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2a. The Interstellar Medium
2.1 Introduction
2.2 Dust in galaxies
• Definitions of opacity, optical depth, colour
excess
• MW extinction law (and LMC,SMC)
• Chemical composition
2.3 Overall Galactic extinction
• Mapping the dust distribution
• Extinction in other galaxies
2b. The Interstellar Medium
2.4 Observations of dust clouds
• Dark clouds (& Bok globules, EGGs)
• Reflection nebulae
2.5 Interstellar gas
• Atomic hydrogen (Lya, HI, local bubble)
• Molecular gas
• HII regions
• Coronal gas
• Stellar remnants - Planetary nebulae,
Supernova remnants
• Cosmic rays
2. The Interstellar Medium
Sources: KKOPD 4.5, 15 (whole chapter); CO 12.1 (very brief); a lecture on dust
by Keel; further reading BM 3.7
2.1 Introduction:
The Galaxy (and other galaxies) consist of stars, dust and gas (and DM)
stars – studied in detail in previous semester
interstellar medium (dust and gas) in this section
Why? The history review of history of determining shape and size of galaxies have
shown the importance of dust absorption and scattering, and Rayleigh scattering
in the determination of the true size of the Galaxy
• distances to object in the Galaxy
• mapping distances to other galaxies
• determination of the total light, hence luminous matter
• apart from being the ingredient in the formation of stars
Interstellar gas is filthy: compressed to density of ordinary air (factor of 1021)
density of smoke would be such that objects would disappear in haze at distance
of much less than 1m!!!
The effects of absorbing material in galaxies were recognized before the physical nature of
galaxies became clear.
A study by H.D. Curtis published in 1918 compared photographs of spirals in an obvious
inclination sequence, showing that a band of obscuring material lies in the disk plane
It became clear, especially from comparing with nearby edge-on systems such as NGC
891, that this layer is exactly what we see in our view of the Milky Way
Lund
panorama
of MW
(blue light)
NGC 891 (red light)
Effects of Dust
Important effects caused by scattering or absorption
Dimming:
scattering (redirection of light out of line of sight)
absorption (photons transformed to heat)
Reddening: as light travels through interstellar dust cloud bluer photons get scattered more
strongly than red ones
Heating: absorption is an important
source of energy for the ISM:
blue light  heats dust 
re-emits in IR
 many galaxies emit
large fractions of their
light in the IR
Dust has been observed in two guises: optical/UV absorption & far-IR emission
With the recent advance in NIR and FIR astronomy, a lot is being learned about dust and grain
properties, but also in mm-observations (molecules, especially CO)
Observations of the ISM
• Optical imaging and spectroscopy
• Absorption of star light (dimming) and reddening
e.g. Dark clouds, patches in galaxies (horsehead
nebulae, Coalsack)
• Scattered light
e.g. Reflection Nebulae from gas around stars
• Line emission from ionized gas (warm gas)
e.g. regions around hot young stars (HII regions),
Planetary nebulae, SN remnants
Observations of the ISM – cont.
• NIR and FIR observations of heated dust (e.g. IRAS,
Spitzer, Herschel, WISE)
• Radio, submm and mm
• Cold gas
• neutral gas (21cm transition of HI)
• molecular gas (CO transition at 2.6mm), also masers
in star forming region in radio (H2O)
• From polarization - magnetic fields (information about
grains)
• X-ray: Hot gas (interstellar and intergalactic)
Phenomena caused by the interstellar medium
(Table 15.2 in KKOPD)
Some recent highlights
Left: M31 in CO (cold molecular gas) obtained with 800 hours of observation at the
30m IRAM mm telescope at the 2.6mm transition line (Astron. Astrophys. July 2006)
Right: optical image
for comparison
?
M31 as seen with Spitzer,
the mid-infrared (3.5 – 8 μm)
Space Telescope launched
08/03)
Two highly obscured (AB = 22 mag) spiral galaxies discovered in GLIMPSE
(the mid-infrared Legacy Survey) obtained with the Spitzer Space Telescope
At
l = 317.04
b = -0.50
l = 316,87 b
= -0.60
At the heart
of
the GA?
2.2 Back to Dust in Galaxies
•Dust is important in star-formation& induced SF through interacting and merging galaxies
•Dust is being condensed in stars’ atmospheres and expelled by stellar winds
•Dust is a useful first proxy for phases of the interstellar medium that are harder to trace,
since cool gas and dust are coupled gravitationally through the drag of atom/grain collisions
NGC 3314 - Hubble Heritage image
AM1316-241
Definitions on extinction and optical thickness
(m – M) = 5 log r – 5
Recall: distance modulus:
(with r in pc)
as distance r increases
brightness decreases as 1/r2
apparent magnitude increases
Except: absorbing (and/or light scattering) medium along the light path  EXTINCTION
How does extinction depend on distance?
-Take source emitting light Lo into solid angle ω
- in the distance interval [r , r + dr], the extinction dL
is proportional - to flux L,
- the distance travelled in the medium
- and the absorbing material through
the opacity 
(Note:  has dimension 1 / m)
For  : zero  transparent
infinity  completely opaque
Definition of optical thickness (dimensionless)
dL = -a L dr
dt = a dr ® dL = - L dt
Integrating from source (L = L0, and τ= 0)

Definitions of extinction and optical thickness
Optical thickness is the depth of a medium in which the intensity of light of a
given frequency is reduced to a factor of 1/e
- roughly 2/3 of the light is absorbed within one optical thickness depth
 L falls of exponentially with increasing optical thickness:
Note: empty space is transparent  thus  = 0  τ does not increase 
L remains const.
With a flux density F0 of the star at its surface R and F(r) the flux density at a
distance r, we get:
R 2 -t
F(r) = F0 2 e
r
relating this to absolute magnitudes (thus for r =10 pc), we can determine the dist. modulus
m - M = -2.5 log
F (r )
= 5 log r - 5 - 2.5 log e - t = 5 log r - 5 + (2.5 log e)t
F (10 pc )
or
m - M = 5 log r - 5 + A
where A is the extinction in magnitudes
For A ≥ 0 is the total extinction in magnitudes due to the entire medium between
the star and observer. For constant opacity:
r
t = a ò dr = ar ® A = ar with a = 2.5a log e ...ext. in mag. per unit distance
0
a = 1.086
The change in magnitude due to extinction is approximately equal to the
optical depth along the line of sight
(since A ≈ r = τ)
Summary of definitions
opacity  [m-1]; characterizes absorbing medium (through reduction of light)
dL
= -a dr
L
For  : zero  transparent; infinity  completely opaque
optical thickness (dimensionless)
dt = a dr ® dL = - dt ® L = L0e - t
L
depth of a medium in which the intensity of radiation of a given frequency is reduced by a factor of 1 / e;
 2/3 of the light is being absorbed within one optical depth;
) L decreases exponentially - with increasing α (thickness of absorbing material
Using flux ratios and the distance modulus we got a value of the total extinction A in
magnitudes due to the entire medium between the star and observer
m - M = 5 log r - 5 + A with A = 2.5(log e)t
r
For constant opacity:
t = a ò dr = ar ® A = 2.5(log e)ar = 1.086ar
0
For a unit distance the extinction in magnitude A = 1.086
The change in magnitude due to extinction is approximately equal to the optical
depth dτ along the line of sight
Atmospheric Extinction (see lecture notes of first semester)
The Earth’s atmosphere also causes extinction: The observed magnitude m
depends on the location of the observer and the zenith distance of the object
- determines the distance light has to travel through the atmosphere.
- to compare different observations: compare the different pathes light has taken
to reach you (your telescope and detector)
- assume atmospheric layer of const. thickness
X = 1 / cos z = sec z
Where X is the airmass. The magnitude
increases (gets fainter) linearly with distance X
m = m0 + kX
Where k is the extinction coefficient
The extinction coefficient can be determined by
observing the same (standard) source several
time during the night with a wide variety in
zenith distances – usually well known for
established Observatories.
First dust extinction measurements by Trumpler
Trumpler 1930: first clear evidence of the existence of dust through the study of the space
distribution of open star clusters:
- absolute magnitudes of brightest stars: from spectra type  distance r could be calculated
from apparent magnitude via the distance modulus (m-M) or fitting of MS stars
- he also determined the linear diameter D from the apparent angular diameter (D = d r)
- more distant cluster appeared to be systematically larger than nearer ones
 Space is not completely transparent: star light is dimmed by some intervening material
m - M = 5 log r - 5 + A
Trumpler obtained in the Galactic plane: apg = 0.70 mag/kpc (recall A = a r)
Today: apg = 2mag/kpc thus 10mag extinction over a path of 5 kpc!!!
Colour Excess - but first:
Recap: Photometry, magnitudes, filters
-mv :visual magnitude (sensitive as in human
eye, yellow at 5500Å)
-mpg : photographic magnitude, more
towards the blue
-mbol : bolometric magnitude, integrated
over all 
-BC : bolometric correction (BC = mbol - m
or Mbol - M), [original: visual, to F5 ]
Colour Excess
Another effect caused by the interstellar medium is reddening: blue light is
scattered and absorbed more then red: the observed colour index (B-V) increases.
E.g. the observed visual and blue magnitude of a star are:
V = MV + 5 log r - 5 + AV and B = M B + 5 log r - 5 + AB
The observed colour index (same for apparent and intrinsic magnitudes) is:
and
( B - V ) = ( M B - M V ) + ( AB - AV )
... or
( B - V ) = ( B - V )0 + E( B - V )
... where
( B - V )0 = M B - M V
… is the intrinsic color of a star
E(B - V) = (B - V) - (B - V)0
… is the colour excess
Studies of the interstellar medium show that the colour excess E(B-V) is
practically constant for stars:
AV
R=
» 3.0
E(B - V )
This makes it possible to determine the visual
extinction if the colour excess is known
Colour Excess
(B -V ) = (B -V )0 + E(B -V )
(B -V )0 = M B - MV
and
E(B - V) = (B - V) - (B - V)0
where
…is the intrinsic color of a star
… is the colour excess
Observed colour
Colour excess
Measure magnitudes of
stars of same spectral
type in different filters to
get wavelength
dependence
A(l ) µ1/ l
Longer λ
Extinction as a function of 1/ λ
for wavelengths longer
than B-band
Colour Excess
(B -V ) = (B -V )0 + E(B -V )
(B -V )0 = M B - MV
and
E(B - V) = (B - V) - (B - V)0
where
…is the intrinsic color of a star
… is the colour excess
2175Å bump
A(l ) µ1/ l
Longer λ
Extinction as a function of 1/ λ
Milky Way Extinction
curve
Colour Excess
• Ratio of colour excess E(B-V)
and Av is ~ constant for stars
AV
R=
» 3.0
E(B - V )
• R does not depend on
properties of star or amount of
dust
• Useful for getting distances
from photometry alone
• Know intrinsic colour for
particular spectral type
• Measure observed colour
AV » 3.0EB-V
V - MV = 5logr - 5+ AV
• A(λ)  0 for very large λ (e.g. radio)
• Measure to optical/NIR to ~ 2μm from earth
• Shorter wavelengths (UV) from space
• Interstellar extinction
largest at shortest
wavelengths
• ~10% of optical in IR
• Negligible in radio
Optically obscured objects
can therefore be studied in
the IR and in the radio.
Extinction curve to shorter wavelengths
for MW, LMC, SMC
Extinction
•Reminder: Extinction caused by both scattering and absorption
• Absorption: radiant energy transformed into heat, re-radiated at
IR wavelengths corresponding to temperatures of dust particles
 reduced intensity and reddening of colour
• Scattering: the direction of light changed as f(λ), leading to
reduced intensity in the original direction of propagation (and
change in colour)
•Extinction caused by dust grains with D near wavelength of light
•Gas can cause extinction but much lower scattering efficiency
 scattering by total amount of gas negligible in interstellar space
(contrary to air molecules in Earth’s atmosphere  atmospheric
extinction)
Scattering
Assume that size a, refraction index m and number density n of particles is known
•all particles are spheres with radius a
 geometrical cross section is π a²
 the true extinction cross section of particles:
Cext = Qext p a 2
where Qext is the extinction efficiency factor (or extinction coefficient)
•Q is dimensionless
•Q depends on composition of dust grains and wavelength
•
consider volume element dV with length dl and cross section dA normal to the direction of
light propagation with particle density n
•
assume that particles don’t overshadow each other  then there are
in the volume element dV
n dl dA
particles
Scattering cont.
n dl dA particles
Cross section
They will cover the fraction
dt of the area dA :
Cext
n dA dl Cext
dt =
= nCext dl
dA
The intensity within the path dl will then be reduced, proportional to fraction covered:
dL = -Ldt with dt the optical depth ® total optical depth observer - star:
r
r
0
0
t (r) = ò d t = ò nCext dl = Cext nr
Where
n
is the ‘average’ particle density along the path.
In magnitudes:
A = (2.5 log e)t = 1.086t ® A(r ) = (2.5 log e)Cext n r = 1.086Cext n r
Note: can be used to determine mean dust density if the other quantities are known
(extinction and distance)
Scattering cont.
For spherical particles with radius a and refractive index m, the extinction efficiency factor Qext
can be calculated exactly.
In general:
Define
Qext = Qabs + Qsca
x = 2pa / l
(Absorbing and scattering efficiency factors)
with λ the wavelength of radiation, then
Qext = Qext ( x, m)
The exact expression is a series expansion in x that converges more slowly for larger values x.
•For x«1 (particles much smaller than λ):
Rayleigh scattering
•Else particles similar or larger than λ:
Mie scattering
Qext ~ const.
Cext µ a 2
Since
Cext = Qext p a 2
Qext µ a / l
Cext µ a 3 / l
Analogy: waves on the surface of a lake
- If waves much smaller than an obstructing object (an island), they are simply blocked
- If waves are much larger than object in their way (grain of sand), the waves pass by almost
completely unaffected (Qλ ~ 0)
Scattering cont.
Example:
Extinction efficiency factor for refraction index
m=1.5 and m=1.33 (water) for increasing
particle size wrt wavelength
For large x (x»1): Qext ≈ 2
Cross section is not just geometrical crosssection (Q = 1) as expected, because of
diffraction of light at edges of particle
• Extinction has wavelength dependence because cross-section for
scattering depends on λ
• Red light scattered less than blue light
• Causes stars to look redder than effective temp. suggests
Chemical composition of dust grains
The UV bump at 2175Å = 4.6μm-1
•Predictions for A from Mie theory work well for longer wavelengths (IR to visible)
Cext µ a3 / l gives A =1.086Cext nr µ1/ l
•but strong deviations in the UV around 2175Å - sharp rise as wavelength
decreases
The bump gives hint about composition of dust:
Graphite: well-ordered form of carbon, interacts strongly with light at 2175Å
•Unclear how big graphite particles can get in ISM, but abundance of carbon and
the resonance around 2175Å suggests graphite is a major component of ISM
NIR observations (ISO)
•Observations of dark absorption bands at 9.7 and 18 μm
•from stretching of silicates in SiO and Si-O-Si bonds (energy levels tend to be
grouped in closely spaced bands – broad features in the spectrum)
 Existence of Silicate grains in dust clouds and diffuse ISM
NIR emission
Observations of unidentified infrared emission bands between 3 and 12μm
 Due to vibrations in C-C and C-H bond (planar molecules with organic benzene
ring-like structures) known as
Polycyclic Aromatic Hydrocarbons (PAHs)
Polarization
Light from interstellar dust tends to be slightly polarized; typically a few percent
(amount depends on wavelength)
 Dust particles can not be spherical but must be elongated
 Dust particles must be somewhat aligned
 Suggests existence of a (weak) magnetic field (difficult to measure as Earth and
Solar B much stronger)
 By studying the direction of polarization in various directions one can map the
direction of the Galactic magnetic field (see next sheet)
The dust in the ISM seems composed of both graphic and silicate
grains ranging from 0.25 μm to several Angstroms, the size of PAHs
 But note: the dominant component of the ISM is gas: HI, HII, H2
Between 1-1.5kpc
Magnetic field of MW
More distant stars:
B is aligned with Gal. Plane
between 100-200pc
Nearer stars:
Plumes of B rise tens of
parsec above the Gal. Plane
From Zeeman splitting of
Doppler shifted 21 cm lines
 B ~ 10-10 – 10-9T
(about 1 millionth of
interplanetary field in solar
system)
NIR emission
Observations of unidentified infrared emission bands between 3 and 12μm
 Due to vibrations in C-C and C-H bond (planar molecules with organic benzene
ring-like structures) known as
Polycyclic Aromatic Hydrocarbons (PAHs)
Polarization
Light from interstellar dust tends to be slightly polarized; typically a few percent
(amount depends on wavelength)
 Dust particles can not be spherical but must be elongated
 Dust particles must be somewhat aligned
 Suggests existence of a (weak) magnetic field (difficult to measure as Earth and
Solar B much stronger)
 By studying the direction of polarization in various directions one can map the
direction of the Galactic magnetic field (see next sheet)
The dust in the ISM seems composed of both graphic and silicate
grains ranging from 0.25 μm to several Angstroms, the size of PAHs
 But note: the dominant component of the ISM is gas: HI, HII, H2
The Milky Way’s ISM
30 Doradus – Tarantula
Nebula in LMC (HST image)
•Extinction curves for Milky Way, LMC, 30 Doradus and SMC
•Increased extinction to shorter wavelengths
•interrupted by local maximum - the so-called 2175-Angstrom bump
The interstellar medium (ISM) of the Milky Way
- About 10% of the mass of the MW consists of interstellar gas
- Interstellar space contains about 1 gas atom / cm3
- about 0.1% of that in dust grains (gas and dust distribution are strongly linked)
- It is strongly concentrated towards the plane of the MW
- What (little) we know about its grain properties comes from
- the analysis of the extinction curve
- the normalized amount of extinction as a function of wavelength,
- derived from looking at pairs of stellar spectra with similar
temperatures but different foreground extinctions.
- The general extinction curve within each of the Milky Way, LMC, and SMC is fairly
well defined
- The overall increase to shorter wavelengths (approximately with absorption in
magnitudes inversely proportional to wavelength) is interrupted by
a local maximum
- The so-called 2200-Angstrom bump
Main properties of interstellar gas and dust
• Dust grains form in the atmospheres of stars of late spectral type (lower mass) and
then expelled by radiation pressure into space.
• Also form during star formation & possibly directly from atoms and molecules in IS
clouds
2.3 Overall Galactic Dust Extinction
(from galaxy counts to gas distribution to the DIRBE/COBE dust maps)
• Images of edge-on galaxies & MW show that dust is concentrated in a fairly
thin disk (~100pc for MW)
• very high optical depth as viewed through the plane
• IR observers quote values of AV=40 magnitudes toward the galactic center
• Disk not very optically thick in the vertical direction (ZOA is relatively narrow)
• Dust distribution is patchy: some areas at high & moderate galactic latitudes
where we can see out (Baade’s window at (ℓ,b) = (1º, -3.9º) with 1.26 < AV <
2.79; Stanek 1996)
• What is the radial and scale-height distribution of dust?
• How does it relate to the spiral structure?
• How does it affect the light from a galaxy's own stars and from background
sources?
Edge-on
spiral NGC
891 (again) in
the red
Recall: the sun is located near the central plane of the galactic dust layer
 Dust in direction of GP is very large (higher towards GB than Gal. Anticentre)
 Dust towards Galactic poles is low (0.1mag)
This is apparent in the distribution of galaxies  A band of about 20º width where hardly any
galaxies are
 ZONE OF AVOIDANCE
The Effects of dust and stars in the Galaxy
on external galaxies → smaller and fainter
Part of Takahiro Nagayama’s
PhD thesis 2004
(Nagoya University)
(1) Assume ‘homogenous’ dust layer that will give rise to a total extinction Δm magnitudes in
vertical direction. Then the total Galactic extinction at Galactic latitude b will be:
AV (b) = Dm(b) = Dm / sin b = Dm csc b
 A galaxy with the apparent magnitude m0 will
be dimmer by that amount
(2) Assume uniform distribution of galaxies in space. The number count per unit solid angle to
apparent magnitude m will increase as:
However, accounting for the Galactic foreground extinction, the observable number of
galaxies will be reduced:
log N (m, b) = log N 0 [m - Dm(b)]= 0.6[m - Dm(b)]+ C
= log N 0 (m) - 0.6Dm(b)
= C '-0.6Dm / sin b = C '-0.6 AV
Where C ' = log N 0 (m) does ‘NOT’ depend on Galactic latitude. So by counting galaxies
at various latitudes b the extinction can be determined. Lick counts: Δmpg = 0.51mag,
Determination in the blue by Sandage (1973):
AB = 0.132(csc b - 1) for b £ 50°
=0
for b > 50°
Relation of dust with interstellar gas (valid for AV < 1 mag):
•E(B – V) is proportional to column density NHI of interstellar hydrogen
atom (HI or H2)
N HI
E(B - V) =
•Bohlin et al. 1978; Kent et al. 1991:
5.8 ´10 25 m-2
E(B - V )
 Indicating constant dust to gas ratio :
= const.
N HI
Constancy down different lines of sight
fixed number and size of dust grains associated with given mass of
hydrogen
- obviously only true as a first approximation
Near the Sun the number density is:
n HI =106 m-3
Along line of sight of length (distance) d in kpc (1 kpc = 3.086 x 1019 m):
d -2
d
m Þ E(B -V ) = 0.53
kpc
kpc
Þ AV =1.6mag/kpc
N HI = 3.1´10 25
DUST MAPS
Burstein & Heiles 1982 - most of the sky lying outside the ZOA (|b| < 10º)
Method:
Count of number of
galaxies in any zone to
some limited
magnitude;
Problem: cluster of
galaxies compensate
for extinction.
Extinction linked to
neutral hydrogen from
21 cm column density
Problem: ionized and
molecular gas cannot
be detected.
Combining two methods
 accuracy of 0.01 in
E(B-V) or 10%
HI column density map
Hartmann et al. 1997; grid of 0.5º, supplemented
with the lower resolution southern sky map constructed by Dickey & Lockman (1990).
Log scale from 1019 to 2 1022 cm-2
Even newer HI data at: http://www.astro.uni-bonn.de/hisurvey/profile)
HI column density map
Parkes Galactic All Sky Survey (GASS) mapped entire Southern sky visible from
Australia (south of dec = 1 degree), 16 arcmin resolution
Log scale from 1019 to 2 1022 cm-2
Even newer HI data from GASS at: http://www.astro.uni-bonn.de/hisurvey/profile
HI column density map
Parkes Galactic All Sky Survey (GASS) mapped entire Southern sky visible from
Australia (south of dec = 1 degree), 16 arcmin resolution LSR -400 km/s and 500 km/s.
Log scale from 1019 to 2 1022 cm-2
McClure-Griffiths et al. 2009
https://www.sciencedaily.com/releases/2016/07/160725151144.htm
Hodges-Kluck, Miller, Bregman.
THE ROTATION OF THE HOT
GAS AROUND THE MILKY WAY.
The Astrophysical Journal, 2016;
822 (1): 21
DOI: 10.3847/0004-637X/822/1/21
DIRBE Dust map
http://irsa.ipac.caltech.edu/applications/DUST/
NASA Extragalactic
Database coordinate
and extinction
calculator
Dust emission
measured by
COBE/DIRBE and
IRAS/ISSA in the
infrared
Schlegel,
Finkbeiner &
Davis 1998
(Schlafly&Finkbein
er update 2011)
More direct than
Burnstein & Heiles
(1982)
The distribution of cataloged galaxies with D ≥ 1.3’
(The diameter limit for which this Aitoff projection is complete)
 Comparison with DIRBE dust extinction maps: hardly any galaxies for AB = 1mag
º
UGC in the north of δ ≥ -17.5° (Nilson 1973); ESO Uppsala south of δ ≤ -17.5°; MCG inbetween (VV &A 1963-197
Extinction in other galaxies
(Galactic or internal extinction Ai)
Classical test for effects of dust traces
back at least to Holmberg's 1958 paper
- using the surface-brightness versus
inclination test
- For transparent galaxies
surface brightness should vary with
apparent axial ratio a/b (same light
concentrated in a smaller area)
- For opaque galaxies we see only a
thin skin (outer layer closest to us)
- the mean surface brightness will be
constant with inclination.
Note: also affects isophotal diameters
Holmberg came to the reassuring
conclusion that dust effects for global
light were minor
Extinction in other galaxies
Can be described by the following law:
for a/b < 4.7
Where a and b are the major and minor axis. The maximum is invoked so that highly
inclined galaxies do not get over-corrections. The slope depends on galaxy type later
spirals having more dust) and strongly on the waveband (absorption diminishes with
increasing λ)
Typical values for the blue magnitudes (B) are (Sandage & Tammann 1981):
BUT observationally, Valentijn (1990 Nature 346, 153) analyzed surface photometry
of the entire ESO/Uppsala galaxy survey to conclude that galaxies are almost
optically thick, out through the optical disk.
More modern result of the study of extinction in a sample of Sc galaxies
Giovanelli, Haynes et al. 1994
Note:
Disks have an intrinsic thickness
and the relation between
inclination i and axis ratio a/b
cos i = b/a
relation does break down:
i = cos-1 [((b/a)²-q0²)/(1-q0²)]1/2
with
q0 ~0.2 (intrinsic thickness of
disk)
Aside on the inclination of spirals
• Holmberg (1958) assumed that disk galaxies can be
represented as oblate spheroids with
(b / a)2 - q02
cos i =
1- q02
2
where i is the inclination, (b/a) is the observed axial
ratio, and q0 is the axial ratio for an edge-on system
(c/a, for c the disk thickness)
• q0 can be a function of morphological type
• This assumes that galaxies have circular isophotes
(clearly not always e.g. M101)
• Less obvious cases
increase uncertainty in i,
particularly at low i, but
variations in q0 may be
equally important
Important not "just" for understanding how galaxies work, but for
• the distance scale (inclination corrections to magnitudes in Tully-Fisher)
• understanding dark matter (if we're seeing only half the starlight, that increases
the relative amount of dark matter)
• the evolution of QSOs (when does the cumulative absorption along a random
line of sight become so large that QSOs will disappear from optical surveys?)
Cardiff workshop: The Opacity of Galaxy Disks (1990)
• improved statistical analyses of various surface photometry samples
• models of radiative transfer in clumpy media
• studies of overlapping galaxies
• new measurements of far-infrared and sub-millimeter dust emission
 overall agreement: disks centrally quite optically thick, falling to transparent at the
edges, with dusty spiral arms and considerable variation among galaxies.
Look at huge Sandage and Bedke 1988 Book: Atlas of Galaxies
You will see galaxies behind the disks of nearby spiral galaxies
2.4. Direct Observations of Dust Clouds
(a) Dark clouds
Recall: dust is seen in spiral arms (in particular in the inner edge)
but we also have individual dust clouds (dark patches) - often close to young star clusters!!
Coalsack (Schirmer)
Horse-head nebula in Orion
(a) Dark clouds
Usefulness (apart from being pretty ):
• of huge interest for star formation and determination of progenitor material of
new stars
• Determination total dust content (total extinction) in cloud from star counts
within certain magnitude interval in and adjacent to a dark cloud
Wolf diagram:
Comparison region: star counts increase
monotonically
In or on dark cloud: increase to certain level,
leveling off, the renewed increase with expected
slope of log N
 Δm is extinction due to dark cloud (2 mag in
figure)
 Stars up to that level are in front of dark cloud
 Star after plateau are behind cloud
 Can be useful for distance estimates to DC
Dark globules
(or Bok globules)
e.g. The Eagle Nebula (M16)
optical
"Pillars of Creation“:
gas & dust sculpted and lit up
by bright, powerful high-mass
stars in the NGC 6611 young
stellar cluster
Sometimes see against
bright nebulae – new SF
Properties of Bok globules:
- high extinction AV ~ 10 mag
- low temp.:
T ~ 10 K
- high density:
n > 104 cm-3
- low masses:
M ~ 1 – 1000 Msun
- small sizes:
r ~ 1pc
Wide-field IR image (144) of the Messier 16 region excellent
spatial resolution (ISAAC instrument, VLT 8.2-m)
penetrates the obscuring dust  light from newly born stars.
2 of 3 pillars have very young, relatively massive stars in their
tips. Another dozen or so lower-mass stars associated with the
small "evaporating gaseous globules” (EGGs) scattered over
the surface of pillars.
HST
Dark globules
(or Bok globules)
e.g. The Eagle Nebula (M16)
optical
"Pillars of Creation“:
gas & dust sculpted and lit up
by bright, powerful high-mass
stars in the NGC 6611 young
stellar cluster
Sometimes see against
bright nebulae – new SF
Properties of Bok globules:
- high extinction AV ~ 10 mag
- low temp.:
T ~ 10 K
- high density:
n > 104 cm-3
- low masses:
M ~ 1 – 1000 Msun
- small sizes:
r ~ 1pc
•Site of formation for binaries & multiple star
systems
•Mostly H2, carbon oxides, Helium, ~1% silicate
dusts
•EGGs – dense gas clouds, shelter from
ionizing flux
HST
Reflection Nebulae
(about 500 known)
Note: Reflection nebulae & emission nebulae
often seen together
- sometimes both referred to as diffuse
nebulae
- clouds of dust which reflect the light of a
nearby star or stars.
The nearby star or stars are not hot enough
to cause ionization in the gas of the nebula
like in emission nebulae but bright enough to
give sufficient scattering to make the dust
visible.
The Witch Head (IC2118) at d~1000 ly, associated with the
bright star Rigel in the constellation Orion (top right – not
visible on image)
•Glows primarily by light reflected from Rigel.
•Fine dust in the nebula reflects the light.
•Blue because Rigel is blue color and dust grains reflect
blue light more efficiently than red.
Reflection Nebulae
(about 500 known)
•spectrum of RN is similar to that of the
illuminating stars (scattered light)
•Dust grains are carbon compounds (e.g.
diamond dust) and compounds of e.g. Fe,
Ni
•Often polarization observed (due to dust
alignment – stars hardly ever have
polarized light)
•Usually blue because the scattering is
more efficient for blue light than red
The Witch Head (IC2118) at d~1000 ly,
associated with the bright star Rigel in
the constellation Orion (top right – not
visible on image)
NGC 2264: the Christmas Tree Cluster
taken at the NSF's 0.9-meter telescope on Kitt Peak with the NOAO Mosaic CCD camera
Open cluster of stars embedded
in a diffuse nebula.
Cone nebula: bottom centre
Fox Fur nebula: upper left
Image created by combining H
(red-orange), oxygen [OIII] in
light-blue and sulfur [SII] in blue
violet
Press release December 2005
“Spitzer Unveils Infant Stars in the Christmas Tree Cluster:”
Spitzer spots a stellar snowflake on
the "Christmas Tree Cluster." The
infant stars appear as pink and red
specks in the snowflake-shaped
cluster at the center of the image.
Astronomers combined light from
Spitzer's IRAC and MIPS cameras
in this mosaic. (Photo:
NASA/JPL/SSC)
Quote: …. a spectacular new picture of a star-forming region called the "Christmas Tree Cluster,“
complete with the first-ever views of a group of newborn stars still linked to their siblings ….
In 1922, E. Hubble found 2 interesting relationships:
Spec. types: OBAFGKM
(1) - emission nebulae (EN) only occur near stars with sp. type earlier than B0
- reflection nebulae (RN) may be found near stars w. sp. Type B1 and later
(2) Relationship between angular size R of nebulae and the
apparent magnitude m of the illuminating star
5 log R = -m + c
 the angular diameter is larger for brighter stars (the const. takes account of the fact
that the longer an exposure the brighter the larger the nebulae will appear)
 Figure by vd Bergh 1966 based on
measurement on POSS:
- Points correspond to RN
- Line is regression with slope 12.0
Theoretical derivation of this relation
by assuming:
-illumination of dust cloud is inversely
proportional to d² of illuminating star
- Dust clouds are uniformly distributed in space
The theoretical derivation predicts a value of the const. c in terms of albedo (ratio of scattered
to incident light) and phase function of grains
Albedo quite high (not well determined as distance between star and RN not well known)
Note: even if DC not visible (anymore) as RN, it will reflect light (certainly if albedo is high)
visible as diffuse Galactic light (indeed observed) - 20-30% of total brightness of MW
2.5 Interstellar Gas
Discovery and Composition
Mass of gas is > 100 times larger than dust, but much less easily observed (does not cause
extinction). Existence surmised early 20th century
-1904 Hartmann: noted abs. lines in binary systems that are not Doppler-shifted by star
motions in accordance to motions of stars and other observed lines:
From interstellar gas clouds (various at various Doppler shifts) between star and observer
In the optical
Most common lines are neutral sodium Na (D1 and D2 at 5898.8 and 5890.0 Angstrom)
e.g. in star spectrum of HD14134:
-Two pairs of D1 and D2 lines, corresponding to
gas clouds in 2 different spiral arms
with
Δv = 30 km/s
… and singly ionized calcium, the H and K lines (Ca II H 3938.7 and Ca II K 3968.5 Angstrom)
More lines exist in UV, but the most common and strongest there is 1216 Angstrom (Ly)
 we now use Ly to trace inter’galactic’ gas clouds in the high redshift universe
in the optical domain at z > 3
Many of the atoms are found to be ionized – due to UV emission of stars, and fact that ISM is
not very dense; hardly any electron around for recombination
Note:
- Mostly
elements up to
Zn (a few heavier
elements have
been detected)
As in stars most
of the mass is
in H (70%) and
He(30%)
But heavier
elements are
less abundant!!
The probably
reside in the
dust grains
(hence not
produce
absorption lines)
(a) Atomic Hydrogen
(a) Lyman  studies of ISM:
Ly  : transition from ground state (n=1) to 1st excitation state (n=2) - 1216Å. Excellent probe
of ISM: the ISM is cool; most HI atoms are in ground state (n=1) and can be excited to n=2.
Note: contrarily to stars, which are hot, thus most H in excited state n=2; resulting for T=
10’000K, in strong Balmer line absorption series
But Ly  observations need to be made from space!! First observations with rocket in 1967.
Example: spectrum of ζOph,
strongest lines is Ly  absorption
(Morton 1975)
Copernicus satellite
Complementary to 21cm line observations, because of the difficulty of determining distances
to nearby hydrogen clouds, while for Ly  the distance to the star is often known.
Some results from Ly :
•The average gas density within 1 kpc: 0.7 atoms/cm3
•But between Sun and Arcturus (11pc) it is:
 Sun is in a low density bubble of ISM:
 The so-called “Local Bubble”!!
0.02 – 0.1 atoms/cm3
Interlude: The Local Bubble
Results from the Extreme UV Explorer EUVE (1992 – 2000)
– updating some very early results from 1975 during a historic link-up of Apollo (USA) and Soyuz (USSR)
observations of very near stars showed underdense hot region  interstellar ‘gas’ is not uniformly
distributed but unevenly, with cool dense clumps interspersed with hot low-density gas shaped like
‘bubbles’ and ‘tunnels’
The Sun seems to reside (travel through) in such
a low density bubble; “Local Bubble”
-A peanut-shaped cavity
- of low-density hot thin gas: ~ 0.001 atoms/cm3
with T ~ 1- 2 106 K (so 100 less dense than
normal ISM and 1000 - 100’000 times hotter)
- Contains about 200’000 stars
- Extent - ~ 100 – 200 pc
- Probably caused by multiple SN explosions that
occurred several 100’000 yrs ago
NASA launched a satellite in January 2003 - the Cosmic Hot
Interstellar Plasma Spectrometer, or "CHIPS" - to study the Local
Bubble.
Similar features observed
in other galaxies:
“Holes and Shells in the
ISM of the nearby dwarf
galaxy IC2574”, Walter &
Brinks, 1999, AJ
Optical image:
H I holes in IC 2574. The gray-scale map is a linear representation
of the H I surface-brightness map.
(b) The Hydrogen 21cm line
The 21 cm line is produced by the reversal of the spin of
the electron relative to proton in nucleus
(corresponding to 2 allowed values of the quantum
number ms = ±½) in the hyperfine forbidden spin-flip
transition (can be in absorption too)
Photon emitted from higher energy level (parallel spins):
λ = 21.049cm
ν = 1420.4MHz
It has extreme low probability: spontaneous de-excitation only every 3.5 1014s or 1.1 107yrs
- In the ISM collisions are more probable.
- They occur on timescales of 100’s of years (can result in excitation and/or de-excitation)
- Still there is enough HI that this transition does occur (typical number density of HI atoms
along line of sight N(HI) ~ 1021 cm-2))
All HI emission comes from outer space: best vacuums produced on Earth have such ‘high’
densities that collisions dominate and spontaneous 21 cm emission never occurs ….
21cm line emission predicted theoretically by H. van der Hulst 1944
 first detection observationally in 1951 by Ewen and Purcell
Important tool in mapping the location and density of HI, measuring velocities (Doppler) and
estimating magnetic field (Zeeman)
Although HI is quite abundant, the rarity of 21 cm emission (or absorption) means that the line
remains optically thin of large interstellar distances
Assuming that the line profile is Gaussian, the optical depth of the line center is given by
t H = 5.2 ´10 -19
Where
NH
TDv
- NH is the column density of HI (in units of atoms cm-2)
- T is the temperature of gas (in K)
- Δv is the full width of the line at half maximum (in km/s)
Since the line width is primarily due to Doppler (due to gas motions within the cloud or of the
cloud as a whole) Δv is usually expressed in km/s
As long as the line is optically thin, the optical depth is proportional to the neutral HI column
density
Studies of diffuse HI clouds indicate
temperatures of 30-80K
densities of 100 - 800 cm-3
masses of 1 – 100 Msun
Comparing the optical depth of HI and dust along the same LOS shows that
N H µ N d when AV < 1
Indicating (as mentioned before) that dust and gas are distributed in the same way – except
for AV > 1, where this relation breaks down
Studies of the 21cm line have revealed more about the properties of the ISM than any other
method. A lot has been learned about the spiral structure and rotation of the Milky Way – and
other galaxies.
Usually the 21cm line occurs in emission. It can be observed in all directions of the sky (due to
the large abundance of H). An observed line of sight is shown below: it reveals high intensity
peaks where we have lots of gas (spiral arms), but the clumps are also blue– or redshifted
 This has provided important clues about the differential rotation of the Milky Way and its spiral
structure
First overall distribution of the neutral hydrogen distribution in the Milky Way
from the Leiden and Parkes Surveys. Densities are given in atoms/cm-3
modern sketch of HI
distribution
including LMC, SMC
Oort, Kerr, Westerhout 1958. MNRAS 118, 379
HII regions
HII region = emission nebulae of ionized hydrogen
 Hot O stars radiate strongly in the UV; if they have (enough) HI around them, this will be
ionized  resulting in a so-called HII-region.
Some examples: The Lagoon Nebulae (M8): 5,000 lyrs from Earth in Sagittarius.
 First light SALT image with SALTICAM
UVI 120  40 s (Sept. 2005)
The hot, central star, O Herschel 36 [upper left], is the
primary source of the illuminating light for the
brightest region in the nebula, called the Hourglass.
The glare from this hot star is eroding the clouds by
heating the hydrogen gas in them. This activity drives
away violent stellar winds that are tearing into the
cool clouds.
The Orion Nebula, Messier 42
- H II region excited by 4 hot stars in the
Trapezium Cluster
- dark regions are opaque dust clouds in
front of the nebula
- Radio and IR show a rich molecular
cloud behind it (see later notes)
- upper part shows the reflection nebula
NGC1977 (The Running Man Nebula)
Close up of the Running Man: RGB combined image …by
Peter Spokes/Adam Block/NOAO/AURA/NSF
• HII regions form after star formation starts with a giant molecular cloud (GMC)
• Hot new stars ionize their surroundings
• HII regions last for a few million years before gas is dispersed
• Clouds of ionized gas typically have spectra dominated by a few emission lines and
a weak continuum - Strong Balmer lines
• An atom in an HII region remains ionized for several 100 years, but neutral only for
a few months after recombination before being ionized again
• when excited HII recombines, it slowly returns to ground-state via various
transitions
• Most recombinations in the optical from
the Balmer series n = 3  2
• Includes H at 6563Å - often use
narrow band filter images around H
(Recall Lyman series n =2  1 in UV: 1215
- 912Å, Paschen n= 4  3 in IR, 18750
- 8210Å)
• when excited HII recombines, it slowly returns to ground-state via various
transitions
• The number of recombinations per unit time and volume :
nrec = a ne ni with a = 3.1´10-13 cm3s -1 for T = 8000K
Where ne, ni are the density of electrons and ions and alpha is the recombination
coefficient (depends on temperature)
• For completely ionized hydrogen we have 
nrec µ ne
2
• Most recombinations lead to H emission from the 3 to 2 transition, so
• The surface brightness of a nebula in H is proportional to the Emission
Measure
EM = ò ne dl
2
(integral along LOS to the nebula)
Ionized Helium He+ or He++ (respectively He II or He III)
Ionization of He requires more energy than for hydrogen
 Only the very hot stars will have He+ regions
 Large HII regions will surround smaller He+ or He++ regions, which can also
be seen in the spectrum
Other strong emission lines (sometimes stronger than H and He)
called Nebulium early 20th century
•thought to be unknown element, but in 1927 I.S. Bowen showed forbidden
lines of ionized O and N
•Main ones are [O II] - 3726 Å, 3729Å (doublet)
[O III] – 4959 Å and 5008 Å
[N II] – 6583 Å (very close to Hα at 6563 Å)
•Don’t see in terrestrial laboratories – too dense so collisions are more likely
Strömgren Sphere:
- Ionization propagates in a sphere from central star
- UV radiation is absorbed very efficiently  very sharp boundary between HII and HI.
- In a homogenous medium, the HII region around a single star will be spherical, forming a
sphere with the size proportional to temperature of star
e.g. B0 V: R = 50 pc
A0 V: R = 1 pc
Take recombination rate nrec  multiply
by volume of HII region (assumed
spherical) = equal to ionizing photons
N produced per second  radius of
sphere:
1/ 3
æ 3N ö -2 / 3
rS @ ç
÷ nH
è 4pa ø
Recall:  proportional to T
As temperature of HII regions is higher
than surrounding gas, it will expand;
After million of years it will become
diffuse and merge with ISM
The Rosetta Emission Nebula
surrounding the open cluster NGC2244
Strömgren
Derivation of Strömgren sphere size
•Assume star embedded in uniform medium of neutral hydrogen.
•A sphere of radius rs around this star will become ionized - the “Stromgren
radius”
•Volume of sphere from setting the rate at which ionized hydrogen
recombines equal to the rate at which the star emits ionizing photons
i.e. all of the ionizing photons are “used up” re-ionizing hydrogen as it
recombines
•The recombination rate density is αn2
α is the recombination coefficient (in cm3 s-1) and n=ne=ni is the number
density of ions or electrons
•The total rate of ionizing photons (in photons per second) in the volume of
the sphere is N*. Setting the rates of ionization and recombination equal to
* ö1/3
one another, we get
so
3
æ
4p rs
a n2 = N *
3
3N
rs = ç
2÷
4
pa
n
è
ø
•Typical values: N* ~1049 photons/s, α ~ 3 x 10-13cm3/s and n ~10cm-3,
imply Strömgren radii of 10 to 100 pc.
Strömgren Sphere:
- Ionization propagates in a sphere from central star
- UV radiation is absorbed very efficiently  very sharp boundary between HII and HI.
- In a homogenous medium, the HII region around a single star will be spherical, forming a
sphere with the size proportional to temperature of star
e.g. B0 V: R = 50 pc
A0 V: R = 1 pc
Take recombination rate nrec  multiply
by volume of HII region (assumed
spherical) = equal to ionizing photons
N produced per second  radius of
sphere:
1/ 3
æ 3N ö -2 / 3
rS @ ç
÷ nH
è 4pa ø
Recall:  proportional to T
As temperature of HII regions is higher
than surrounding gas, it will expand;
After million of years it will become
diffuse and merge with ISM
The Rosetta Emission Nebula
surrounding the open cluster NGC2244
Strömgren
Because of Galactic extinction, only a few nearby HII regions are possible to
observe optically
 go to radio or IR:
• e.g. recombination of H, He
• Some fraction of recombinations lead to transitions between high energy
levels -> radio emission e.g. from E110 – E109 for H (5.01GHz)
• provides data on Gal. structure and rotation through velocities (Doppler)
and distances to these regions
• Not only line emission but also radio continuum emission
Bremsstrahlung or free-free emission from electrons
• produced by the acceleration of a charged particle, such as an electron,
when deflected by another charged particle (e.g. atomic nucleus)
• Bremsstrahlung has a continuous spectrum with intensity
I µ EM = ò ne dl
2
• strong IR continuum emission from thermal radiation of dust inside nebulae
2b. The Interstellar Medium
2.4 Observations of dust clouds
• Dark clouds (& Bok globules, EGGs)
• Reflection
nebulae
NASSP
application
2.5 Interstellar gas
Deadline 31 Aug
• Atomic hydrogen (Lya, HI, Local Bubble)
http://www.star.ac.za/
• HII regions
• Molecular gas
• Coronal gas
• Stellar remnants - Planetary nebulae,
- Supernova remnants
• Cosmic rays
Interstellar Molecules
1937, 1938: First molecular absorption lines discovered in stellar spectra: CH, CH+, CN
1970 others in UV: H2 and CO (first discovered in radio at 2.6 mm)
 molecular hydrogen H2 is the most
abundant molecule, followed by carbon
monoxide CO
The detection of molecular hydrogen:
one of the biggest breakthroughs in UV
astronomy
•strong absorption band at 1050Å
• first observed with rocket in 1970
(Carruthers)
•large fraction of ISM hydrogen is
molecular
•fraction increases strongly for denser
clouds (and higher extinction)
 for AV > 1 mag most of the
hydrogen is molecular
Radio Spectroscopy
Note: absorption can only be observed if bright star is ‘behind’
molecular cloud
•because of dust extinction observations of molecules can not be made
in optical and UV!
•Until 1960s it was thought that at most di-atomic molecules existed
(gas too diffuse and UV radiation to strong to allow existence of more
complicated molecules)
•But 1963: discovery of hydroxyl radical OH
•As of 2002: about 130 molecules detected
•The heaviest is HC11N with 13 atoms
• Molecular lines in the radio regime occur in either absorption or emission
• Three kinds of transitions (see Table 15.4 KKOPD, 2 slides back, for examples):
1. Electron transitions: correspond to changes in the electron cloud of the
molecule (similar to single atoms) - found in optical and UV
2. Vibrational transitions: changes in vibrational energy of molecule generally in the IR
3. Rotational transitions: changes in rotational energy of molecule
(molecules in ground state do not rotate; they have zero angular
momentum; but they may be excited by collisions - generally in mm and
radio
• Many only been discovered in the densest clouds (like Sag B2 cloud in Galactic
centre) while others are more common
• The most common (H2) cannot be observed in the radio
• next best are CO, OH and NH3 (ammonia)
• Despite their relative abundance being only a small fraction of H2, it is
more than sufficient for detection in dense clouds
(e.g. Sag B2 clouds contains enough ethanol C2H5OH for 1028 bottles of
Vodka)
Example: radio map of the distribution of 13C16O in the very rich molecular cloud near the Orion
Nebula. The contours are lines of constant intensity (Kutner et al 1976, ApJ); right where we
have the dark clouds in the optical image on sheet 53)
Most of the molecules detected in dense molecular clouds near HII regions, not in
the actual HII regions (dissociation by high T and strong UV radiation)
Three types of molecular sources have been detected near HII regions:
1.large gas and dust envelopes around HII region
2.small dense clouds within these envelopes
3.very compact OH and H2O maser sources
The large envelopes have been discovered primarily in CO; also OH and H2CO
(formaldehyde)
•Large size and density (n ~ 103 -104 molecules/cm3)
•Masses of 105-106 Msun (e.g. Sgr B2)
Among the most massive objects in the MW
IR observations of thermal dust radiation:
•Peak at 10-100μm  T ~ 30-100K
Masers
Some IS clouds contain small maser
sources
• OH, H2O and SiO emission a few million
times stronger than elsewhere
(amplified stimulated emission)
• radiating regions of only a few AU (5-10)
• related to star forming regions
Central part of Orion Nebula: large crosses
indicate OH masers, small crosses H2O
masers
Note: compared to earlier image, this
region would only be a few mm in size
Interlude: Relation of dust and gas to the formation of
Protostars
•
Milky Way: M = 1011 Msun and age = 1010yrs  average SFR : 10Msun / yr
•
This is an upper limit for present rate. Recall that O stars have a lifetime of a million years
(106 years) only - earlier SFR must have been higher
•
Counts of O stars indicate current SF rate at 3 Msun / yr
 stars form in dense IS clouds (mostly located in spiral arms)
 contraction of cloud under its own gravity  fragmentation into parts  protostars
•
Observations: stars are not formed individually but in loose associations (recall Christmas
tree siblings) of a few 100s of stars born ~ simultaneously
•
Theory: calculations show it is difficult to form single stars:
 contraction of IS cloud: if mass (gravity) is larger than pressure
1920 James Jeans: cloud with certain T and density can only collapse if mass is
above a certain limit, the so-called Jeans limit:
3
M J » 3´10 4
T
M sun
n
With n = the density in atoms/m3, and T = temperature
In typical ISM Hydrogen cloud: n = 106, T = 100K  MJ ~ 30 000 Msun
- only in the densest and coldest clouds with n = 1012 and T = 10 K do we get 1 Msun
• SF starts in clouds of a M ~ x 1000 Msun and diameter of D ~10pc
• Contraction but no heating (optically thin – liberated energy gets carried away)
Density increases – Jeans mass decreases  separate condensation nuclei
form, each contracting individually
3
 cloud fragments
M J » 3´104
T
n
M sun
• Fragmentation is enhanced by ↑ rotation (conservation of ang. momentum L with
contraction)
• Continuation of this process – until individual condensations become optically
thick
Liberated energy does not get carried away anymore:
• ↑ T  ↑ MJ
•Fragmentation ceases because pressure increases and contraction stops 
protostars
(except for some rapidly rotating protostars – may split up  binary systems)
Accepted view – BUT
Fragmentation process conjectural
•effects of rotation, magnetic field and energy input not well understood
Start of contraction?
• compression of gas due to passage through spiral arm
• expanding HII regions or SN explosion
More insight is coming forth from IR observations
•condensing clouds have 100-1000K while IR radiation can escape even densest
clouds (see p 63 Fig. with maser sources BN)
In Summary: the 5 phases of interstellar gas
Hot Corona of Milky Way
•1956 Lyman Spitzer showed: MW had to be surrounded by a very hot gas
•20 years later: Copernicus satellite found evidence: Galactic coronal gas
(analogy to hot solar corona)
 emission lines of highly ionized O, N and C (O VI, N V, and C IV)
 requires temp. of 100000 – 1 000 000K, also from broadness of lines
•Galactic gas is evenly distributed through whole MW - extends several kpc from GP (up to
70 kpc according to recent FUSE measurements)
•Density is low 10-3 atoms/cm3 (recall mean density in GP is 1 atom/cm3)
 forms kind of background sea in which denser and cooler ‘clouds’ may form
•1980s: IUE satellite observed similar corona in LMC and in spiral arm of M100
 probably common and important form of matter in galaxies
•Source:
•
SN explosions most likely  form hot expanding bubbles that will permeate galaxies
•
stellar winds from hot stars
Planetary Nebulae (PN)
(PN have nothing to do with planets; nomenclature due to similarity in appearance)
Stellar evolution:
•bright regions of gas also occur around stars in ‘late’ stages of their evolution
•PN are gas shells expelled from the star (pulsations, stellar winds) leaving small hot
blue cores; whole outer atmosphere is being ejected into space.
• Estimated
number in MW:
50 000
• Observed 2000
• Sizes of a few
arcsec to
degrees
• Typical
physical sizes:
0.3 pc
AAO Student Fellowship Program
The Australian Astronomical Observatory provides opportunities for undergraduate students
to participate in research projects. Students will spend 10-12 weeks in the period Dec 2016
- Feb 2017 working at the Australian Astronomical Observatory in Sydney on research
projects under the supervision of AAO staff astronomers and engineers. Students will have
the opportunity to participate in a field trip to visit the telescopes at Siding Spring
Observatory.
Please encourage your undergraduate students to apply.
The deadline for applications is:
*** 31 August 2016 ***
Details are available here:
http://www.aao.gov.au/science/research/students/fellowships
The stipend is A$700 per week.
How to Apply
Applications are required to be sent by e-mail. Please send your application as a single
Word or PDF document attachment to the AAO Student Fellowship Coordinator, Prof.
Andrew Hopkins
([email protected]). The application should include the following:
• gas shell expanding at 20-30 km/s around the core of the original star (of 50000 –
100 000K)
• generally more symmetrical in shape and expand more rapidly (than HII regions)
(e.g. expansion rate is measurable between current images and photographs
from 50 yrs ago for Lyra – see image below)
• PNe will dissipate into ISM within at most 50 000 years (very quick compared to
lifetime of a star), while central stars cool into white dwarfs
• Expanding gas ionized from UV radiation of central star (like HII regions)
• causes atoms to become excited
• when electrons cascade back to lower energy levels, photons emitted in the
visible
• Often have bright emission lines from
forbidden transitions
• [OIII] at 4959Å and 5007Å – green
• H (6563Å; Balmer) and forbidden
ionized nitrogen ([NII] at 6583Å) –
red
PN Lyra (M57)
Other examples of PN:
Cats Eye (NGC 6543)
HST WFPV2 images (1995)
New H2 Ring image
(from rotational transitions of molecule)
From Spitzer IRAC observations
Helix (NGC 7293) by David Malin (AAT)
Closest PN: at a distance of d=213 pc; in the constellation Aquarius;
size: 16 arcmin (1/2 full moon)
N
•
Tadpole-like objects in
ring: "cometary knots” glowing heads and
gossamer tails resemble
comets (but see HST
image on next slide)
•
Each gaseous head at
least twice the size of our
solar system
•
each tail stretches about
1,000 AU
•
Gaseous knots
supposedly result from
collision between gases
•
The ejected hot gas
collides with the cooler
gas (ejected 10,000 years
before)
Hubble Space Telescope Helix: collision of two gases near a dying star
Red light: [NII] 6584Å emission; green: H, 6563Å; blue, [OIII] 5007Å
Evolution in clumps seems to correspond with age of the nebula (O’Dell 2002)
clumps in older nebulae are smaller and well-formed
clumps in younger nebulae larger and less sculpted
Supernova Remnants
•End point of evolution of massive stars
•Collapse of stellar core leads to violent ejection of outer layers - remain as
expanding gas cloud
•About 120 SNR known in MW
•Some optically visible as ring or irregular nebulae (see next 2 slides), but most are
only visible in radio (because radio not susceptible to extinction)
•In radio they are extended sources, similar to HII regions – BUT with polarized
radiation
• Emission of HII regions is thermal: free-free
emission of hot plasma  intensity grows
or remains constant with increasing
frequency, whereas SNR radiation falls off
linearly
• In SNR it is synchroton radiation from
electron moving in spiral orbits around the
magnetic field lines (not ionized by central
star)
 Continuous spectrum over all wavelengths
The Crab Nebula M1 (located in Taurus)
Explosion observed in 1054 by Chinese; Crab is still expanding at a rate of 1450 km/s!!
Luminosity: L = 8 x 104Lsun; mostly highly polarized synchroton radiation (thus indicating
presence of relativistic electrons spiraling around magn. field lines)
The continuing high luminosity and source of electrons remained a puzzle until the discovery of
a pulsar at its centre - all irregular SNR have pulsars at their center and are long-living
Crab Nebula looks blue due
to optical synchroton emission,
red filaments from H
NOTE:
Ring-like SNR have no pulsar:
all their energy comes from
SNR explosion
Vexp ~ 10’000-20’000 km/s
Forms shell as ejected material
starts to sweep up ISM
Expansion slows down and
shell cools and merges with
ISM after ~100’000yrs
The Vela SNR (image UK Schmidt)
Explosion about 11,000 years ago (may well be the first ever observed by human beings)
The optical photograph (below - left) is suggestive of shock waves:
• Debris from ejected SN material encounters material in ISM producing shock fronts
several AU wide.
• The shocks excite and ionize the ISM, causing the observed emission:
When the electrons recombine with these atoms, light in many different colors and
energy bands is produced
A spherical, expanding shock wave is also visible in X-rays (right); ROSAT PSPC finds hot gas
bubble of various million degrees (similar to Local Bubble)
After
Ring-like SN: SN 1987A
Exploded on 23 February 1987, in LMC, visible by eye
It shone as bright as 100 million suns for several
months
Direct observations of expansion of shell
HST images 1994:
Before
•
•
•
Temperature of the hot spots surges from a few thousand to a million degrees
Note fading of central star: about a factor of 1 million in intensity
Glowing debris of the central star heated by radioactive elements (principally titanium 44)
created in SN explosion and will continue to glow for many decades.
The progenitor of SN1987A did lose some mass  resulting in very unusual structure
Three rings:
•innermost ring has d = 0.42pc (a bit over 1 lyr); lies in a plane with SN at its center (slightly
inclined)
•
glow in the visible in OIII emission heated by UV radiation from SN reaching material
supposedly was ejected 20 000 years before the explosion
•2 larger rings: do not lie in plane containing the SN, but in front and behind; origin uncertain:
•
did star reside near NS or BH?  as this source wobbles it could create jets with
some kind of bipolar flow in 2 planes; possible source found at 0.1pc
•
large rings produce by hot fast stellar wind from blue supergiant progenitor, overtaking
the slower cooler wind emitted when star was still a red supergiant
•An aside: very exciting observations were based on its neutrinos  the first ever observed
neutrinos from another source than the Sun!
 neutrino burst was recorded over period of 12½ sec – 3hrs before arrival of first photons
12 events detected with Japan’s Kamiokande II Cerenkov detector
8 events in California, IMB Cerenkov detector.
Confirmation of basic SN core-collapse theory (“seeing” formation of NS out of collapsed iron
core)
 determination of upper limit of rest mass of electron neutrino:
me £ 16eV
Cosmic Rays
Composed mainly of bare H nuclei (protons), roughly 90%, and -nuclei (9%), rest of heavier
atomic nuclei and electrons
They occur throughout IS space; energy density of the same order as the radiation of stars
 important for ionization and heating of IS gas
Because they are charged  interact with magnetic field  change of propagation direction – no
info on source of origin
Energies: most have
: < 10 9 eV; number decreases with increasing energy
most energetic: 1020 eV (rare – 1 proton could lift a book by 1cm
largest particle accelerators reach ‘only’ 1012 eV)
The origin of low energy cosmic rays is difficult to disentangle as they are mixed with solar
cosmic rays
But gamma and radio observation allow mapping of distribution of cosmic rays in the MW:
- collision of cosmic ray protons with IS HI  pions which decay and form a
gamma ray background
- cosmic ray electrons emit synchroton radiation in IS magnetic field in radio
Both emissions show strong concentration to GP; with peaks that coincide with location of SNR,
such as Crab and Vela
Apparently SN form cosmic rays: - SN explosion  energetic particles
- if SNR has a pulsar  accelerate particles
- shock waves of expanding SNR shell  relativistic particles
The end of
ISM, dust and
pretty pictures
Revised PN Model
(from presentation of Sarah Eyerman on
The Nature and Origin of Molecular Knots in Planetary Nebulae)
CO
OI
O+
O2+
O3+
He2+
He+
He0
HI
Molecular
Clumps
shadow
H2
dust
H2
CO
H0 H+
star
light
How does this relate to what we see in other galaxies?
Strong correlations among total far-IR emission, H emission, and CO vs. total gas mass
The CO-FIR relation is illustrated by Fig. 7 of the review by Young and Scoville 1991 (ARA&A
29, 581, reproduced from the ADS), with the CO data transformed into H2 masses:
But Kennicutt (The Interstellar Medium in
Galaxies, 1990):
the same CO-FIR relation is also followed
by a burning cigar, his Jeep, the
Yellowstone forest fire, and the observable
Universe .
This may be another manifestation of the
well-known astrophysical principle that big
galaxies are big and little ones are little, so
that differences among scale-linked
properties are second-order effects.