, the pressure increases, which in turn gives rise to a further density decrease because the region is then bound to expand!

In detail ....


Lower left: heating cooling


Element Tcool [K]* phase

FeVIII 6374 Å 106 hot ionized

OI 5003 Å 2 ·105 ionized

CI 610 m 24 neutral371 m 39 “

CII 158 m 92 “

OI 185 m 98 “ 63 m 228 “

HI 21 cm 20 “

H2 28 m 500 molecular

CO 2.6 mm > 5.5 “

Main coolants:

* optimum temperature for coolant to work


Extended to three-phase model: McKee & Ostriker (1977):

additional hot medium withT 106 K, fV 70%

system of hot bubbles ‘Swiss cheese’


Consistent with observations:

• numerous HI holes found in M31 (Brinks 1981), M33 (Deul & den Hartog 1990)

• meanwhile many other galaxies (see Walter & Brinks 1999) for a comprehensive review

Brinks & Walter (1999)

• hot gas filling the halos of galaxies




7. Galactic winds

Galactic winds:

• winds play an important role in the evolution of (small) galaxies (Matteucci & Chiosi 1983); may explain- metal deficiency of dwarf galaxies- (part of) enrichment of IGM- magnetization of the IGM (Kronberg et al.

1999)

• modern numerical simulations (e.g. Mac Low & Ferrara 1999;Ferrara & Tolstoy 2000):

for mechanical luminosity L = 1038 erg s-1

blow-out occurs in 109 M galaxy only ~30% metals retained


e.g. M82: high star formation rate high SN rate huge amount of mechanical

and radiative energy deposited in the ISM overpressure

e.g. M82: - LFIR = 1.6 ·1044 erg s-1 - LX = 2.0 ·1044 erg s-1 - SFR ~ 2 yr-1 - SN ~ 0.1 yr-1


Weiß et al. (2001)

Evidence for overpressured regions: expanding molecular superbubble in M82, broken out of the disk

result of high ambient pressure and dense ISM

main contributor to high-brightness X-ray outflow!

vexp 45 km s-1 Ø 130 pc

M 8 ·106 M Einp 1054 erg kin 106 yr SN ~ 0.001 yr-1

10% of Einp hot X-ray gas10% of Einp expansion of molecular shell


e.g. NGC1569:

LFIR = 8 · 1041 erg s-1

LX = 3 · 1038 erg s-1

SN ~ 0.01 ··· 0.001 yr-1

SFR 0.5 M yr-1

starburst ceased ~5 ··· 10 Myr ago

(Israël & de Bruyn 1988; Greggio et al. 1998):

Martin (1999)

partly vw vesc - H velocities (Martin 1998) - X-ray temperature (Della Ceca et al. 1996;

Martin 1999)

),(2 zResc


The hot gaseous halo of NGC4631

Wang et al. (2000)


Klein et al. (1991)

Magnetic fields

• B-fields in dwarf galaxies exhibit less coherent structure

• low-mass galaxies may have strong winds less containment for CRs (Klein et al. 1991)

Mühle et al. (in prep.)


Wind models for dwarf galaxies

Mc Low & Ferrara (1999):

- dwarfs with masses 106 M M 106 M, - mechanical luminosities L ~ 1037 ··· 1039 erg s-1 (over 50 Myr)- significant ejection of ISM only for galaxies with M 106 M - efficient metal depletion for galaxies with M 109 M

... Many more models

Recchi et al. (2001)

D’Ercole & Brighenti (1999)

.

.

.

Mac Low & Ferrara (1999) t = 100 Myr


8. Gas mass and b

20

00

3

8

H

G

Simplest form of Friedman equation:

m + k + = 1

Matter density:

m = DM + B

B tied to baryon/photon ratio

= nB/n = 2.88 ·10-8 · B · h2

h = H0/(100 km s-1 Mpc-1) = 0.72 0.08

pretty well-determined from helium (and deuterium) abundance, n well known from CMB

B · h2 = 0.019 0.0012 B = 0.04


HI is easily recovered, H2 more tricky

In galaxies: Mgas/M* 0.1 ··· 0.7 (massive ··· dwarf galaxies) not: ellipticals and dwarf ellipticals)

still uncertain, however, because of unknown H2 (see Chapter 3)

Total mass: eventually need to reconcile observations and theory; DM density profiles like, e.g., ‘NFW’ Navarro, Frenk & White 1997)

Blais-Ouellette et al. (2001)

Baryonic dark matter: perhaps numerous cold molecular clouds (Combes & Pfenniger 1997); X-ray absorption and eventually ALMA should disclose it ....

But this would work for galaxies only, where the problem is not all that severe!

And it turns out that galaxies contribute little baryonic mass on large scales


dr

T d

dr

d

m G

r T kr M

Htot

) (log ) (log) (

2

What about clusters of galaxies?

Total masses from- v of galaxies- gravittional lensing- X-rays

gas masses precisely derived from X-ray brightness profiles B

assuming hydrostatic equilibrium, the total mass is derived:

m = DM + B

• mass of galaxies insignificant!

• Hot gas: B/ (DM + B) 0.17

• forget about ‘good old M/L = few hundred’ in clusters of galaxies: It’s the hot gas in galaxy clusters that dominates the baryonic matter

Böhringer (1995)


Frenk et al. (1999)

N.B.: in galaxy clusters the relative contribu-tion of baryonic matter (the hot gas) increases with radius!

in galaxies it is the DM that increases with galactocentric distance!



Consider volume element with surface A in the galaxy plane and hight z vertical to it

since CR « gas , we only need to consider gas density gravitational force on the gas then:

zAgF gaszz

gz = component of the gravitational field plane treating the gas and CRs as a fluid, the force exerted by the pressure P can be written

AzPzPzzPzzPF CRgasCRgasP )]()()()([

Accounting for the magnetic field, we have for pressure balance

zgasmagCRgas gdz

dP

dz

dP

dz

dP

dz

dP

where gz < 0.

Appendix: pressure balance in galaxies

back


Solution via 3-D magneto-hydrodynamics ... or

... replace pressure gradient by difference in pressure between ‘upper and lower edge’ of the disk, devided by its half-thickness h:

2

0h

zzgas

zmagCRgasg

h

PPP

Observations reveal that this is roughly fulfilled. Test? use Poisson equation to estimate gravitational acceleration:

G42

where = gravitational field and * = mass density of stars (which provide the disk potential); since

G

zRRR

RR4

112

2

2

2

2

zRz,

We arrive at

Gz

42

2

where zgz

back


Within |z| 100 pc, the stellar density is roughly constant, so that

G

z

g z 4

implying

zzGg z 4

Observations yield * = 0.15 M pc-3 = 1.0 ·10-23 g cm-3 10-29 s-1 observed gas density gas = 0.05 M pc-3 = 3.3 ·10-24 g cm-3 and we know that

0

0z

z

gasgas e

z0 250 pc

gas(z=h/2) = 2.7 ·10-24 g cm-3 (h 100 pc) so that

gas(z=h/2) · gz(h/2) = gas(z=h/2) ·(-·h/2) = 1.3 ·10-12 g cm-3 s-2

back


2

0h

zzgas

zmagCRgasg

h

PPP

back

We had to evaluate

which means

2122

103.122

cmdyn

hhPPP gasmagCRgas

Now evaluate pressures on the left-hand side of the equation:

Pgas = 1/3 · gas · v2 , v

2 8 km s-1

Pgas = 7 ·10-13 dyn cm-2

PCR Pmag = B2/8 = 3 ·10-13 dyn cm-2 (B = 5 G)

(Pgas +PCR + Pmag)z=0 10-12 dyn cm-2

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Gaseous Phases in Galaxies

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