A&A 435, 17-27 (2005)
DOI: 10.1051/0004-6361:20041655
D. Reimers1 - E. Janknecht1 - C. Fechner1 - I. I. Agafonova2 - S. A. Levshakov2 - S. Lopez3
1 - Hamburger Sternwarte, Universität Hamburg,
Gojenbergsweg 112, 21029 Hamburg, Germany
2 -
Department of Theoretical Astrophysics,
Ioffe Physico-Technical Institute, 194021 St. Petersburg, Russia
3 -
Departamento de Astronomia, Universidad de Chile,
Casilla 36-D, Santiago, Chile
Received 13 July 2004 / Accepted 19 January 2005
Abstract
HE 0141-3932 (
= 1.80)
is a bright blue radio-quite quasar with
an unusually weak Ly
emission line.
Large redshift differences (
)
are observed
between high ionization and low ionization emission lines.
Absorption systems identified at
= 1.78, 1.71, and 1.68
show mild oversolar metallicities (
)
and can be attributed
to the associated gas clouds ejected from the circumnuclear region.
The joint analysis of the emission and absorption lines leads to the
conclusion that this quasar is seen almost pole-on.
Its apparent luminosity may be Doppler boosted by
10 times.
The absorbing gas shows a high abundance of Fe, Mg and Al
(
)
along with underabundance of N (
).
This abundance pattern is at variance with current chemical evolution models
of QSOs predicting
and
at
.
Key words: cosmology: observations - line: formation - line: profiles - galaxies: abundances - quasars: absorption lines - quasars: individual: HE 0141-3932
In the course of a high-resolution study of the Ly
forest at
intermediate redshifts (
)
in bright quasars
from the Hamburg/ESO Survey with
the UV-visual echelle spectrograph (UVES) at the VLT, the
discovery spectrum of the QSO HE 0141-3932 (
,
B = 16.2, Wisotzki et al. 2000)
attracted our attention by two facts: it appeared to have no or only very
weak Ly
emission along with clearly recognizable
Mg II, and it showed several
absorption line systems at
ranging from 1.78 to 1.68 with very complex
and strong metal profiles which allow us to
suggest that these systems may originate in the ejected gas.
Since the identification spectrum of HE 0141-3932 was of only moderate
quality, we took further low-resolution spectra with EFOSC 2 with the
ESO 3.6 m telescope to improve the redshift measurement and
we found a third peculiarity, namely large redshift differences
(
)
between high ionization and low ionization emission
lines.
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Figure 1: EFOSC 2 spectra of HE 0141-3932. The spectral resolutions are indicated in the panels. All identified emission features are labeled. The parameters of the prominent emission lines are listed in Table 2. |
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Emission line spectra similar to that of HE 0141-3932
are observed in some high redshift BL Lacertae (BL Lac) objects
(Urry & Padovani 1995).
However, HE 0141-3932 is a radio-quiet quasar, i.e. not a blazar.
A few other blue radio-quiet quasars are known with
apparently missing or weak Ly
emission but clearly present
metal lines:
PG 1407+265 (
= 0.94, McDowell et al. 1995),
Tol 1037-2703 (
= 2.20, Srianand & Petitjean 2001),
PHL 1811 (
= 0.192, Leighly et al. 2004),
and several cases from SDSS (Fan et al. 1999, 2003).
The nature of their peculiar emission line spectra is not clear.
In particular, Leighly et al. (2004) suggest a high accretion rate which powers the UV emission from an optically thick accretion disk, while suppressing the formation of a hot corona. However, recent radio observations of PG 1407+265 on the milliarcsecond scale revealed a relativistic jet of moderate power beamed toward us (Blundell et al. 2003). This means that the quasar is seen almost pole-on and its emission line spectrum may be diluted by continuum radiation both from the jet and the accretion disk.
In the present paper we analyze both the emission and absorption line spectra of HE 0141-3932. We expect that this joint consideration may shed some light on the unusual properties of this quasar.
The paper is organized as follows:
observations are described in Sect. 2,
the emission lines are studied in Sect. 3, the analysis of
the
absorbers is given in Sect. 4,
the results are discussed and summarized in Sect. 5.
Table 1: Log of spectroscopic observations of HE 0141-3932.
HE 0141-3932 was observed with the UVES at the 8 m ESO VLT on Paranal over 7 nights in July/August 2001. Eleven individual exposures with integration times of 60 min were made using the dichroic mode in standard settings (Table 1). With a slit width of 1'', a resolution of 41 000 (7.3 km s-1) in the blue and 38 000 (7.9 km s-1) in the red is achieved.
The data reduction was performed at the Quality Control Group in Garching using the UVES pipeline Data Reduction Software (Ballester et al. 2000), the vacuum-barycentric corrected spectra were co-added. The resulting signal-to-noise ratio, S/N, is typically 75.
Since both the Ly
and the Mg II emission line ranges
are covered only by UVES spectra, we performed an absolute flux calibration
using an appropriate master response curve and the respective airmasses
during the observations. The procedure is described on the ESO web page
www.eso.org/observing/ dfo/quality/UVES/qc/response.html.
Further spectra were taken with EFOSC 2 at the ESO 3.6 m telescope
on October 2, 2003, to improve our knowledge
about the redshifts of the emission lines. Details are given in Table 1.
Wavelength and flux calibration was performed according to standard
procedures. The resulting spectra are shown in Fig. 1.
Notice that due to the rapidly decreasing sensitivity below 3500 Å,
the EFOSC 2 spectra only partially cover the wavelength range expected for
the Ly
emission line.
The combination of EFOSC 2 spectra with roughly flux calibrated
UVES spectra allows us to estimate redshifts
and equivalent widths of all emission lines between
Ly
and Mg II
Å.
Equivalent widths (EW) and redshifts were measured by fitting
Gaussian profiles to the data.
Table 2 presents the results.
The Ly
EW is relatively uncertain since in the calibrated
UVES data the continuum is difficult to determine and the blue
wing of the line is incomplete (see Figs. 1 and 2).
The identification of the emission line at 3400 Å is ambiguous:
it could be either Ly
at z = 1.80 or N V at
z = 1.75 or blend of both lines.
There are two arguments in favor of the N V identification:
it has a width comparable to C IV and its redshift - if identified
as N V - is equal to that of the C IV,
while in QSOs with significant redshift differences between high
ionization lines (C IV) and low ionization lines
(Mg II, H
), the Ly
line typically is found
at the redshift of the
high ionization BLR emission lines (Gaskell 1982; Espey et al. 1989).
On the other hand, the emission line at 3400 Å is quite strong for
N V, although there are quasars known with
N V strength comparable to that of C IV (Hall et al. 2004;
Baldwin et al. 2003b).
Even if this line would be entirely due to emission of neutral hydrogen,
we can conclude that Ly
in HE 0141-3932 is unusually weak.
Table 2: Estimated equivalent widths, redshifts and relative velocities for emission lines of HE 0141-3932.
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Figure 2: Emission lines of HE 0141-3932 (spectra from UVES and EFOSC 2). v = 0 km s-1 corresponds to z = 1.80. For orientation, at v = 0 km s-1 and v = -5000 km s-1 dashed lines are plotted. |
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In Fig. 2 we display the emission line profiles on a velocity scale relative
to z = 1.80. The redshift of the low ionization lines of
z = 1.80 appears to be roughly the systemic redshift,
since the UVES spectrum shows that the Ly
forest starts
at
z = 1.8086.
While the velocity shift in the emission lines is extremely large,
how unusual are the emission line intensities?
This can be best seen by a comparison with the histogram of equivalent
widths of the LBQS-quasars, a well selected sample
of QSO given by Francis (1993).
The Ly,
C IV, C III] and
Mg II equivalent widths are among the bottom 2%, 1%, 13%
and 1.5%, respectively, of the Francis (1993) distribution.
At the same time,
the Al III 1856/64 and Fe III 2075 lines appear
unusually strong.
In order to investigate the physical properties of gas
which could produce such an unusual emission line spectrum
we compared the observed equivalent widths of HE 0141-3932
with the compilation
of quasar BLR rest-wavelength emission line equivalent widths given
as functions of column density, incident ionizing
spectrum, and metal abundance of the emitting gas clouds
(Korista et al. 1997).
While Ly
,
Mg II and C III]
have roughly the same velocity (redshift) and could originate in the
same volume, we were unable to find a parameter combination which met
the constraints on Ly
and Mg II line strengths simultaneously -
the theoretical Ly
is always much too strong for a
parameter combination that matches the Mg II equivalent width.
The C IV/Ly
EW ratio can be met
assuming high gas densities
[log n(H)
12], but
the discrepant redshifts exclude formation in the same location.
A mean BLR parameter set appears to be unable to reproduce the
measured EWs of emission lines in HE 0141-3932.
In order to estimate the physical parameters of the absorption systems we used the Monte Carlo Inversion (MCI) method. Detailed description of the MCI is given in Levshakov et al. (2000, 2002, 2003a). Here we outline briefly its basics needed to understand the results presented.
The MCI is based on the assumption
that all lines observed in the absorption system are formed in
a continuous absorbing
gas slab of thickness L
where the gas density,
,
and velocity, v(x),
fluctuate from point to point giving rise to complex profiles
(here x is the space coordinate along the line of sight).
We also assume that within the absorber the metal abundances are constant, the gas is optically thin for the ionizing UV radiation, and the gas is in thermal and ionization equilibrium. The intensity and the spectral shape of the background ionizing radiation are treated as external parameters.
The radial velocity v(x) and gas density
are considered as two continuous random functions which are
represented by their sampled values at equally spaced intervals
.
The computational procedure is based on
adaptive simulated annealing. The fractional ionizations of
different elements at each space coordinate x,
,
are computed
with the photoionization code CLOUDY (Ferland 1997).
The following physical
parameters are directly estimated by the MCI procedure:
the mean ionization parameter U0,
the total hydrogen column density ,
the line-of-sight velocity dispersion,
,
and
density dispersion,
,
of the bulk material
[
],
and the chemical abundances
of all elements
involved in the analysis.
With these parameters we can further calculate
the column densities for different species
,
and the mean kinetic temperature
.
In general, the uncertainties on the fitting parameters U0, ,
,
,
and Za are about 15%-20%
(for data with
)
and the errors of the estimated column densities are less than 10%.
However, in individual absorption systems
the accuracy of the recovered values can
be lower for different reasons like partial blending of the line profiles,
saturation, or absence of lines of subsequent ionic transitions.
The MCI can be supplemented with an additional procedure
aimed at restoring the shape of
the background ionizing spectrum. A formal
description of this procedure is given in the Appendix.
Its accuracy depends significantly on the number of
unsaturated lines of the subsequent
ionic transitions of different elements
(e.g. Si II/Si III/Si IV,
C II/C III/C IV) available for the analysis.
The absorber at
= 1.7817 (described below)
reveals enough such lines and we use them to estimate
the spectral shape of the background UV radiation in the range E > 1 Ryd.
All calculations are carried out with the laboratory wavelengths and oscillator strengths taken from Morton (2003). Solar photospheric abundance for carbon is taken from Allende Prieto et al. (2002), for silicon, nitrogen and iron - from Holweger (2001).
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Figure 3:
Hydrogen and metal absorption lines associated with the
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Many unsaturated lines of different ionic transitions along with a
single saturated hydrogen line Ly
are detected in this system (Fig. 3)
separated by 2000 km s-1 from the central source.
The blue wing of Ly
is partially blended by the adjacent
hydrogen line [N(H I)
cm-2 ] from an
intervening absorber (only a weak C IV
doublet is detected in this system).
C III
is
beyond the wavelength coverage and Si III
is blended, probably
with Ly
absorption from the
= 1.7606.
We started the analysis assuming standard ionizing backgrounds such as a
power law
(with different indexes
),
the mean intergalactic spectrum (at z = 1.8)
of Haardt & Madau (1996),
and the AGN-type spectrum of Mathews & Ferland (1987, hereafter MF).
None of these spectra was consistent with the observed intensities of
the C II, C IV, Si II and Si IV lines: all
trials underproduced C II/C IV and overproduced
Si II/Si IV,
but the MF spectrum provided the lowest .
All spectra gave the mean ionization parameter in the range of
.
In spite of the saturation, the neutral hydrogen column density can be
estimated with a sufficiently high accuracy (
20%) since the velocity
dispersion of gas is determined by numerous metal lines detected in this
system (the procedure to restore partly blended profiles is
described in Sect. 3 in Levshakov et al. 2003a).
All runs with different model spectra
showed a rather high metallicity - slighty above solar value, and
an almost solar ratio of Si/C.
Taking into account this preliminary information we can assume that the UV
spectrum to be found should maximize the product of ratios
C II/C IV and Si IV/Si II
in the range of
for solar
metal content and solar element abundances.
The spectral shape of the MF spectrum can be taken as
a first approximation (for more details see Appendix).
Applying the adjustment procedure as described in Appendix,
we estimated a new shape of the ionizing spectrum which ensured
a self-consistent description of all lines observed in the
= 1.7817 system. This spectrum
is shown in Fig. 4 by the dashed line, whereas the initial MF
spectrum is the solid line. As seen in Fig. 4, the
restored spectrum is softer in the range 2 Ryd < E < 6.5 Ryd
but much harder above 6.5 Ryd.
It shows a break at the He II ionization edge with the
amplitude
(A)/
(B)
corresponding to the optical depth
(He II)
1.
Physical parameters estimated with this modified ionizing background are
listed in Table 3, Col. 2, and the corresponding synthetic profiles
are plotted in Fig. 3 by the smooth curves.
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Figure 4:
Spectral
shape of a typical AGN
ionizing continuum from Mathews & Ferland (1987) shown by the solid line
and its modification (dashed line) estimated for the
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This absorption system spans the velocity range of 500 km s-1 and
is located in velocity space at a distance of 9600 km s-1 from the QSO.
It reveals many lines of different ionic transitions (Fig. 5).
The structure of the Ly,
C IV, Si IV and N V profiles indicates that the system may be divided into
two subsystems: one at
-200 km s-1 < v < 180 km s-1 (A) and another at
180 km s-1 < v < 320 km s-1 (B).
The subsystem B has an unsaturated hydrogen line and hence
allows us to estimate the hydrogen column density
with a sufficiently high accuracy.
The physical parameters computed with the UV background deduced for the
= 1.7817 system are given in Table 3, Col. 4.
The obtained extremely high metallicity
- almost
- is striking.
Since silicon lines Si III
and
Si IV
are very weak and the element ratios ([Si/C], [N/C]) not known a priori,
the mean ionization
parameter
represents a lower limit
consistent with the observed intensities of the
C IV lines and the absence of the C II
absorption.
Formally this subsystem could be fitted
with any higher ionization parameter producing even higher
metallicity and higher [Si/C] and [C/N] ratios.
To confirm the high metallicity in this subsystem
we repeated calculations assuming other ionizing backgrounds - power
laws with
ranging between 1.0 and 1.8 and the MF spectrum.
These trials also delivered metallicity of
about 10 solar.
In principle,
such high metal enrichment of the circumnuclear gas is predicted
in some models of chemical evolution of QSO/host galaxies
(Hamann & Ferland 1999).
However, another explanation of our results is that the absorbing
gas is not in equilibrium with the ionizing background, e.g.,
is still cooling and recombining.
In this case high metallicities can be caused by a longer recombination
time of hydrogen as compared to C IV and N V
(e.g., Osterbrock 1989).
The ionization parameter U and, hence, the total hydrogen column density
is determined by the observed
line intensities of different ions.
Since hydrogen is ionized higher than it would be in the ionization equilibrium
with the same parameter U, the total hydrogen becomes underestimated
leading to an artificially high metallicity.
We may expect that the adjacent subsystem A (centered at v=0 km s-1)
would help to clarify the
physical conditions in the
= 1.7103 absorber.
The subsystem A shows complex profiles of low ionization
lines C II, Mg II, Si II, and
Fe II along with Al III, Si III,
Si IV, N V and saturated C IV (see Fig. 5).
The computations had been carried out with the ionizing background from the
= 1.7817 system. The presence of the subsequent ionic transitions guarantees
the accuracy of the mean ionization parameter of
10%
(for a given ionizing background).
The assumption of constant metallicity inside the
absorber turned out to be inconsistent with the observed blue
wing of Ly
and the absence of any metal absorption in the
range -200 km s-1 < v <-100 km s-1.
A single available hydrogen line does not allow us
to conclude whether the blue wing is blended or
there is indeed a gradient of metal content.
The former possibility looks more probable since
there is a non-zero flux in
the Ly
profile at -110 km s-1 < v < -90 km s-1
with the mean intensity
(marked by the arrow in Fig. 5).
The constant metallicity throughout this sub-system
seems to be appropriate since very steep blue wings of several
ions (Si III, Si IV, C IV) do not indicate
a progressive dilution.
Thus, the calculations were carried out with the assumption of
constant metallicity. The neutral hydrogen column density and the shape of
the blue wing of the synthetic Ly
shown in Fig. 5
were calculated with the velocity
and density fields estimated from the observed metal line profiles
(for details see Sect. 4.3 in Levshakov et al. 2003a).
The obtained physical parameters are given in Table 3, Col. 3.
We also tried
other ionizing backgrounds (different power laws and the MF spectrum), but
none of them was consistent with the observed relative
intensities within the C II/C IV
and Si II/Si III/Si IV profiles.
Table 3:
Physical parameters of the
= 1.7817, 1.7103 and 1.6838 metal absorbers
toward HE 0141-3932 (
= 1.80) derived by the MCI procedure
(limits are given at the 1
level).
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Figure 5:
Same as Fig. 1 but for the
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Figure 6:
Same as Fig. 1 but for the
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In spite of all uncertainties intrinsic to this subsystem several
parameters were steadly reproduced in all runs. These are (1)
slightly undersolar nitrogen abundance, [N/C] =
;
(2) significantly
lower ratios of Si/C and Al/C as compared to their solar values,
[Si/C] =
,
[Al/C] =
;
and (3) extremely high
overabundance of iron [Fe/H]
.
The extremely high overabundance of iron as well as underabundances of
silicon and aluminium may be caused by non-equilibrium ionization.
When an absorbing gas comes to equilibrium through cooling and recombining
its ionization parameter can be overstated due to different
cooling and recombination
times of the observed ions
stemming from density fluctuations (so-called "hot photoionization'').
A higher density gas has shorter cooling and recombination times and
hence comes faster to equilibrium. Thus, the sub-system Amay be described as consisting of dense
gas clumps that are already close to equilibrium
(seen in low ionization absorptions)
and ambient rarefied gas still far from equilibrium
(responsible for most of C IV and N V absorptions).
The equilibrium ionization parameter should be lower,
probably ranging between 0.001 and 0.003.
With this U0, only a
mild overabundance of iron,
,
would be enough to describe
the observed intensity of the Fe II lines.
Thus, we estimate a metallicity of
in
the higher ionized subsystem B and
in the lower ionized subsystem A(both metallicities are referred to silicon lines).
This strong metallicity difference is
another argument in favor
of non-equilibrium ionization in sub-system B.
It probably has lower gas density than in A which leads
to longer cooling and recombination times.
We suggest that some time ago the whole absorbing complex
(A+B) was either
exposed to a much more intense radiation
or shock-heated up to the
temperatures when collisional ionization becomes significant (Klein et al.
1994, 2003; Levshakov et al. 2004).
This system separated by
km s-1 from the QSO shows
a partially saturated Ly
and many metal lines
with complex profiles (Fig. 6).
Si II
and N V
are blended with
Ly
forest absorptions, and at the position
of Si II
a clear continuum
window is seen.
The apparent structure of the Ly
profile suggests that
this system can be divided
into two subsystems:
one at -250 km s-1
< v < -180 km s-1 (A) and another
at -180 km s-1 < v < 150 km s-1 (B) which
were analyzed separately.
Calculations were carried out using both the ionizing spectrum
restored for the
= 1.7817 system and several power law spectra.
Physical parameters
listed in Table 3, Cols. 5, 6 correspond to the
= 1.7817 background.
The red wing of the synthetic profile of Ly
in the subsystem Bis calculated simultaneously with
metal lines assuming a constant metallicity inside the absorber.
The ionization parameter given in Table 3 for the subsystem A is determined
from the observed intensity of C IV and the noise level
at the expected position
of the C II
line and should be considered
as a lower limit.
The almost solar ratio of [Si/C] =
and a strong
underabundance of nitrogen [N/C] < -0.5 obtained for
both subsystems indicate that photoionization is probably
close to equilibrium.
However, slightly higher metallicity in the sub-system A and
the remaining flux
at the
shallow bottom of Ly
may indicate that hydrogen has not yet reached its equilibrium
with the ionizing background (see sub-system B in the foregoing section).
Although the absolute values of the abundances are uncertain (they depend on the
assumed background and the mean ionization parameter),
the solar to oversolar carbon content,
the ratio
and a significant underabundance of nitrogen are
constantly reproduced in all trials.
Other systems with unusually strong and complex C IV profiles
are identified in the spectrum of HE 0141-3932.
Unfortunately, they cannot be analyzed by the MCI because their
hydrogen lines are unavailable, and we describe them only qualitatively.
The system at
= 1.7365 (separated by
7000 km s-1 from the QSO) shows also a clear N V doublet
(Ly
profile is blended).
The apparent ratio C IV/N V
is very much like that in the sub-system Aat
= 1.6838 and the
= 1.7365 system probably has similar
physical parameters.
The system at
= 1.4978 (
km s-1) exhibits
a strong and complex Si IV absorption
and a weak N V doublet as well as C II
(L
is out of range),
and in these features
resembles the sub-system B at
= 1.6838.
The oversolar metallicity obtained in the absorbers described above
place them in a class of associated systems.
Furthermore, their extremely large radial velocities
imply that they originate in gas ejected from
the circumnuclear region of the QSO/host galaxy.
This is also supported by the relatively strong emission flux in
Fe II and Fe III lines and at the same time by
the Fe II absorptions in the associated systems.
Note that Fe II lines in
absorbers with N(H I)
cm-2 are extremely rare; probably this is the first such detection.
The distance between the absorbing cloud and the light source
can be estimated from
the photoionization model and
the column density ratios of C II/C II or
Si II
/Si II.
For an ion in the interstellar (intergalactic) medium, the ratio
of excited to ground-state population is equal to the ratio of the
collisional excitation rate
to the spontaneous
transition probability
(Bahcall & Wolf 1968):
We do not detect Si II
or
C II
lines in the associated systems, therefore
clear "continuum windows'' at the expected positions of
C II
and
Si II
were used to
set upper limits on the column densities of
C II
and Si II
(Table 3).
For the 3
upper limits on
N(Si II
)
and N(C II
), Eq. (1) provides
cm-3 in the
= 1.78, and 1.71 (A) systems, and
cm-3 in the
= 1.68 (B) system.
Since the degree of ionization in these systems is
high (
), the upper limits on
the total gas density
are the same, i.e.,
cm-3 and
cm-3 , respectively
(the contribution of
the ionized helium is ignored since it has a small effect).
To estimate the distance, the QSO continuum luminosity
at the Lyman limit
must be known.
Absolute spectrophotometry is not available for HE 0141-3932.
Therefore we estimated
the intrinsic luminosity
from (i) the comparison of the QSO B magnitude
(assuming the QSO spectral energy distribution is a MF-type, i.e.,
in the range 912 Å
Å)
with the
specific flux of a star having
mB = 0.0 outside the Earth's atmosphere
(
erg s-1 cm-2 Hz-1) and from (ii)
the empirical formula given by Tytler (1987, Eq. (17))
which is obtained from a fit to
data on
over 60 QSOs with a wide range of redshifts.
The observed B magnitude of 16.09 (corrected for extinction)
translates to the flux
(4400 Å) =
erg s-1 cm-2 Hz-1which in turn leads to the apparent luminosity near the Lyman limit
erg s-1 Hz-1(H0 = 71 km s-1 Mpc-1,
,
and
are used to calculate
the luminosity distance of 13 Gpc).
The second method gives
erg s-1 Hz-1which shows that both values are in reasonable agreement.
In the following we use the
second one since it is based on the observational data.
Given the upper limits on ,
the distance from the QSO to the absorbing cloud
can be calculated from the estimated ionization parameter U
which is defined as
With the estimated Lyman continuum luminosity
erg s-1 Hz-1,
one finds
photons s-1,
assuming
,
and
in the range
.
A substitution of the numerical values in Eq. (2) provides
r1.68 > 100 kpc,
r1.71 > 280 kpc,
and
r1.78 > 450 kpc.
The characteristics of both the broad emission lines and associated
absorptions of HE 0141-3932
can be best explained by the almost pole-on view of this quasar.
The weakness of emission lines is then due to dilution by direct
radiation from the accretion disk, whereas
the velocity shifts of the associated systems and their distances from the
light source can be caused by the entrainment into the large-scale
outflowing jet.
However, powerful jets propagating to distances of 0.5 Mpc
imply the existence of significant radio radiation, but
HE 0141-3932 is a radio-quiet QSO.
An upper limit to its radio flux of 1 mJy translates into
the intrinsic radio luminosity
erg s-1 Hz-1.
This is a radio power of a typical FR I source. Jets in the FR I-type
structures are known to be subsonic or slightly oversonic with velocities
1000-10 000 km s-1, heavier than and isobaric with the external
medium (e.g., Hughes 1991).
These characteristics are in line with the suggestion that the
observed absorptions originate in the entrained clouds.
It is also known that FR I jets do not show an extended radio emission.
This means that radio observations of such jets, especially when they are
seen at a small angle to the line of sight,
need a high spatial resolution and high sensitivity.
Nevertheless, the calculated
distances of several hundreds of kiloparsecs seem to be 4-5 times
overestimated since the typical jet length of the FR I object is below
100 kpc. We cannot explain the source of this discrepancy.
It should be noted, however,
that distances of hundreds of kiloparsecs
for the associated systems are not exceptional
(see, e.g., Morris et al. 1986; Tripp et al. 1996;
D'Odorico et al. 2004).
The second fact that is hard to explain is the large velocity excess
between low- and high-ionization emission lines.
According to a model of the quasar atmosphere (e.g., Elvis 2004),
low-ionization emission
lines (H,
O I, Fe II, Mg II, C III])
come from the outer region of the accretion disk which is
well shielded from the central source radiation and optically
thick for Ly
.
Ly
and high-ionization emission lines (C IV, Si IV,
N V)
are formed in a cool phase of the wind arising from the inner parts of
the accretion disk.
The emission at 3400 Å observed in HE 0141-3932 cannot be
interpreted unambiguously, but if it is (even partially) due to
Ly
at z = 1.80, then
its redshift and strength is inconsistent with
this scenario since C IV and Si IV are seen at z=1.75.
The same situation (i.e.,
)
is observed in PG 1407+265.
Apparently, a complex geometrical model of the broad
emission line region is necessary to explain the observations.
As already mentioned, the relativistic jet beamed toward us
was discovered in the bright (B = 15.7) and radio-quiet PG 1407+265
(Blundell et al. 2003).
There are indications that HE 0141-3932
may also have a similar small-scale
relativistic jet.
The Lyman limit luminosity of HE 0141-3932 is
erg s-1 Hz-1, whereas
its radio luminosity is less than
erg s-1 Hz-1.
This implies either an unusually flat radiation continuum (typically for
QSOs is
at
Hz and
above), or that the apparent optical luminosity is
Doppler-boosted due to the relativistic motion of the light source.
Another argument in favor of boosting comes from the metallicity
estimations.
HE 0141-3932 is a bright source. It is supposed that
the QSO luminosity is determined by the accretion rate which requires
a large amount of circumnuclear gas. This, in turn, supposes
a large mass of the contributing stellar population and, hence,
a high metal enrichment of the accreting gas.
For reference, all QSOs with luminosities
above
erg s-1 Hz-1in the sample of Dietrich et al. (2003)
have metallicities
with the mean value of
.
The metallicity of
has been also measured in the associated
system of a very bright QSO HE 0515-4414 (Levshakov et al. 2003b).
Gas in HE 0141-3932
has metallicity
,
which supposes
a relatively low stellar population involved in the enrichment, low
accretion rate and, hence, low luminosity.
Boosting can be caused by a relativistic jet seen at a small viewing angle.
Taking into account the measured equivalent widths of the emission lines,
we may assume that the luminosity is amplified by
10 times,
i.e. the intrinsic luminosity of HE 0141-3932 might be only
erg s-1 Hz-1.
Another issue of interest is the relative abundances
obtained for the absorbing gas.
The analyzed absorption systems show high iron content,
[Fe/C] =
,
[Fe/Mg] =
(
= 1.78),
but at the same time nitrogen
is strongly underabundant, [N/C]
(
= 1.68).
Although these values were estimated in different absorption systems
they are representative for the bulk of circumnuclear gas for the
following reason.
The mass of the stellar population involved in the
enrichment of a quasar's circumnuclear region is
(Baldwin et al. 2003a) and hence large metallicity
gradients and sharp discontinuities
due to enrichment by only a few stars
are unlikely.
In the Hamann & Ferland (1993, 1999) models of QSO chemical evolution,
solar metallicity is reached after
0.2 Gyr and is characterized by
a relative overabundance of nitrogen,
,
and an underabundance of iron,
.
Due to the delay of 1 Gyr in Fe enrichment expected from
the longer evolution of SNe Ia which are the main sources of iron,
the emission line ratios Fe II/C IV and
Fe II/Mg II were proposed as a clock to constrain the QSO ages.
However, in these models, large values of [Fe/C] and [Fe/Mg]
are always associated with a considerable overabundance of nitrogen,
.
We do not observe such
behavior in our systems and, hence, cannot confirm this "iron clock''. This
is in line with the result of
Matteucci & Recchi (2001), who showed that the time scale
for enrichment by SNe Ia is not unique but
a strong function of the adopted
stellar lifetimes, initial mass function and star formation rate and can
vary by more than an order of magnitude.
Acknowledgements
The work of S.A.L. and I.I.A. is supported by the RFBR grant No. 03-02-17522 and by the RLSS grant 1115.2003.2. C.F. is supported by the Verbundforschung of the BMBF/DLR grant No. 50 OR 9911 1. S.L. acknowledges support from the Chilean Centro de Astrofísica FONDAP No. 15010003, and from FONDECYT grant N.
In general, the ionizing UV background is produced by the attenuated UV flux of QSOs and/or other sources (like young galaxies) and by the recombination radiation from intergalactic gas clouds. To estimate the shape of the UV continuum from the observed profiles of ions in the metal systems we use the response function methodology from the theory of experimental design (see, e.g., Box et al. 1978, Chap. 15).
At first, the shape of the UV continuum in the range E > 1 Ryd
is to be parameterized
by means of k variables (factors).
For instance, the AGN-type spectrum
can be described by a broken power-law defined by
the following factors (see Fig. 4):
f1 - the power law index (slope)
in the range 1 Ryd < E < f2 Ryd; f2 - the coordinate of the
first fracture;
f3 the slope between f2 and the point of the second fracture f4;
f5 the value
of the break after f4 (decimal logarithm of the intensities ratio)
with the slope f6;
f7 the slope in the high UV range.
Formally we fix the starting point of the X-ray break at E = 100 Ryd
and the slope after it at -1.5
since this spectral range affects very weakly the fractional ionizations of
ions we are interested in.
In this factor space, the Mathews-Ferland (MF) spectrum is represented by a
point with the coordinates
.
At the start of the procedure we select
a basic UV spectrum (e.g. power law, HM, or MF) and
carry out the MCI calculations using the fractional
ionizations
produced by this background.
If the trial
spectrum turns out to be inconsistent with the observed line intensities, we
assign its parameters to a
"null point'' in the factor space and vary
the factors around this null point according to the chosen set
of treatments. Thus, for each treatment we obtain a new trial UV spectrum.
The fitness of a particular spectrum is measured by the
value of the response function which is defined individually for
each absorption system in such a way as
to ensure the self-consistent description
of all absorption lines detected in the system.
The best UV spectrum is that which provides
a maximum value of
within the constraints
(per degree of freedom) for each
individual line.
The information needed to construct the response function
is obtained in several test runs with the "null spectrum''.
The following example illustrates this step of the procedure.
Let us assume that an absorption system exhibits lines of
Si II, Si IV and C II, C IV. Trial calculations
with a basic UV continuum show that this continuum overpredicts the
intensity of Si II and underpredicts Si IV whereas
the behavior of carbon lines is opposite - underpredicted C II
and overpredicted C IV. In this case it is conceivable to search
for such a background that will produce a maximum value of the product
Si IV/Si II and C II/C IV.
The line intensities (and column densities)
are proportional to the fractional ionizations and
we can calculate the response function simply as
After the response function has been evaluated at each point of the
experimental plan, factor weights (influences)
can be estimated. For this purpose a standard polynomial model
is used:
After the first iteration of the spectral shape adjustment is completed,
we repeat the MCI calculations with the designed UV background
and, if necessary, go back to the adjustment procedure
until a satisfactory result (i.e. normalized
for
all profiles included in the analysis) is achieved.