D. Reimers1 - R. Baade1 - R. Quast1 - S. A. Levshakov2
1 - Hamburger Sternwarte, Universität Hamburg,
Gojenbergsweg 112, 21029 Hamburg, Germany
2 -
Department of Theoretical Astrophysics,
Ioffe Physico-Technical Institute,
194021 St. Petersburg, Russia
Received 30 June 2003 / Accepted 15 August 2003
Abstract
A new molecular hydrogen cloud is found in the sub-damped
Ly
absorber
[
(H I) =
]
at the redshift
= 1.15 toward the bright quasar HE 0515-4414 (
= 1.71).
More than 30 absorption features in the Lyman band system of H2are identified in the UV spectrum of this quasar
obtained with the Space Telescope Imaging Spectrograph (STIS)
aboard the Hubble Space Telescope.
The H2-bearing cloud shows a total H2 column density
N(H2)
cm-2 and
a fractional molecular abundance
derived from the H2 lines arising from the J = 0-5 rotational levels
of the ground electronic vibrational state.
The estimated rate of photodissociation at the cloud edge
s-1 is much higher
than the mean Galactic disk value,
s-1.
This may indicate an enhanced star-formation activity in the
z = 1.15 system as compared with molecular clouds
at
where
.
We also find a tentative evidence that the formation rate
coefficient of H2 upon grain surfaces at z = 1.15is a factor of 10 larger
than a canonical Milky Way value,
cm3 s-1.
The relative dust-to-gas ratio estimated from the [Cr/Zn] ratio
is equal to
(in units of the mean Galactic disk value),
which is in good agreement with a high molecular fraction in this system.
The estimated line-of-sight size of
pc may imply that the H2is confined within small and dense filaments embedded in a more rarefied gas giving
rise to the z = 1.15 sub-damped Ly
absorber.
Key words: cosmology: observations - quasars: absorption lines - quasars: individual: HE 0515-4414
The most abundant interstellar molecule in the universe, H2, is currently observed not only
in the Milky Way disk (e.g., Rachford et al. 2002) and halo (e.g., Richter et al. 2003a),
but also in the Magellanic Clouds (e.g., Tumlinson et al. 2002)
and in more distant regions
of the universe such as intervening damped Ly
absorbers (DLAs)
seen in spectra of background quasars (QSOs).
The DLAs are the systems with neutral hydrogen column densities
N(H I)
cm-2. They are believed to originate
in protogalactic disks (Wolfe et al. 1995). The systems with lower hydrogen column densities,
1019 cm-2
(H I)
cm-2, are formally called sub-DLAs.
The sub-DLAs may also be related to intervening galaxies.
At the moment there are known 9 molecular hydrogen systems detected in DLAs and sub-DLAs in the redshift range
from
= 1.96 to 3.39 (see Table 1).
In this paper, we present results from the analysis of
a new 10th H2 system detected at
= 1.15 in the sub-DLA toward
the bright quasar HE 0515-4414.
This is the first detection of H2 with the STIS/HST at an intermediate redshift - a cosmological epoch when an enhanced star formation rate (SFR)
is observed in young galaxies.
The SFR shows a peak at
over the redshift interval
(see, e.g., Hippelein et al. 2003 and references therein).
Molecular hydrogen, being an important coolant for gravitational collapse of gas clouds
at
K, is known to play a central role in star formation processes, and thus
one may expect that the SFR and the fractional abundance of H2 are correlated.
Studying H2-bearing cosmological clouds leads to better understanding of the physical
environments out of which first stellar populations were formed.
Table 1:
Molecular hydrogen abundances
and dust contents
in cosmological H2-bearing clouds.
Spectral data of the quasar HE 0515-4414 (
= 1.71, V = 14.9; Reimers et al. 1998)
in the UV range were obtained
with the HST/STIS (Reimers et al. 2001).
The medium resolution NUV echelle mode (E230M) and a
aperture provides a resolution power of
(FWHM
km s-1).
The overall exposure time was 31 500 s.
The spectrum covers the range between 2279 Å and 3080 Å
where the signal-to-noise ratio (S/N) per resolution
element varies from S/N
to
5.
The data reduction was performed by the HST pipeline completed
by an additional inter-order background correction and by coadding the
separate sub-exposures.
The spectral portion where the H2 lines occur suffers from a poor S/N ratio (20).
An additional problem arises from the limited wavelength
accuracy. The MAMA detectors produce an absolute wavelength definition
between 0.5-1.0 pixel (
limit as given by Brown et al. 2002).
For our data 1 pixel corresponds to 0.038 Å. The spectral
overlap of successive echelle orders allows to examine the wavelength
errors from order to order. Using well-defined line profiles we find
relative wavelength shifts of 0.02-0.05 Å (see Fig. 1 for an
example).
Additional echelle spectra of HE 0515-4414 were obtained during
ten nights between October 7, 2000 and January 2, 2001 using
the UV-Visual Echelle Spectrograph (UVES) installed
at the VLT/Kueyen telescope. These observations were carried
out under good seeing conditions
(0.47-0.70 arcsec) and a slit width of 0.8 arcsec giving
the spectral resolution of
(FWHM
km s-1).
The VLT/UVES data have a very high S/N ratio
(
50-100 per resolution element) which allows
us to detect weak absorption features.
The high resolution VLT/UVES data reveal two narrow components
in the fine-structure C I lines associated with the sub-DLA at
= 1.15 (de la Varga et al. 2000). The stronger
component at
= 1.15079 is separated from the weaker one
by
km s-1, and shows 2.8 times higher column density
(Quast et al. 2002, hereafter QBR).
Exactly at the redshift of C I lines we identified more
than 30 absorption features in the Lyman band system of molecular hydrogen H2(see Fig. 2).
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Figure 1:
Typical example for a relative wavelength shift between
successive orders. The H2 L2-0 P(3) line shows a difference of
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Figure 2:
Continuum-normalized STIS spectra (histograms) of the quasar HE 0515-4414
(individual echelle orders marked by a- e)
and over-plotted synthetic H2 profiles (smooth lines) arising from the rotational levels
J = 0 to J = 5 of the lowest vibrational level v = 0 in the ground electronic state
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In this section we describe the measurements
of the neutral hydrogen column density, metal and dust
content and the H2 abundances in the
= 1.15 sub-DLA.
These values are well known to be physically related.
The formation and maintenance of diffuse H2 in the Milky
Way clouds is tightly correlated to the amount of interstellar
dust grains, which provide the most efficient H2 formation
on their surfaces (see, e.g., Pirronello et al. 2000 and
references therein).
In order to estimate the column density of atomic hydrogen contained
in the sub-DLA, particular care has to be taken. Since the Doppler core of the Ly line is completely saturated, only the Lorentzian part
gives information about the line profile (Fig. 3). Moreover, the
Lorentzian part is less pronounced than for typical DLAs and hence less
distinguishable from the quasar continuum. Therefore, we
simultaneously optimised the continuum
and fitted the spectral features using
standard Voigt profile fitting technique. The
continuum is modelled as a linear function, and the
Voigt function is calculated using the pseudo-Voigt approximation
(Thompson et al. 1987).
Our optimized model (Fig. 3) reveals some additional
absorption in the blue part of the damped Ly
line at
km s-1.
This additional absorption is H I Ly
which
is seen also in metal lines. The whole sub-DLA system is spread over 700 km s-1 (Quast et al. 2003).
This line together with
other narrow absorption lines seen in the wings of the damped Ly
were included in the Voigt
fitting.
The derived column density of atomic hydrogen
in the sub-DLA is (H I)
,
where we estimated the standard deviation by varying the column density
until the resulting model profile is apparently inconsistent with the
observed data (Fig. 3).
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Figure 3:
Part of the H I Ly![]() ![]() ![]() ![]() |
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Figure 4:
Parts of the UVES observations showing absorption arising
from the ions Cr II and Zn II (histograms).
The solid and dashed lines represent our optimised model and its
deconvolution, respectively. Mg I
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To measure the metal abundance in the main sub-component of the
= 1.15 system
we used the Zn II
lines (Fig. 4).
The presence of dust grains in DLAs is usually estimated from the abundance ratio
[Cr/Zn]
assuming that Zn is undepleted (Pettini et al. 1994).
In our high S/N spectrum, only a weak Cr II
line was detected
at
km s-1 (Fig. 4). Other Cr II lines (
)
are too weak to be visible. Their oscillator strengths scale as
f2056:f2062:f2066 = 1:0.74:0.50 (Bergeson & Lawler 1993).
The column densities for Zn II and Cr II
were calculated by Quast et al. (2003):
(Cr II) =
and
(Zn II) =
.
We used these values to estimate the dust-to-gas ratio in Sect. 4.2.
Molecular hydrogen at
= 1.15079 is detected in the J = 0 up to J = 5 rotational levels.
At a spectral resolution of
10 km s-1 it is not possible
to resolve the internal structure in the H2 lines
observed in the C I absorption (see Fig. 1 in QBR).
The H2-bearing gas may also be distributed over a wider velocity range
as compared with C I which is easily ionised by UV photons in optically
thin zones.
However, for a good approximation one can assume that H2traces, in general, the volume distribution of C I since such correlation is indeed
observed in the Milky Way (e.g., Federman et al. 1980).
Therefore, in our H2 analysis we used a two-component model based on
the observations of C I by QBR. We note that the C I data were obtained with
higher spectral resolution (FWHM
km s-1) and considerably higher signal-to-noise
ratio (up to S/N
for the parts of the spectrum with C I lines).
Panels a-e in Fig. 2 present echelle orders of the STIS spectra
of HE 0515-4414 (histograms) in the wavelength regions of the H2 Lyman 0-0 to 4-0 bands,
together with a two-component Voigt profile fit of the data (smooth lines).
It is seen that some of the identified H2 transitions are
contaminated by the Ly
forest or blended with metals from different intervening
systems.
This hampers significantly the measurements of
accurate equivalent widths and their analysis through the curve of growth.
Moreover, the noise level (shown by the dashed line) is rather high for the available
STIS data and this may explain why some of H2 features are inconsistent with others.
For instance, the observational profile of L4P4
in panel a differs from those of L3R4 (a), L3P4 (b), L2R4 and L2P4 (c),
and L1R4 (d). Relative strengths of the L0R0 and L0R1 lines from different
echelle orders (panels d and e) are not consistent (L0R0 is partly blended
with Fe II
at z = 0).
The apparent depths of the close pairs L0P2 + L0R3 (e) and L0P1 + L0R2 (e and d),
as well as the single lines L1R2, L2R4 (c) are deeper than those
calculated from the simultaneous fit to all H2 lines.
Under these circumstances a standard -square fitting cannot
provide a statistically valuable measure of
goodness-of-fit. To estimate model parameters we required that the calculated
spectra were within 1
uncertainty range for the majority
of the unblended H2 profiles
or their unblended portions which match the data.
We tried to optimize a set of the H2 column densities for the two-component model with the components located at the measured redshifts of the C I lines z1 = 1.15079 and z2 = 1.15085. The broadening b-parameters were fixed at b1 = 2.0 km s-1, b2 = 3.5 km s-1 (as deduced from the C I lines by QBR), but a column density ratio between the sub-components, N2/N1, was a free parameter ranging from 0.03 to 0.36 (the latter corresponds to the C I column density ratio found in QBR).
For a given H2 component, the same b-parameter was used independently of the rotational level. The column density in each J level was derived from several calculations of the Voigt profiles with a fixed value of N2/N1 and different N(J) which match the observational spectra. The limiting values of N(J) (an adjustable minimum and maximum) were chosen to estimate the uncertainty interval for column densities.
All identified transitions from J = 0 and 1 are optically thick, but
the apparent central intensities of the L0R0 and L0R1 lines (d) are not zero
(we consider the L0R0 line in panel e as corrupted by a bad merging of different
spectra). These lines restrict N(0) and N(1) by, respectively,
and
cm-2 at
N2/N1 = 0.03.
On the other hand, we observe neutral carbon which is usually shielded
in molecular clouds from the background ionising UV radiation by the H2 absorption
arising from the J = 0 and 1 levels. An essential shielding in H2 lines
occurs when N(H2)
cm-2.
This gives us a hint at a possible range of N(J=0,1).
For the lower value of
N2/N1 = 0.03, the contribution from the second H2 component is negligible, but the
synthetic profiles are systematically narrower as compared with the data.
The presence of the second component is, therefore, important. On the other hand, the maximum
value of
N2/N1 = 0.36 provides too wide synthetic profiles even with
N(0) = N(1) = 1016 cm-2. We found that with
an optimal set of the H2 column densities may be deduced. An example of such solution
is shown in Fig. 2. The obtained results are given in Table 2.
We investigate now the physical conditions in the
= 1.15 H2-bearing cloud
by considering the processes and parameters that balance the formation and dissociation
of molecular hydrogen. Since our observations show a relatively high metallicity in
this sub-DLA, [Zn/H] =
[i.e.,
-
],
and the dust content is approximately similar to the mean value for the cold gas in the
Galactic disk (
), we consider catalytic reactions on the surfaces of
dust grains (Hollenbach & Salpeter 1971) as the dominant H2 formation process,
whereas ion-molecular gas phase reactions (Black 1978) are less efficient.
The measured column densities can be used to estimate the
kinetic temperature,
,
the gas density,
,
the photodissociation rate, I, and the rate of molecular formation on grains, R.
However, in view of the large uncertainties
in the column densities N(0) and N(1),
we can only provide an order-of-magnitude estimate for these parameters.
Another obstacle in the H2 analysis is that the balance equation is related to the space
densities of H I and H2. In case of homogeneous clouds one can assume that
n(H I)/n(H2) (H I)/N(H2).
However, this assumption may not be correct
for DLAs where multiphase structures and complex profiles
are usually observed.
Observations show that with each step in increasing spectral resolution the profiles
break up into subcomponents down to the new resolution limit.
Table 2:
The H2 column densities for different rotational levels
from the
= 1.15 sub-DLA toward HE 0515-4414.
The sub-DLA at
= 1.15 reveals,
for example, transitions from neutrals
and low ions
(as C I, O I, C II, Si II)
to highly ionized ions
(as C IV, Si IV) spread over
km s-1 (Quast et al. 2003) which implies that the neutral H2-bearing
cloud(s) is embedded in a lower density, higher temperature gas.
Neutral hydrogen H I
can be spread over all gas phases that contain neutral gas with and without molecules.
The H2 on the other hand may have a very inhomogeneous
distribution in DLAs and concentrate in small clumps (Hirashita et al.
2003). Thus, only a fraction of the total H I may be relevant to
the formation of H2.
In our Galaxy, for instance, "tiny-scale atomic structures'' (TSAS)
and "small-area molecular structures'' (SAMS)
in the ISM are observed (e.g., Lauroesch & Meyer 1999;
Heithausen 2002). They show very high densities (
cm-3)
and very small sizes (
a few AU).
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Figure 5:
H2 rotational excitation in the
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To take this uncertainty into account, a scaling factor
for the
H I column density can be introduced. Following Richter et al. (2003b), who observed
similar complex structures in the Milky Way halo molecular clouds, we define
The kinetic temperature of the gas is usually estimated through
the excitation temperature T01 describing the
relative populations of the J=0 and J=1 levels.
This temperature is proportional to the negative inverse of the slope of the excitation diagram
drawn through the points of the respective J levels in a plot
versus E(J) shown in Fig. 5. Here, E(J) is the excitation energy
of the rotational level J relative to J=0, and g(J) is its statistical weight.
Figure 5 shows that the value of T01
is rather uncertain in our case because of large errors in
N(0) and N(1). Its mean value
K corresponds
to the excitation diagram shown by
the dotted line, and its upper limit is about 270 K, which represents,
probably, an upper limit for
of the gas in the main sub-component of the
= 1.15 system.
For levels with J=2,3,4, and 5 the accuracy of the column densities is higher and we find
K. The difference between
and
is not significant and the points in Fig. 5 can be fitted, in principle,
to a single excitation diagram.
But the previous analysis of the fine-structure level populations of C I,
where the most probable value for
of 240 K was found (QBR), indicates that these
two temperatures may not be equal.
This is also in line with results on the H2 study in the Milky Way which revealed
that single excitation diagrams fit usually only
optically thin lines with N(H2) < 1015 cm-2 (Spitzer & Cochran 1973).
For higher column densities, there is "bifurcation to two temperatures,
depending on the J levels'' (Jenkins & Peimbert 1997).
According to our calculations presented in Sect. 3.1,
the total H I column density in the main sub-component is
N(H I)
cm-2.
With N(H2)
cm-2,
the ratio of H nuclei in molecules to the
total H nuclei is
Listed in Table 1 are the molecular hydrogen fractions in all known H2 systems.
The
values were derived in the standard way assuming the scaling factor
.
This may imply that
the listed molecular hydrogen fractional abundances are systematically underestimated.
In Fig. 6 we compare these abundances with the dust-to-gas ratios, ,
estimated from
(Vladilo 1998):
Figure 6 demonstrates an apparent correlation between
and
in the range
which supports
the assumption that molecular hydrogen
abundances in quasar absorbers are governed by the dust content
similar to that observed in the Galaxy.
This conclusion rises the question: why is the H2 detection
in QSO absorbers in this case so rare (lower than 30% according to Ledoux et al. 2003)?
Following Hirashita et al. (2003), we
suppose that
a relative paucity of H2 observations in DLAs may be caused by a bias against finding H2in dense molecular clumps that have
a small angular extent and thus a small volume filling factor.
DLAs are mainly associated with diffuse clouds that have large volume filling factor and low
molecular fractions,
but they may also contain a small size dense filaments like the above mentioned
TSAS or SAMS.
Besides,
low metallicity of the QSO absorbers can also significantly
suppress H2 formation (Liszt 2002).
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Figure 6:
Relation between H2 fractional abundance
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One point in Fig. 6 (Q 0551-366) shows an unrealistic high
dust-to-gas ratio, about 2 times the
Galactic value. This large value may be, probably, explained by
systematic errors in the measurement of the Zn II column density.
For instance, the relative
abundance of Si, [Si/H]
,
differs significantly
from [Zn/H]
according to Ledoux et al. (2002).
The fraction of these elements in dust in the Milky Way is approximately identical,
and
(V02).
At [Fe/H]
,
Ledoux et al. measured
[Si/Fe]
which is in line with other observations (see Fig. 1 in V02).
This means that [Zn/H] is
most likely
overestimated in the
= 1.962 system.
In equilibrium between formation on grains with rate coefficient R (cm3 s-1)
and photodissociation with rate I (s-1),
we may write that (Jura 1975b)
To estimate the formation rate of H2 upon grain surfaces
,
we use approximation described by Jura
(1975b). It assumes that (i) the levels J = 4 and J = 5 are populated by direct
formation pumping and by UV pumping from J = 0 and J = 1, (ii) the self-shielding
in the levels J = 0 and J = 1 is about the same, (iii) the upper levels
J = 4 and J = 5 are depopulated by spontaneous emission (which is valid if
cm-3). We do not consider additional rotational excitation of H2caused by a shock because restrictions on the gas density
(
cm-3) and kinetic temperature (
K)
set by the observations of C I, C I
,
and C I
(QBR) show that collisional excitation of the levels J = 4 and J = 5 is not
significant.
Using the cascade redistribution probabilities
p4,0 = 0.26 and
p5,1 = 0.12,
calculated by Jura (1975a), and assuming
K, we can re-write Eqs. (3a)
and (3b) from Jura (1975b) in the form
To estimate the photodissociation rate at the cloud surface I0, the shielding
effect is to be taken into account. The shielding factor, S, depends on line overlap,
self-shielding of H2, and continuum absorption. Lee et al. (1996) showed that these
various factors can be well represented by the H2 column density. They calculated the
values of S as a function of N(H2) for a turbulent velocity of 3 km s-1,
which suits well for our case.
From their Table 10 we find
for
N(H2) =
cm-2, respectively.
This gives us a
rough estimate of
(H I)/SN(H2)
s-1 (with the uncertainty of about 120%), or
s-1.
The result obtained should be considered, however, as an upper limit on
since in our estimations we assumed that the H2 is one singe gas cloud.
If, in reality, the H2 is inhomogeneously distributed among several cloudlets,
the value of
should be lower.
In the Milky Way, the mean value for
s-1
(e.g., Richter et al. 2003b)
and, thus, we may conclude that the H2in the
= 1.15 sub-DLA is
probably exposed to a radiation field with the intensity much higher than
the mean Galactic value.
For comparison, molecular clouds in the LMC and SMC also reveal
10-100 times more intense UV radiation field than the Galactic one (Browning et al. 2003).
High rotational excitation of the H2 observed in the LMC/SMC gas and
at
= 1.15 is compared with Galactic data in Fig. 7.
We note that the upper limit on the local UV field at
= 1.15 of
about 100 times the Galactic value was independently determined
from the analysis of the C I fine-structure lines by QBR.
We can also estimate the photodissociation rate
produced by the intergalactic UV background (UVB) field
on the surface of a cloud (Hirashita et al. 2003):
Star-formation activity is not, however, intense
in the H2-bearing clouds at higher redshift. For example,
the UV radiation fields in the
= 3.025 H2 absorber toward Q 0347-3819
and in the Galactic ISM are very much alike (Levshakov et al. 2002). Other examples
may be found in Ledoux et al. (2003).
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Figure 7:
Total column density of H2 vs. excitation ratios
N(4)/N(2) and N(5)/N(3), for LMC, SMC, Milky Way (data
are taken from Browning et al. 2003), and sub-DLA at
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We now consider constraints on the volumetric gas density
and the H2formation rate coefficient R stemming from the foregoing estimation of
s-1.
A useful reference point is provided by the J = 2 level.
The population of this level is more directly affected by collisional processes, since
it has a longer radiative lifetime as compared with J=3 and other levels.
The critical density,
,
at which the probability of collisional and
radiative de-excitation of J=2 are equal is 200 cm-3, if
= 100 K, and
cm-3, if
= 240 K (the collisional de-excitation
rate coefficients qjj' are taken from Forrey et al. 1997).
If
,
collisional de-excitation becomes important.
To estimate ,
the grain formation rate of H2 should be known or vice versa.
R is a complex function of the gas and dust temperature and other poorly known parameters
(Hollenbach & McKee 1979).
According to their calculations,
cm3 s-1.
Using this value, we find
cm-3,
if
.
However,
for such large density, the J=0 and J=2 levels as well as the J=1 and J=3levels should be in thermal equilibrium, which we do not observe. Moreover, the upper limit
on the gas density set from the excitation of C I is 110 cm-3 (QBR).
If we tentatively adopt
cm-3, then
.
This value of R is larger than predicted in theoretical calculations,
but it is conceivable
that R may vary in space since
we observe considerable variations in the UV extinction among
different clouds. A low grain formation rate of H2 (
)
was, for instance, recently estimated in the LMC and SMC by Browning et al. (2003).
Although our high value of R is consistent with the best determinations of upper
limits to R toward
Peg,
Sco,
Sco etc. (Jura 1974),
we cannot without further observations
conclude that this result is certain.
To test whether or not the grain formation rate coefficient R
exceeds the value of
,
a higher accuracy for the N(0) and N(1) column densities is needed
to verify that
.
For
cm-3, the linear thickness of the H2-bearing cloud is small,
pc. Similar characteristics of molecular hydrogen small structures are found
in intermediate-velocity clouds (IVC) in the Milky Way halo (Richter et al. 2003b).
Acknowledgements
We thank our referee P. Richter for his comments and remarks. S.A.L. acknowledges the hospitality of Hamburger Sternwarte, Universität Hamburg. The work of S.A.L. is supported in part by the RFBR grant No. 03-02-17522. R.Q. is supported by the Verbundforschung of the BMBF/DLR under Grant No. 50 OR 9911 1.