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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 : M dust 0.1M gas 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 10 3 to 10 9 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) 1

Chapter 3 THE INTERSTELLAR MEDIUM Introduction

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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:
' 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 ejected by stars
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-
Heavy elements in the gas can be depleted into dust
The very low density of gas allows detection of
forbidden lines :
– low transition probabilities
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:
III – doubly ionised positive ion
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
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 to image. The dark areas close to the Galactic plane in the optical waveband represent the 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 largely unaffected by dust and gives a more accurate map of the distribution of stars, including the bright central bulge. Radio continuum images indicate hot ionised gas and high energy radiation from supernovae. Molecular hydrogen maps molecular clouds - the raw materials for star formation, Atomic hydrogen extends over almost the full 360 longitude range and is useful for mapping the outer reaches of the Galaxy, particular its rotation curve at large radii [image credit: Jodrell Bank, Leiden Dwingeloo, Max Plank Institute, IRAS, COBE, A. Mellinger, J. Friedlander, S. Digel, ROSAT, NASA Goddard Flight Center].
Cold/Warm Gas: the 21 cm Line of Neutral 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
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:
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-
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
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 right is the UV image only showing near-UV in yellow and far-UV in blue. Each image is approximately 100 kpc by 100 kpc. [Credit: NASA/JPL/Caltech/VLSA/MPIA].
T ∼ 10 to 100K, relatively high density (cold
dusty molecular clouds)
detect directly
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
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
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
ting photons (radiative decay)
– free/bound transitions → continuum radiation
– bound/bound transitions → emission lines
first excited level (n = 2) give the Balmer series
Transitions down to ground state (n = 1) give Lyman
series (in UV)
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 ( ” )
Hβ 4861 A ( ” )
Hγ 4340 A ( ” )
Hδ 4102 A ( ” )
Hε 3970 A ( ” )
– 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-
Some of the most prominent optical lines of HII re-
gions are:
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
sorbed by a neutral H atom → ionises H atom to
produce a free electron
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
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
• 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
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
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
late-type galaxies and are valuable for use in
measuring redshifts
Figure 3: The Orion Nebula, M42. The most famous example of an HII region. The gas fluoresces because of the UV radiation from the hot young stars, recently formed in a dense region of gas [Hubble Space Telescope: NASA, ESA, M. Robberto (Space Telescope Science Institute/ESA), Hubble Space Telescope Orion Treasury Project Team.]
Figure 4: The optical spectrum of the Orion Nebula, showing very strong emission lines from Hα (red/pink), [OII] (blue), and [OIII] and Hβ (both green).
– 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)
– 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 Helix Nebula, right. Gas has been ejected from a hot, evolved star and the ultraviolet radiation from the star ionises the gas. [Images from the Hubble Space Telescope. Ring Nebula: produced by the Hubble Heritage Team (AURA/STScI/NASA). Helix Nebula: produced by NASA, NOAO, ESA, the Hubble Helix Nebula Team, M. Meixner (STScI), and T.A. Rector (NRAO).]
Supernovae eject material at very high velocities
into the interstellar medium
significantly 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 the Veil Nebula, right. The Crab Nebula is a very young supernova remnant, produced by 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 and A. Loll (Arizona State University). Veil Nebula image from the 0.9m Burrell Schmidt Telescope at Kitt Peak National Observatory, Arizona: NOAO/AURA/NSF.]
stars or old evolved stars can show maser emission
– density ∼ 1014 m−3
– transitions are in the radio
– the overpopulated state decays by stimulated
emission → maser emission
– coherent emission - polarised
– OH and H20 masers are observed (e.g. in Orion)
– useful kinematic tracers
trons moving relativistically in a magnetic field
– observed in both optical and radio
– photons are emitted in the instantaneous direction
of electron motion
Spectacular sources of synchrotron emission are sys-
tems with jets → young stellar objects with bipo-
lar outflows, or active galactic nuclei, lobes of radio
If interstellar gas is seen in front of a continuum
source, light from the source is absorbed at certain
are seen in absorption
in the IR – probably caused by carbon molecules,
possibly polycyclic aromatic hydrocarbons (PAHs)
Cold interstellar CN molecules:
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
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
but at temperatures T ∼ 103 − 104 K and lower
peratures T ∼ 104 K and lower densities
• hot ionised medium – ionised gas (HII) at very
high temperatures T ∼ 105 − 106 K but very low
the long term
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
– 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
nebulae, zone of avoidance for galaxies at
low galactic latitudes
– it allows dense molecular clouds to cool - facilitates
star formation
ular hydrogen
called extinction
passing through an element of interstellar space will
experience a change dIλ in intensity Iλ due to extinc-
This is related to the change dτλ in the optical depth
τλ at the wavelength λ that the light experiences
along its journey by
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 τλ is 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 )
the magnitude system)
Iλ above) :
= C − 2.5 log10 Iλ 0 − 2.5 log10( e − τλ)
= C − 2.5 log10 Iλ 0 − 2.5 ln( e − τλ)
ln(10) τλ
So the observed magnitude m is related to intrinsic
magnitude m0 by:
mλ = m0 + Aλ
that the star would have in the absence of interstel-
lar extinction, and A is the extinction in magnitudes,
given by:
Note that Aλ depends on the photometric band
centred at 5500 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
The interstellar extinction law. The extinction caused by dust is plotted against wavelength and extends from the ultraviolet through to the near-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:
= (B0 − V0) + (AB − AV )
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
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)
Extinction is the absorption and scattering of light
by dust
AV is the extinction in the visual waveband (for ex-
E(B − V ) ≡ EB−V is the colour excess or ‘reddening’
the visual band ∼ 3.0 (see the graph on page 26)
Spinning dust grains tend to align with their long
axes perpendicular to the local magnetic field (Davis-
Greenstein Effect)
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
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
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
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
Figure 7: A star-forming HII region within M16, The Eagle Nebula. The blue-green colour from the mostly ionised gas is caused by the light of [OIII] and Hβ emission lines from neutral hydrogen atoms. The gas is being ionised by ultraviolet radiation from hot, young stars off the top of the picture. The dark pillars, in contrast, are regions 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 the absorption 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 image produced by NASA, ESA, STScI, J. Hester and P. Scowen (Arizona State University).]
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
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
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
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
galaxies means that star formation occurs preferen-
tially in the arms