ISM Lecture 10

The Three-Phase Interstellar Medium

10.1 Why necessary?

Two-phase model (see Sect. 8.3) FGH model: Heating by cosmic rays => p.e. effect UV from stars: (T) Cooling by [C II] fine-structure emission: (T)

Thermal equilibrium: (T)=(T) => two stable solutions Cold Neutral Medium T=50-100 K, nH=10 cm-3

Warm Neutral Medium T~8000 K, nH=0.2 cm-3

But….. SN explosions destroy FGH model on timescales <106 yr => need new theory which includes SN activity

Also, vertical scale height of CNM,WNM can only be explained if the medium is turbulent, ~7 km s-1 => something must be stirring up the ISM continuously (mechanical rather than radiative energy)

Note that all theories of ISM are global theories, which try to explain ISM on average, but not locally

Overview of SNR evolution

Why do SN destroy 2-phase ISM?

Frequency S of SN events (not well known)

Probability Q that SN#2 goes off before #1 disappears

Recall

FGH model: no=0.2, vS,fade=10 km/s => tfade=5.2106 yr, Rfade=138 pc => Q=6.(S/10-13)

S rate

volumein Galactic disk

yr

kpcpcpcyr

31 150

1530010 2

13 /

() t E n

R E n

o

o

fade

fade

yr

76 pc

2 9 106510 32 0 27

510 32 0 16

. . .

. .

=> SN events destroy FGH model in 106 yr => need better theory of ISM!

Q R St4

33 fade fade

10.2 Overview of three-phase model

Q: what is ‘steady-state’ (in global statistical sense) of ISM subject to SN explosions at rate S10-13 pc-3 yr-1?

MO paper is difficult to read; simplified summaries of main ideas given in Appendix

McKee & Ostriker 1977, ApJCox 2005, ARAATielens Chap. 8.5-8.6

Effect of supernovae on ISM Overlapping SNRs create “tunnel” system of hot

ionized (“coronal”) gas threading through HI clouds, much like the “holes”in a “swiss cheese”

SNR evolves in isolation until it intersects a tunnel and connects to coronal gas; then pressure drops suddenly as SNR “vents” to tunnel system and contributes to pressure of coronal gas

Characterize topology by a “porosity” parameter Q with the fraction of the ISM occupied by coronal gas is f1-exp(-Q), with Q1

Supershells in 21-cm survey of high-latitude gas => large SNR

Colomb et al. 1980

Three-phase MO model of ISM

1. Cold H I clouds 1. CNM: filling factor f~0.025

2. Warm cloud envelopes 1. WNM f~0.1, ionization by very soft X-rays (?)2. WIM f~0.2, ionization by hot B stars

3. Hot ionized medium1. HIM f~0.7

All three phases are in approximate pressure equilibrium P/k=nT3000 cm-3 K

There is no explicit discussion of dense molecular clouds in MO model, since these have filling factor f10-3

Cross section of a characteristic small cloud

x = ne / n (ionization fraction)

Cold cloud cores give rise to optical absorption lines

Warm neutral medium has low ionization fraction, which is maintained by soft X-ray background (?)

Outer layer is largely ionized by stellar UV background

Small-scale structure of the ISM

A supernova blast wave is expanding into region shown

Small clouds within SNR evaporate

Large-scale structure of the ISM

Only SNRs and large clouds are shown

Altogether 9000 clouds occur in this volume

10.3 Physical ideas underlying MO modela. Pressure-balance

Various phases are in pressure equilibrium Pressure is maintained by SNRs

Suppose pressure is initially low SNRs overlap with internal pressure larger than ISM pressure pressure in ISM increases until SNRs overlap just when they fade

Pressure balance requires

Note that actual value of Q uncertain especially if magnetic fields taken into account; current estimates range form 0.25-0.75

This equation can be used to determine PISM

Q R St 4

313 fade fade

Calculation PISM

Use expressions for Rfade and tfade with vfade=cISM=velocity dispersion ISM in 105 cm/s; no=nH= density into which SNR expands =>

PISMnHx1.4mHc2 => Use

4

32 1 103 15

511 26 1 48

52 6

R t E n cofade fade3pc yr . . . .

P k n cISM o/ 170 cm K -352

n c n c co o 1 48

52 6

52 1 48

50 352. . . .( )

4

31 103 13

131R t S Sfade fade

3 pc yr /

Result PISM

Substitute in above equations =>

This result is about a factor of two higher than the average observed insterstellar pressure of 3500 cm-3 K. Above discussion did not include effects of evaporation of clouds which will slow SNR

P k S E cISM / . . . 6400 13

0 68510 86

50 24 cm K3

b. Mass balance

The various phases of the MO model exchange mass. In steady state, these mass flows must balance

Cloud material is lost by “thermal evaporation” due to heat conduction into clouds located within hot young SNR

Cloud material is created in dense cool shell of radiative SNR: clouds make holes in shell and collect material by condensation

WIM HIM

Cool dense shell

Cold cloud cores

WNMRadiative SNR

Cloud evaporation

Shell fragmentation

Heating

Ionization

Thermal evaporation

Coldcloud

Conductive heat flux into cloud

Mass flow away from cloud

Hot medium

If cloud small => evaporates If cloud large => survives + condensation of material on top of it => There is a critical radius for which evaporation = condensation Rcrit strong function of T: RcritnT2

For T>107 K, Rcrit~1020 cm => 30 pc => most clouds evaporate For T<106 K, Rcrit much smaller => less evaporation MO assume cloud size distribution

What happens to cold cloud inside SNR?

Supernova density vs. diameter

T higher in young SNRs

More clouds evaporate in young SNRs

n(HIM) should be greater in young SNRs than in old (larger) SNRs

c. Energy balance

MO assume that most SN energy leaves system as radiation. This is satisfied if SNRs enter cooling (radiative) phase before they overlap (which they do, both by radiation and cloud evaporation)

A small fraction of SN energy remains as kinetic energy of shells at time of SNR overlap. This source of energy input serves to maintain random motions of clouds. Ultimately, this energy is also converted to radiation following cloud-cloud collisions

d. SummaryInput parameters for MO three-phase model

1. Supernova rate S

2. Energy per supernova E51

3. Mean density of gas in disk

4. Ionizing EUV photon production rate per volume

5. Cloud size distribution and geometries

There are uncertainties associated with all these parameters

10.4 Successes of the MO model

Predicts pressure of the ISM correctly “ab initio” n vs. age (or R) relation observed for SNRs:

younger SNRs are denser O VI absorption lines: probably produced in

cloud evaporation interfaces, where cloud material is heated to T 5105 K

Successes of the MO model (cont’d)

Observed soft X-ray background in rough agreement with T 5105 K expected for cooling SNR gas = HIM

Electron densities in WIM envelopes 0.04 cm–

3 agree with observations (pulsar dispersion measures)

Predicted cloud velocity dispersion 7 km s–1 agrees with observations

10.5 Problems with MO model

MO model does not predict enough WNM; prediction f 0.1 vs. observation f = 0.2 - 0.5

Cold H I distribution is observed to be much smoother than in MO model

Clouds are more sheet-like than spherical Type II SNe are not random but occur in clusters

of 10 - 40 stars => superbubbles which create galactic fountain (see Sect. 10.6)?

Large fraction of type I SNe occur in disk only and overall SN rate uncertain

Problems with MO Model (cont’d)

For Q 1, expect 3000 supershells covering 50% of disk area on sky – where are they?

O VI absorption lines from cloud interfaces require large energy injection rate – where does it go?

Grain destruction in shocks very effective – how does enough dust survive?

Q 1 requires mean gas density n < 1 cm–3 – is n really that low?

10.6 Recent developments and modifications of MO model

Scale height of WNM and HIM larger than thought previously Expect ~20 pc from pressure equilibrium vs. ~200

observed (see following tables) McKee 1990: This increased “weight” of ISM can

be maintained if porosity Q is high, fHIM 0.5 Filling factors may vary with R and z in the

Galaxy, with mostly WIM at large z, not HIM? External galaxies (e.g., NGC 891) have thick

ionized disk

Filling factors with z

WIM constrained by new Hobservations by Reynolds (1999) at high latitudes => scale height ~1 kpc. Need O stars whose photons can travel over large distances (kpc) to ionize such a thick layer within MO model

Scale height ISM

Blitz 1990

“The Galactic corona”

Savage 1995

Recent developments (cont’d)

COBE / FIRAS all-sky survey surprisingly strong [N II] fine structure lines, consistent with large filling factor and scale height of WIM

Fixsen et al. 1999

Three-phase model revisited: disk

Cox 2005

Purple= mol. cloudsSolid green= cold H I cloudsHatched green= warm H IHatched green on yellow= diffuse warm H IIOrange= Hotter gas with OVIRed=Hot gas emitting X-raysBlue= star

Top three panels show common(mis)conceptions, bottom panelsshow features that need to beadded.

Three-phase ISM: vertical

Same colors as before;Top three panels havedifficulties.

Three-phase ISM: global

Same colors as before;Top two panels haveproblems

Summary filling factors in ISM

Little doubt that CNM, WNM, WIM, and HIM are all present in ISM, but filling factors are still controversial fWNM 0.1 - 0.5

fWIM 0.2 - 0.5

fHIM 0.8 - 0.1MO FGH

10.6 Gaseous Galactic halo and fountain

Some SNe occur well above Galactic plane and cannot cool efficiently if Thalo 106 K Can also happen with large OB clusters

This SN energy will be pumped into halo; it can only be removed mechanically Galactic wind

Thus there may be a Galactic wind with 1 Msun/yr and a halo with Thalo 106 K, nhalo 510–4 cm–3

Halo would contribute to X-ray background

/E M

E M E

FHG

IKJFHG

IKJ10

2 1

10 20051

51erg / 200yr

v yr

700 km / s

vesc2

Sun

esc

Observational evidence

High velocity clouds: large number of H I clouds at high latitudes with large velocities falling into Galaxy? HVCs may condense from Galactic wind and fall back towards its

point of origin

Presence of ions: C3+, Si3+, N4+, … lines seen towards stars in LMC and SMC (see Sect. 6.2) originate in boundary layers of halo gas and condensing clouds? Or in tenuous gas of T 12000 K?

Bregman 1980

10.7 The local ISM (within 200 pc from the Sun)

Interstellar absorption lines towards nearby stars very little neutral matter within 100 pc of the Sun, N(H) 1019 H atoms cm–2

Confirmed by polarization data n < 0.3 cm–3 Soft X-ray emission mostly of local origin O VI absorption lines also indicate that there is

a lot of hot gas locally

=> The Sun is inside a “local bubble” of hot, ionized gas!

Local bubble

Where did bubble come from? SN 107 yr ago? Is Sun located in unique surroundings? Is there cold cirrus (as seen by H I, CO, dust) inside bubble? If

so, how are they formed and how long do they survive?

Galactic neighborhood

100 pc

Orange=Mol. clouds

P. Frisch

10.8 Summary regions in ISM1. H II regions

Traditional H II regions surrounding early type stars T 104 K, n 0.1-104 cm-3, f small Heated and ionized by photons with h>13.6 eV, <912 Å Cooled by forbidden lines of atoms + ions [O III], [O II], [N II], …. Observed by optical lines of atoms and ions, radio continuum, radio

recombination lines Coronal gas: very hot, tenuous gas pervading ISM (=HIM)

T3x105 K, n 0.003 cm-3, f 0.6? Heated by shocks from SN and collisionally ionized Cooled by adiabatic expansion and X-ray emission Observed by X-ray emission, optical emission, non-thermal radio

emission, UV absorption lines of OVI

H II regions (cont’d)

Warm ionized gas: warm diffuse gas throughout ISM (=WIM) T 8000 K, nH 0.25 cm-3, ne/nH 0.7, f0.2?

Heated and ionized by stray photons from O and B stars<912 Å

Cooled by forbidden lines of atoms and ions Observed by broad optical and radio recombination

lines (e.g. H166) and pulsar dispersion measures

2. H I regions

Cold neutral clouds: H I clouds throughout ISM (=CNM) T 80 K, nH 40 cm-3, f=0.025, ne/nH 10-4

Heated by UV photons with >912 Å through p.e. effect Cooled by [C II] fine-structure emission at 157 m Observed by H I 21cm emission + absorption, optical and UV

absorption lines of atoms (e.g. Na) Warm neutral gas: warm diffuse gas throughout ISM

(=WNM) T 8000 K, nH 0.4 cm-3, ne/nH 0.15, f 0.1-0.5? Heated + partly ionized by soft X-rays + stray photons Cooled by atomic lines Observed by H I 21 cm emission

3. H2 regions

Diffuse molecular clouds: includes translucent and high-latitude clouds T 30-80 K, nH 100-103 cm-3, ne/nH 10-4, f0.01 Heated by UV photons with >912 Å through p.e. effect Cooled by [C II] fine-structure emission Observed by H I 21 cm emission, CO mm emission, optical + UV

absorption lines, IRAS 100 m Dense molecular clouds: dark clouds + GMCs

T 10-100 K, nH 103-106 cm-3, ne/nH 10-6, f0.0005 Heated by cosmic rays + newborn stars Cooled by mm emission from molecules such as CO Observed by mm lines from molecules, FIR and submm continuum

emission from dust

Most regions are in approximate pressure equilibrium. Traditional H II regions may be expanding. Dense molecular clouds are self-gravitating