Chapter 3

THE INTERSTELLAR MEDIUM

Introduction

The interstellar medium (ISM): the gas and dust

distributed between stars in a galaxy

In the Milky Way:

mass of gas � mass of dust : Mdust ' 0.1Mgas

ISM is generally a small fraction of a galaxy’s lumi-

nous mass:

' 0 % for an elliptical

' 1− 25 % for a spiral (increases from Sa to Sd)

' 15− 50 % for an irregular

Very diffuse: in the plane of the Galaxy, particle

number density ' 103 to 109 atomic nuclei m−3

Mixture of:

• gas remaining from the formation of the galaxy

• gas ejected by stars

• gas accreted from outside (such as infalling dif-

fuse gas or the ISM of accreted galaxies)

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Important (1) in galaxy evolution - gas promotes

star formation in denser regions – absorption by dust

allows molecular clouds to cool

Important (2) observationally – emission lines from

gas are prominent and can be used to observe dy-

namics

Chemical composition is about 90 % H, 9% He, plus

a trace of heavy elements (expressed by numbers of

nuclei)

Heavy elements in the gas can be depleted into dust

grains

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Spectroscopy of Interstellar Gas

Gas in the ISM readily emits detectable radiation

The very low density of gas allows detection of

forbidden lines :

– spectral lines not normally observed in the lab

– low transition probabilities

– in the lab the excited states get collisionally

de-excited before they can radiate

– but in the ISM, although collisional times are

� the lifetimes of the forbidden states, the huge

number of atoms in the ISM means that these

are commonly observed

In astronomy we use notation such as HI, HII, HeI

HeII and HeIII where:

I – neutral atom

II – singly ionised positively charged ion

III – doubly ionised positive ion

etc.

So, HI is H0, HII is H+, HeI is He0, HeII is He+,

HeIII is He2+, LiI is Li0, etc.

A negatively charged ion, such as H−, is indicated

only as H−, although few of these are encountered

in astrophysics

Square brackets indicate a forbidden line e.g. [OII]

3

Figure 1: The Milky Way within +/- 10 deg of Galactic plane (360 deg in longitude)in various wavebands. Note how the appearance of the Galaxy varies from image toimage. The dark areas close to the Galactic plane in the optical waveband representthe obscuring effect of dust. However, the same areas are bright in the infrared images,showing that blue light is preferentially attenuated by dust. The near infrared is largelyunaffected by dust and gives a more accurate map of the distribution of stars, includingthe bright central bulge. Radio continuum images indicate hot ionised gas and highenergy radiation from supernovae. Molecular hydrogen maps molecular clouds - theraw materials for star formation, Atomic hydrogen extends over almost the full 360longitude range and is useful for mapping the outer reaches of the Galaxy, particularits rotation curve at large radii [image credit: Jodrell Bank, Leiden Dwingeloo, MaxPlank Institute, IRAS, COBE, A. Mellinger, J. Friedlander, S. Digel, ROSAT, NASAGoddard Flight Center].

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Cold/Warm Gas: the 21 cm Line ofNeutral Atomic Hydrogen

Atomic hydrogen, HI , emits at 21 cm wavelength

(radio) from hyperfine splitting of ground state

– cool/warm ISM – T ∼ 10 to 100K in high density

regions, 103 to 104K in lower density regions

– ‘spin-flip’ transition – coupling of nuclear and

electron spin – forbidden line

Energy difference between the two spin states:

∆E = 9.4× 10−25 J = 5.9× 10−6 eV

producing emission with a rest wavelength:

λ0 = hc/∆E = 21.1061 cm

and rest frequency:

ν0 = ∆E/h = 1420.41 MHz

5

Transition probability, A = 2.87× 10−15 s−1, so lifetime

of an excited state is ' 1/A = 11 million years

21cm transition itself cannot be observed in a labo-

ratory

In the ISM, the 21cm line is observed primarily in

emission, but can also be observed in absorption

against a background radio continuum source

HI observations have many uses: one critically im-

portant application is to measure the orbital mo-

tions of gas to determine rotation curves in our own

Galaxy and in other galaxies

HI observations can map the distribution of gas in

and around galaxies - the 21 cm radio emission pen-

etrates dust

6

Figure 2: An example of 21cm radio emission in M83, The Southern Pinwheel Galaxy,type SABc, showing its extended disk at 21 cm radio wavelength in red on the left,with the UV image superimposed (near-UV in green, far-UV in blue). On the rightis the UV image only showing near-UV in yellow and far-UV in blue. Each image isapproximately 100 kpc by 100 kpc. [Credit: NASA/JPL/Caltech/VLSA/MPIA].

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Cold Gas: Molecules

Molecular hydrogen and other molecules

T ∼ 10 to 100K, relatively high density (cold

dusty molecular clouds)

Molecular hydrogen, H2, is rare and very difficult to

detect directly

– abundance controlled by ‘sticking’ of HI atoms

to dust grains

– no radio lines, so no direct way of tracing H2 in

cold dense gas clouds using radio observations

– but some H2 band absorption in the far UV can

be detected

– H2 can be photo-dissociated by UV radiation into

atomic hydrogen, HI

Other molecules do emit in radio/microwave

– they act as an indirect tracer of cold dense H2 gas

– CO is particularly useful - has strong lines at

1.3 mm and 2.6 mm from rotational transitions

– CO and H2 densities are similar so we can use CO

as a tracer for H2

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Hot Gas: HII Regions

Hot gas is readily observed in the optical via emis-

sion lines from largely ionised gas

HII regions are regions of partially ionised hydrogen

around hot young stars (O or B type), with T ∼ 104 K

These stars emit strongly in UV

Any UV (Lyman continuum) photons with wave-

lengths λ < 912 A will photoionise hydrogen produc-

ing H+, i.e. HII ions

The ions and electrons recombine to produce excited

hydrogen atoms

The electrons then cascade down energy levels, emit-

ting photons (radiative decay)

Photons emitted in UV, optical, infrared and radio

– free/bound transitions → continuum radiation

– bound/bound transitions → emission lines

Prominent optical lines from transitions down to

first excited level (n = 2) give the Balmer series

Transitions down to ground state (n = 1) give Lyman

series (in UV)

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Each series is labelled α, β, γ, δ, ..., in order of increas-

ing energy

Transitions from n to n − 1 levels are the strongest

i.e. α lines are the strongest

Lyman series lines of hydrogen are:

Lyα λ = 1216 A (in ultraviolet)

Lyβ 1026 A ( ” )

Lyγ 973 A ( ” )

Balmer series lines of hydrogen are:

Hα λ = 6563 A (in optical)

Hβ 4861 A ( ” )

Hγ 4340 A ( ” )

Hδ 4102 A ( ” )

Hε 3970 A ( ” )

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Collisional excitation can also occur

– not for H - no levels accessible at collision

energies characteristic of HII regions (T ∼ 104 K)

– but possible for NII, OII, SII, OIII, NeIII

[OIII] lines at 4959A and 5007A are particularly promi-

nent

Some of the most prominent optical lines of HII re-

gions are:

[OII] 3727 A [OIII] 4959 A

[NeIII] 3869 A [OIII] 5007 A

Hε 3970 A HeI 5876 A

Hδ 4102 A [NII] 6548 A

Hγ 4340 A Hα 6563 A

Hβ 4861 A [NII] 6584 A

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In HII regions one Balmer photon is produced for

each Lyman continuum photon from the hot star

– so observations of Balmer lines of nebula gives

number of UV photons from star

This happens because most H atoms are in the ground

state, and are therefore opaque to Lyman photons

but transparent to others

A Lyman continuum photon from star will be ab-

sorbed by a neutral H atom → ionises H atom to

produce a free electron

The free electron is then recaptured (free-bound tran-

sition), emitting a continuum photon depending on

which state it is captured into :

• If captured into the ground state (n = 1) → emits

another Lyman continuum photon – back to where

we started

• If captured into n = 2 → emits Balmer contin-

uum photon in going to n=2 – one UV photon

produces one Balmer photon – then decays to

n = 1 emitting a Lyα line photon which will al-

most certainly be absorbed again

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• If captured to n > 2 it can then decay to n=2

emitting a Balmer line photon, or directly to

n = 1 – but a transition to n = 1 emits a Lyman

line photon that can excite another ground-state

H atom, so the process repeats, eventually pro-

ducing a Balmer line photon

So each ionising UV photon from star (λ < 912 A)

will produce on Balmer (line or continuum) photon

HII regions and planetary nebulae also produce ther-

mal continuum radiation – free-free emission:

– the free electrons in the HII can interact with

protons without recombination

– electrons are accelerated, producing radiation

The resulting spectrum is not blackbody because the

gas is transparent to free-free photons: there is no

redistribution of the energy of the free-free photons.

In fact the spectrum is quite flat at radio frequencies

Strengths of the emission lines from HII regions can

provide information on temperature, density and

chemical composition of the interstellar gas

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Colour optical images of HII regions show strong

red/pink and green colours:

– the red and pink is produced mainly by the Hα line

– the green is produced by [OIII] and Hβ

– HII regions are seen prominently in images of

spiral and irregular galaxies

– their emission lines dominate the spectra of

late-type galaxies and are valuable for use in

measuring redshifts

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Figure 3: The Orion Nebula, M42. The most famous example of an HII region. Thegas fluoresces because of the UV radiation from the hot young stars, recently formedin a dense region of gas [Hubble Space Telescope: NASA, ESA, M. Robberto (SpaceTelescope Science Institute/ESA), Hubble Space Telescope Orion Treasury ProjectTeam.]

Figure 4: The optical spectrum of the Orion Nebula, showing very strong emissionlines from Hα (red/pink), [OII] (blue), and [OIII] and Hβ (both green).

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Hot Gas: Planetary Nebulae

Like other HII regions except:

– compact regions around hot evolved stars

– gas is ejected by star through mass loss

– UV photons from star ionise gas

– gas emits photons like HII regions (similar

emission process)

– relatively luminous, with prominent emission lines

– also observed in other galaxies

– useful for tracing distribution & kinematics of stars

Figure 5: Examples of planetary nebulae: the Ring Nebula (M57), left, and the HelixNebula, right. Gas has been ejected from a hot, evolved star and the ultravioletradiation from the star ionises the gas. [Images from the Hubble Space Telescope.Ring Nebula: produced by the Hubble Heritage Team (AURA/STScI/NASA). HelixNebula: produced by NASA, NOAO, ESA, the Hubble Helix Nebula Team, M. Meixner(STScI), and T.A. Rector (NRAO).]

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Hot Gas: Supernova Remnants(SNRs)

Supernovae eject material at very high velocities

into the interstellar medium

– this gas shocks, heats and disrupts the ISM

– low density components of the ISM can be

significantly affected

– dense molecular clouds are less affected

– hot gas from supernovae can even be ejected out

of the Galactic disc into the halo

Supernova remnants have strong line emission. They

expand into and mix with the ISM

Figure 6: Examples of supernova remnants: the Crab Nebula (M1), left, and part of theVeil Nebula, right. The Crab Nebula is a very young supernova remnant, producedby a supernova observed in the year 1054. The Veil Nebula is an older example.[Crab Nebula image from the Hubble Space Telescope: NASA, ESA, J. Hester andA. Loll (Arizona State University). Veil Nebula image from the 0.9m Burrell SchmidtTelescope at Kitt Peak National Observatory, Arizona: NOAO/AURA/NSF.]

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Hot Gas: Masers

Some very high density HII regions around young

stars or old evolved stars can show maser emission

– density ∼ 1014 m−3

– population inversion between certain energy states

of molecules due to radiative excitation

– transitions are in the radio

– the overpopulated state decays by stimulated

emission → maser emission

– coherent emission - polarised

– very narrow emission lines of high intensity

– OH and H20 masers are observed (e.g. in Orion)

– useful kinematic tracers

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Hot Gas: Synchrotron Radiation

Broad-band non-thermal radiation emitted by elec-

trons moving relativistically in a magnetic field

– observed in both optical and radio

– photons are emitted in the instantaneous direction

of electron motion

– polarised perpendicular to the magnetic field

Spectacular sources of synchrotron emission are sys-

tems with jets → young stellar objects with bipo-

lar outflows, or active galactic nuclei, lobes of radio

galaxies.

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Absorption Line Spectra

If interstellar gas is seen in front of a continuum

source, light from the source is absorbed at certain

wavelengths

A number of interstellar lines and molecular bands

are seen in absorption

Some absorption features are not well-understood

Particularly problematic are diffuse interstellar bands

in the IR – probably caused by carbon molecules,

possibly polycyclic aromatic hydrocarbons (PAHs)

Cold interstellar CN molecules:

– CN has rotational modes which produce radio

lines, like most heteronuclear molecules

– radio lines can be observed directly, but more

interesting are the optical lines that are split

because of the rotational modes

Optical observations of absorption by cold CN in

continuum spectra of background stars show rela-

tive populations of the rotational modes (from line

strengths) and hence the temperature of the CN

– temperature ' 2.7 K, i.e., these cold clouds are in

thermal equilibrium with the CMB

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Components of the Gaseous ISM

It’s convenient to divide diffuse gas in the ISM into

distinct components – also called phases:

• cold neutral medium – neutral hydrogen (HI)

and molecules at temperatures T ∼ 10 − 100 K

and relatively high densities

• warm neutral medium – neutral hydrogen (HI)

but at temperatures T ∼ 103 − 104 K and lower

densities

• warm ionised medium – ionised gas (HII) at tem-

peratures T ∼ 104 K and lower densities

• hot ionised medium – ionised gas (HII) at very

high temperatures T ∼ 105 − 106 K but very low

densities

The phases are pressure-confined and are stable in

the long term

Ionisation by supernova remnants is an important

mechanism in producing the hot ionised medium

Cold neutral medium makes up ∼ 50% of the ISM’s

mass, but very small fraction by volume

Supernova remnants, planetary nebulae, and giant

molecular clouds not normally included in these phases

because they are not in pressure equilibrium with

the other components

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Interstellar Dust

Consists of particles of silicates or carbon compounds

– relatively small, but broad range in size

– largest ' 0.5 µm in size with ∼ 104 atoms

– some have . 102 atoms – like large molecules

Dust has a profound observational effect – absorbs

and scatters light – extinction

– diminishes light of background sources e.g.dark

nebulae, zone of avoidance for galaxies at

low galactic latitudes

Dust in galaxies is important because:

– it allows dense molecular clouds to cool - facilitates

star formation

– it catalyses the formation of molecules e.g. molec-

ular hydrogen

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Interstellar Dust - extinction

The absorption and scattering of light by dust is

called extinction

Light of wavelength λ, specific intensity Iλ (i.e. flux)

passing through an element of interstellar space will

experience a change dIλ in intensity Iλ due to extinc-

tion

This is related to the change dτλ in the optical depth

τλ at the wavelength λ that the light experiences

along its journey by

dIλIλ

= − dτλ

Integrating over the line of sight from a light source

to an observer, the observed intensity is

Iλ = Iλ 0 e − τλ

where Iλ 0 is the light intensity at the source and τλis the total optical depth along the line of sight

What is the loss of light in magnitudes?

Magnitude m in some photometric band is related

to the flux F (i.e. intensity) in that band by:

m = C − 2.5 log10(F )

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where C is a calibration constant (see introductory

material in Lecture 1 for more information about

the magnitude system)

In the presence of extinction, for a particular wave-

length λ we have (substituting for the expression for

Iλ above) :

mλ = C − 2.5 log10( Iλ 0 e − τλ)

= C − 2.5 log10 Iλ 0 − 2.5 log10( e − τλ)

= C − 2.5 log10 Iλ 0 −2.5 ln( e − τλ)

ln(10)

= C − 2.5 log10 Iλ 0 −2.5

ln(10)τλ

= C − 2.5 log10 Iλ 0 + 1.086 τ

So the observed magnitude m is related to intrinsic

magnitude m0 by:

mλ = m0 + Aλ

where the intrinsic magnitude, m0, is the magnitude

that the star would have in the absence of interstel-

lar extinction, and A is the extinction in magnitudes,

given by:

Aλ = − 2.5 log10(e−τ ) = + 1.086 τ

Note that Aλ depends on the photometric band

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For example, for the V (visual) band (yellow-green,

centred at 5500 A) :

V = V0 + AV

For the B (blue) band (centred at 4400 A):

B = B0 + AB

Aλ is a strong function of wavelength and scales as

Aλ ∼ 1/λ (not as strong as Rayleigh Law ∼ 1/λ4)

There is much stronger absorption in the blue than

in the red → reddening by interstellar dust

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The interstellar extinctionlaw. The extinction causedby dust is plotted againstwavelength and extends fromthe ultraviolet through to thenear-infrared.[Based on data from Savage& Mathis, Ann. Rev. As-tron. Astrophys., 1979.]

Colour indices, e.g. B − V , are reddened so that the

observed value is:

(B − V ) = (B0 + AB)− (V0 + AV )

= (B0 − V0) + (AB − AV )

≡ (B − V )0 + EB−V

where (B − V )0 is the intrinsic value (no extinction)

and EB−V = AB − AV is the colour excess or redden-

ing, which tells us how reddened a source is, based

on the extinctions in the two magnitudes

For the V photometric band, AV ' 3 EB−V (as shown

in the plot above)

If the intrinsic colour, (B − V )0, can be predicted

from spectrum, then EB−V can be calculated using

EB−V = (B − V ) − (B − V )0

EB−V data can then be used to map the dust distri-

bution in space

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Extinction gets less severe for λ & 1µm as the wave-

length gets much longer than the grains

For sight lines through the Galaxy at the Galactic

poles: AV ' 0.00 to 0.05 mag

At intermediate galactic latitudes: AV ' 0.05 to 0.2 mag

In Galactic plane, extinction can be many magni-

tudes in V and UV (less in IR). Distribution can be

patchy (e.g. Baade’s Window, in the bulge)

Towards the Galactic Centre: AV � 20 mag

X-rays can pass through dust grains (AK ' 3 mag)

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Summary of terminology:

B ≡ mB, V ≡ mV etc

B − V is a colour

Extinction is the absorption and scattering of light

by dust

Aλ is the extinction at some wavelength, λ

AV is the extinction in the visual waveband (for ex-

ample)

E(B − V ) ≡ EB−V is the colour excess or ‘reddening’

AVEB−V

is the ‘reddening ratio’ for interstellar dust in

the visual band ∼ 3.0 (see the graph on page 26)

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Interstellar Dust: polarisation

Dust grains are not spherical

Spinning dust grains tend to align with their long

axes perpendicular to the local magnetic field (Davis-

Greenstein Effect)

Preferentially block light perpendicular to the mag-

netic field: extinction produces polarised light

Polarisation will tend to be parallel to the magnetic

field: polarisation measurements of starlight reveal

the direction of the magnetic field

Dust also reflects light, with some polarisation: ob-

servable as reflection nebulae, where faint diffuse

starlight can be seen reflected by dust

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Interstellar Dust: Radiation byDust

Dust absorbs light - warms the dust - re-emits the

light as black-body radiation (approximately)

So dust has diffuse black-body emission superim-

posed on reflected starlight spectrum

Wien’s displacement law states that the maximum

of the Planck function Bλ of a black body at a tem-

perature T is found at a wavelength

λmax =2.898× 10−3

TK m

This predicts that the peak of the black-body spec-

trum for dust at a temperature of T = 10 K will be

at a wavelength λmax = 290 µm

for dust at T = 100 K will be at a wavelength λmax =

29 µm

and for T = 1000 K will be at λmax = 2.9 µm

So radiation emitted by dust will found in the in-

frared, given the expected temperatures of dust

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Star Formation in the ISM

Stars form by collapse of dense regions of the ISM

under their own gravity

i.e. in cores of molecular clouds, where gas is cold (∼10 K) and densities relatively high ≥ 1010 molecules

m−3

Figure 7: A star-forming HII region within M16, The Eagle Nebula. The blue-greencolour from the mostly ionised gas is caused by the light of [OIII] and Hβ emissionlines from neutral hydrogen atoms. The gas is being ionised by ultraviolet radiationfrom hot, young stars off the top of the picture. The dark pillars, in contrast, areregions of cold, dense molecular hydrogen gas in which star formation is occurring.They are dark because the cold molecules emit virtually no light and because of theabsorption of light by dust mixed with the gas. The ultraviolet radiation is ‘burning’away the surface of the cold gas by photoionisation. [Hubble Space Telescope imageproduced by NASA, ESA, STScI, J. Hester and P. Scowen (Arizona State University).]

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A region of cold gas collapses when its gravitational

self-attraction > hydrostatic pressure support

For gas of uniform density ρ, the Jeans length λJ is

the diameter of a region of gas just large enough for

gravitational force to exceed pressure support:

λJ = cs

√π

Gρ

where cs is the speed of sound in the gas

The Jeans mass is the mass of a region that has a

diameter equal to the Jeans length:

MJ =

(4

3π(λJ/2)3

)ρ =

π

6ρ λ 3

J

The free-fall collapse time Tff , is the time taken for

a static cloud to collapse under its own gravity in

the absence of gas pressure

The free-fall collapse time for a spherically symmet-

ric distribution of mass with a total mass M and

initial radius R to collapse from rest is

Tff = π

√R3

8GM=

√3π

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1

Gρ

where ρ is the mean density before the collapse starts

(see also the introductory material in lecture 1)

Star formation can be self-propagating

Stars form, heat up and ionise cold molecular gas →outward flow of gas compresses gas ahead of it

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Causes instabilities – locally collapse to form new

stars

The enhanced density in the spiral arms of spiral

galaxies means that star formation occurs preferen-

tially in the arms

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