![]() | Annu. Rev. Astron. Astrophys. 2013. 51:
207-268 Copyright © 2013 by Annual Reviews. All rights reserved |
Molecular hydrogen, H2, is the most abundant molecule in the
universe. With the possible exception of the very first generations
of stars, star formation is fueled by molecular gas. Consequently,
H2 plays a central role in the evolution of galaxies and stellar
systems (see the recent review by
Kennicutt
& Evans 2012).
Unfortunately for astronomers interested in the
study of the molecular interstellar medium (ISM), cold H2 is not
directly observable in emission. H2 is a diatomic molecule with
identical nucleii and therefore possesses no permanent dipole moment
and no corresponding dipolar rotational transitions. The lowest energy
transitions of H2 are its purely rotational quadrupole
transitions in the far infrared at
= 28.22
µm and shorter
wavelengths. These are weak owing to their long spontaneous decay
lifetimes
decay ~
100 years. More importantly, the two
lowest para and ortho transitions have upper level energies E /
k
510 K and
1015 K above ground
(Dabrowski
1984).
They are thus only excited in gas with T
100 K. The lowest
vibrational transition of
H2 is even more difficult to excite, with a wavelength
= 2.22 µm
and a corresponding energy E / k = 6471 K.
Thus the cold molecular hydrogen that makes up most of the molecular
ISM in galaxies is, for all practical purposes, invisible
in emission.
Fortunately, molecular gas is not pure H2. Helium, being
monoatomic, suffers from similar observability problems in cold
clouds, but the molecular ISM also contains heavier elements at the
level of a few × 10-4 per H nucleon. The most abundant of
these are oxygen and carbon, which combine to form CO under the
conditions prevalent in molecular clouds. CO has a weak permanent
dipole moment (µ
0.11 D = 0.11
× 10-18 esu cm) and
a ground rotational transition with a low excitation energy
h
/ k
5.53 K. With this
low energy and critical density
(further reduced by radiative trapping due to its high optical
depth), CO is easily excited even in cold molecular clouds. At a
wavelength of 2.6 mm, the J = 1 → 0 transition of CO falls
in a fairly transparent atmospheric window. It has thus become the workhorse
tracer of the bulk distribution of H2 in our Galaxy and beyond.
As a consequence, astronomers frequently employ CO emission to measure molecular gas masses. The standard methodology posits a simple relationship between the observed CO intensity and the column density of molecular gas, such that
![]() |
(1) |
where the column density, N(H2), is in cm-2 and the integrated line intensity, W(CO) 1, is in traditional radio astronomy observational units of K km s-1. A corollary of this relation arises from integrating over the emitting area and correcting by the mass contribution of heavier elements mixed in with the molecular gas,
![]() |
(2) |
Here Mmol has units of
M and
LCO is usually
expressed in K km s-1 pc2. LCO
relates to the observed integrated flux
density in galaxies via LCO = 2453
SCO
v
DL2 / (1 + z), where
SCO
v is the
integrated line flux density,
in Jy km s-1, DL is the luminosity distance
to the source in Mpc, and z is the redshift (e.g.,
Solomon &
Vanden Bout 2005,
use Eq. 7 to convert between W(CO) and SCO
v). Thus
CO is simply
a mass-to-light ratio. The correction for the contribution of heavy
elements by mass reflects chiefly helium and amounts to a
36% correction
based on cosmological abundances.
Both XCO and
CO are
referred to as the "CO-to-H2 conversion
factor." For XCO = 2 × 1020
cm-2(K km s-1)-1 the corresponding
CO is 4.3
M
(K km s-1 pc2)-1. To translate integrated
flux density directly to molecular mass, Equation 2 can be
written as
![]() |
(3) |
For convenience we define
![]() |
(4) |
We discuss the theoretical underpinnings of these equations in Section 2.
Note that the emission from CO J = 1 → 0 is found to be consistently optically thick except along very low column density lines-of-sight, as indicated by ratios of 12CO to 13CO intensities much lower than the isotopic ratio. The reason for this is simple to illustrate. The optical depth of a CO rotational transition is
![]() |
(5) |
where J and NJ are the rotational quantum
number and the
column density in the upper level of the J → J - 1
transition, is the frequency,
Tex is the excitation
temperature (in general a function of J, and restricted to be
between the gas kinetic temperature and that of the Cosmic Microwave
Background),
v
is the velocity width, µ is the dipole
moment, gJ = 2J + 1 is the statistical weight
of level J, and h and
k are the Planck and Boltzmann constants respectively. Under
typical conditions at the molecular boundary,
1 for the
J = 1 → 0 transition requires N(H2)
2-3 ×
1020 cm-2 for a Galactic carbon gas-phase abundance
AC ~ 1.6 × 10-4
(Sofia et
al. 2004).
At the outer edge
of a cloud the carbon is mainly C+, which then recombines with
electrons to form neutral C (Fig. 1). Carbon
is converted to CO by a series of reactions initiated by the
cosmic-ray ionization of H or H2 (e.g.,
van Dishoeck
& Black 1988)
and becomes the dominant carrier of
carbon at AV ~ 1-2. The CO
J = 1 → 0 line turns optically thick very quickly after CO
becomes a significant carbon reservoir, over a region of thickness
AV
~ 0.2-0.3 for a typical Galactic dust-to-gas ratio (c.f., Eq. 21).
![]() |
Figure 1. Calculated cloud structure as a
function of optical depth into the cloud. Top panel shows the
fractional abundance of HI, H2, C+, C, and
CO. Middle panel shows their integrated column densities from the
cloud edge. Bottom
panel shows the emergent line intensity in units of K km s-1
for [CII] 158 µm, [CI] 609 µm, and CO J
= 1 → 0. The grey vertical bar shows where CO J = 1 →
0 becomes optically thick. At the outer edge of the cloud gas is mainly
HI. H2 forms at AV ~ 0.5 while the carbon
is mainly C+. The C+ is converted to C at
AV ~ 1 and CO dominates at AV
|
Equations 1 and 2 represent highly idealized,
simplified relations where all the effects of environment, geometry,
excitation, and dynamics are subsumed into the XCO or
CO coefficients.
A particular example is the effect that spatial
scales have on the CO-to-H2 conversion factor. Indeed, for
the reasons discussed in the previous paragraph, XCO
along a line-of-sight through a dense molecular cloud where
AV
10 is
not expected to be the same as
XCO along a diffuse line-of-sight sampling mostly
material where AV < 1 (see, for example,
Pineda et
al. 2010,
Liszt & Pety
2012).
Thus on small spatial scales we expect to see a large variability in
the CO-to-H2 conversion factor. This variability will
average out on the large spatial scales, to a typical value
corresponding to the dominant environment. Because of the large
optical depth of the CO
J = 1 → 0 transition the velocity dispersion giving rise to
the width of
the CO line will also play an important role on W(CO), and
indirectly on XCO. Indeed, there is not one value of
XCO that is
correct and applicable to each and every situation, although there are
values with reasonable uncertainties that are applicable over large
galactic scales.
The plan of this work is as follows: in the remainder of this section we provide a brief historical introduction. In Section 2 we present the theoretical background to the CO-to-H2 conversion factor. In Section 4 we review the methodology and measurements of XCO in the Milky Way, the best understood environment. We characterize the range of values found and the underlying physics for each measurement technique. In Section 5 we review the literature on XCO determinations in "normal" star-forming galaxies and discuss the techniques available to estimate XCO in extragalactic systems. In Section 6 we consider the effect of metallicity, a key local physical parameter. In Section 7 we review the measurements and the physical mechanisms affecting the value of the CO-to-H2 conversion factor in the starburst environments of luminous and ultraluminous galaxies. In Section 8 we consider the explicit case of XCO in high redshift systems, where a much more restricted range of observations exist. In Section 3 we discuss the results of recent calculations of molecular clouds including the effects of turbulence and chemistry. Finally, in Section 9 we will offer some recommendations and caveats as to the best values of XCO to use in different environments, as well as some suggestions about open avenues of research on the topic.
1.1. Brief Historical Perspective
Carbon monoxide was one of the first interstellar medium molecules
observed at millimeter wavelengths.
Wilson, Jefferts
& Penzias (1970)
reported the discovery of intense CO emission from the Orion nebula
using the 36 foot NRAO antenna at Kitt Peak, Arizona. Surveys of
molecular clouds in the Galaxy (e.g.,
Solomon et
al. 1972,
Wilson et
al. 1974,
Scoville &
Solomon 1975,
Burton et
al. 1975)
established molecular gas to be widespread in the inner Milky Way with
a distribution that resembles giant HII regions more closely than that
of atomic hydrogen gas. The combination of CO and
-ray
observations demonstrated that H2 dominates over HI by mass in
the inner Galaxy
(Stecker et
al. 1975).
By the end of the following decade,
these studies extended to complete the mapping of the Galactic plane
(Dame et al. 1987).
The first extragalactic detections of CO occurred in parallel with these early Galactic surveys (Rickard et al. 1975, Solomon & de Zafra 1975). They found CO to be particularly bright in galaxies with nuclear activity such as M 82 and NGC 253. The number of extragalactic CO observations grew rapidly to include several hundred galaxies over the next two decades (Young & Scoville 1991, Young et al. 1995), and CO emission was employed to determine galaxy molecular masses (Young & Scoville 1982). By the late 1980s, the first millimeter interferometers spatially resolved molecular clouds in other galaxies (Vogel, Boulanger & Ball 1987, Wilson et al. 1988). Such observations remain challenging, though powerful new interferometric facilities such as the Atacama Large Millimeter Array (ALMA) will change that.
The first detection of CO at cosmological redshifts targeted ultraluminous infrared sources and revealed very large reservoirs of highly excited molecular gas (Brown & Vanden Bout 1991, Solomon, Downes & Radford 1992). Because of the deep integrations required, the number of high redshift CO detections grew slowly at first (Solomon & Vanden Bout 2005), but this field is now developing rapidly driven by recent improvements in telescope sensitivity (see the review by Carilli & Walter in this issue). An increased appreciation of the roles of gas and star formation in the field of galaxy evolution and the according need to determine accurate gas masses provides one of the motivations for this review.
Under average molecular cloud conditions, CO molecules are excited
through a combination of collisions with H2 and radiative
trapping. They de-excite through spontaneous emission and collisions,
except at very high densities where collisions are extremely
frequent. Neglecting the effect of radiative trapping, radiative and
collisional de-excitation will balance for a critical density
ncr,J
AJ /
J(Tkin)
(neglecting the effects of stimulated emission), where
Tkin is the kinetic gas temperature.
Thus, for n≫ ncr,J and excitation
temperatures Tex,J
≫ EJ / k
5.53 J
(J + 1) / 2 K the upper level of the J →
J - 1 transition will be populated and the molecule will emit
brightly. In these expressions AJ is the Einstein
coefficient for spontaneous emission (only transitions with
|
J| = 1 are
allowed), AJ = 64
4
J3
µ2 J / (3 h c3
gJ) (A1
7.11 ×
10-8 s-1). The parameter
J(T)
is the corresponding collisional coefficient
(the sum of all collisional rate coefficients for transitions with
upper level J), which is a weak function of temperature. For CO,
1 ~ 3.26 × 10-11
cm3
s-1 for collisions with H2 at Tkin
30 K
(Yang et al. 2010).
Tex,J refers to the excitation
temperature, defined as the temperature needed to recover the relative
populations of the J and J-1 levels from the Boltzmann
distribution. In general, Tex, J will be different for
different transitions.
The critical density for the CO J = 1 → 0 transition is
n(H2)cr,1 ~ 2200 cm-3. Higher
transitions require rapidly
increasing densities and temperatures to be excited, as
ncr,J
J3 and EJ
J2. The high optical depth of the CO emission
relaxes these density requirements, as radiative trapping
reduces the effective density required for excitation by a factor
~ 1 /
J (the
precise factor corresponds to an escape probability and is dependent on
geometry).
The Rayleigh-Jeans brightness temperature, TJ, measured by a radio telescope for the J → J - 1 transition will be
![]() |
(6) |
The final term accounts for the effect of the Cosmic
Microwave Background at the redshift, z, of interest. Note that
frequently the Rayleigh-Jeans brightness temperature is referred to as
the radiation temperature. Observations with single-dish telescopes
usually yield antenna temperatures corrected by atmospheric
attenuation (TA*), or main beam
temperatures (TMB), such that TMB =
MB
TA* where
MB
is the main beam
efficiency of the telescope at the frequency of the observation (see
Kutner &
Ulich 1981
for further discussion). The Rayleigh-Jeans
brightness temperature TJ is identical to
TMB for compact
sources, while extended sources may couple to the antenna with a
slightly different efficiency.
Extragalactic results, and measurements with interferometers, are frequently reported as flux densities rather than brightness temperatures. The relations between flux density (in Jy) and Rayleigh-Jeans brightness temperature (in K) in general, and for CO lines, are
![]() |
(7) |
where is the
half-maximum at full width of the telescope beam (in arcsec), and
is the wavelength of
a transition (in mm).
The Rayleigh-Jeans brightness temperature, TJ, excitation
temperature, Tex,J, and kinetic temperature,
Tkin, are
distinct but related. The 12CO transitions usually have
J ≫ 1,
making the Rayleigh-Jeans brightness temperature a probe
of the excitation temperature, TJ ~
Tex,J for Tex,J ≫
5.53 J K. In general Tex,J can be shown to be
in the range
Tcmb ≤
Tex,J ≤ Tkin. At densities
much higher than
ncr,J the population of the levels J and lower
will approach a
Boltzmann distribution, and become "thermalized" at the gas kinetic
temperature, Tex,J
Tkin. The
corresponding Rayleigh-Jeans brightness can be computed using
Eq. 6. When Tex,J < Tkin, usually
Tex,J / Tex,J-1 < 1 for lines
arising in the same parcel of
gas, and the excitation of the J level is "subthermal" (note that
this is not equivalent to TJ
/ TJ-1 < 1, as is sometimes used in the literature).
1 Henceforth we refer to the most common 12C16O isotopologue as simply CO, and unless otherwise noted to the ground rotational transition J = 1 → 0. Back.