A&A 374, 871-877 (2001)
DOI: 10.1051/0004-6361:20010777

High-resolution OVI absorption line observations
at 1.2 $ \leq z \leq $ 1.7 in the bright QSO HE 0515-4414[*]

D. Reimers 1 - R. Baade 1 - H.-J. Hagen1 - S. Lopez2

1 - Hamburger Sternwarte, Universität Hamburg, Gojenbergsweg 112, 21029 Hamburg, Germany
2 - Departamento de Astronomia, Universidad de Chile, Casilla 36-D, Santiago, Chile

Received 7 March 2001 / Accepted 25 May 2001

STIS Echelle observations at a resolution of $10\,{\rm
km\,s^{-1}}$ and UVES/VLT spectroscopy at a resolution of $7\,{\rm
km\,s^{-1}}$ of the luminous QSO HE 0515-4414 ( $z_{\rm em} =
1.73$, B = 15.0) reveal four intervening O VI absorption systems in the redshift range $1.21 \leq z_{\rm abs} \leq 1.67$(1.38503, 1.41601, 1.60175, 1.67359). In addition, two associated systems at z = 1.69707 and z = 1.73585 are present. Noteworthy is an absorber at z = 1.385 with $\log\,N_{\rm H{\sc i}} =
13.9$ and strong O VI (N(O VI)/N(H I) $\approx$ 1) and C IV doublets, while a nearby much stronger Ly $\alpha$ absorber (log  $N_{\rm H{\sc i}} = 14.8$, $\Delta v =
123\,{\rm km\,s^{-1}}$) does not reveal any heavy element absorption. For the first time, high resolution observations allow one to measure radial velocities of H I, C IV and O VI simultaneously in several absorption systems (1.385, 1.674, 1.697) with the result that significant velocity differences (up to $18\,{\rm km\,s^{-1}}$, are observed between H I and O VI, while smaller differences (up to $5\,{\rm km\,s^{-1}}$) are seen between C IV and O VI. We tentatively conclude that H I, O VI, and C IV are not formed in the same volumes and that therefore conclusions on ionization mechanisms are not possible from the observed column density ratios O VI/H I or O VI/C IV. The number density of O VI absorbers with $W_{\rm rest} \geq 25\,$mÅ is ${\rm d}N/{\rm d}z
\leq 10$, roughly a factor of 5 less than that found by Tripp et al. (2000) at low redshift. However, this number is uncertain and further lines of sight will be probed in the next HST cycle. An estimate of the cosmological mass-density of the O VI-phase yields $\Omega_{\rm b}(O{\sc vi}) \approx
0.0003\, h^{-1}_{75}$ for $\rm [O/H] = -1$ and an assumed ionization fraction O VI/O = 0.2. It should be noted that this result is subject to large systematic errors. This corresponds to an increase by roughly a factor of 15 between $\bar{z} = 1.5$ (this work) and the value found by Tripp et al. (2000) at $\bar{z} = 0.21$, if the same oxygen abundance $\rm [O/H] = -1$ is assumed. Agreement with the simulations by Davé et al. (2001) can be obtained, if the oxygen abundance increases by a factor of $\sim$3 over the same redshift interval.

Key words: cosmology: observations - intergalactic medium - quasars: absorption lines -
quasars: individual: HE 0515-4414

1 Introduction

Recent observations of intervening O VI absorbers in HST-STIS Echelle spectra of bright, low redshift QSOs have provided strong evidence that in the local universe a considerable fraction of baryonic matter might be "hidden'' in a warm ($\sim$$10^{5}\,$K) intergalactic medium (Savage et al. 1998; Tripp et al. 2000; Tripp & Savage 2000). This observation is in accordance with models of hierarchical structure formation by Cen & Ostriker (1999) and Davé et al. (2001) which predict that a considerable fraction of all baryons reside in a warm-hot phase of the intergalactic medium (WHIM) shock-heated to temperatures of $10^{5}{-}10^{7}\,$K. The same models predict that the fraction of baryons residing in this WHIM increases strongly with decreasing redshift from less than 5% at z = 3 to 30-40% at z = 0. Can this prediction be verified or disproved by observations, or can observations even impose constraints on the models? This appears difficult for various reasons. First of all, the WHIM is difficult to detect (cf. Davé et al. 2001), both as diffuse X-ray emission of the hotter parts or in absorption through the O VI doublet. In addition, the temperature distribution of the WHIM varies with redshift so that a complete census would require the detection of all components as a function of redshift. The warm O VI component has the additional complication that both the oxygen abundance and the ionization process cannot be determined from O VI observations alone. While at low redshift (z < 0.3) collisional ionization is the most probable process since the ionizing extragalactic UV background is diluted, O VI can be produced easily by photoionization at redshifts $\geq 2$ and has been observed to be ubiquitous in the low-density IGM (Schaye et al. 2000). On the other hand O VI is not expected to be produced by photoionization for $z \geq 3$ since the reionization of He II is incomplete (Reimers et al. 1997; Heap et al. 2000) and the IGM therefore opaque to photons with energies above 4 Rydberg. The intermediate redshift range remains which for z < 1.9 requires high-resolution UV-spectroscopy of a bright, high-redshift QSO. In this paper, we present combined high-resolution HST/STIS observations of O VI absorption supplemented by ESO-VLT/UVES spectroscopy of the accompanying H I and C IV lines in the brightest known intermediate redshift QSO HE 0515-4414 ( $z_{\rm em} =
1.73$, B = 15.0) discovered by the Hamburg/ESO Survey (Reimers et al. 1998). The data have been taken mainly with the aim of studing the evolution of the Ly $\alpha$ forest and its metal content in the range z = 1 to 1.7. In this first paper we concentrate on the intervening O VI absorption.

2 Observations and data reduction

2.1 Hubble Space Telescope observations

HE 0515-4414 was observed with STIS for 31500 s on three occasions between January 31 and February 2, 2000 with the medium resolution NUV echelle mode (E230M) and a $0.2 \times 0.2$ aperture which provides a resolution of $\sim$30000 ( $FWHM \simeq 10\,{\rm
km\,s^{-1}}$). We used the HST pipeline data with an additional correction for inter-order background correction (Rosa, private communication). The spectrum covers the range between 2279Å and $\sim$3080Å. The coverage at the red end guarantees overlap with the UVES spectra which extend shortwards to $\sim$3050Å.

2.2 VLT/UVES spectroscopy

Echelle spectra of HE 0515-4414 were obtained during commissioning of UVES at the VLT/Kueyen telescope. The observations were carried out under good seeing conditions (0.5-0.8arcsec) and a slit width of 0.8arcsec was used. A summary of the observations and of the detectors used is given in the ESO web pages http://www.hq.eso.org/instruments/uves.

The spectra were extracted using an algorithm that attempts to reduce the statistical noise to a minimum. After bias-subtracting and flat-fielding of the individual CCD frames, the seeing profiles were fitted with a Gaussian in two steps. In a first step the three parameters of the Gaussian - width, amplitude, and offset from the previously defined orders - were unconstrained; in the second step only the amplitudes were allowed to vary, with width and offset held fixed at values found by a $\kappa\sigma$-clipping fit along the dispersion direction to the values obtained in the first step. Flux values were assigned with a variance according to the Poisson statistics and the read-out noise, while cosmic-ray shots were assigned with infinite variances. Thus, the extraction procedure recovers the total count number even at wavelengths where the spatial profile is partially modified by cosmic-ray hits.

The extracted spectra were wavelength calibrated using as reference Th-Ar spectra taken after each science exposure. All wavelength solutions typically were accurate to better than 1/10 pixel. The wavelength values were converted to vacuum heliocentric values and each spectrum of a given instrumental configuration was binned onto a common linear wavelength scale (of typically 0.04 Å per pixel). Finally, the reduced spectra were added, weighting by the inverse of the flux variances.

2.3 Line profile analysis

Our analysis was carried out using a multiple line fit procedure to determine the parameters $\lambda_{\rm c}$ (line center wavelength), N (column density), and b (line broadening velocity) for each absorption component. We have written a FORTRAN program based on the Levenberg-Marquardt algorithm to solve this nonlinear regression problem (see, e.g., Bevington & Robinson 1992). We have included additional parameters describing the local continuum curvature by a low order Legendre polynomial. A free floating continuum is a prerequisite for an adequate profile decomposition in the case of complex line ensembles.

To improve the numerical efficiency we have to provide adequate initial parameters. In some cases the success of the fitting depends on good starting parameters, since the algorithm tends to converge to the nearest, not necessarily global, minimum of the chi-square merit function. A first approximation can be found neglecting the instrumental profile and converting the flux profile into apparent optical depths using the relation

\begin{displaymath}\tau_{\rm a}(\lambda) = \ln [F_{\rm c}(\lambda)/F_{\rm obs}(\lambda)]\,,
\end{displaymath} (1)

where $F_{\rm c}$ and $F_{\rm obs}$ are the continuum level and the observed line flux, respectively. If the instrumental resolution is high compared to the line width, $\tau_{\rm
a}(\lambda)$ will be a good representation of the true optical depth $\tau(\lambda)$. However, an ill-defined continuum level or saturation effects may produce large uncertainties. The apparent optical depth can be automatically fitted with a sum of Gaussians, each having a variable position, amplitude, and width. An a priori line identification is not necessary at this stage of the analysis.

Having obtained first-guess parameters we proceed with Doppler profile fitting using artificial test lines with z=0 and f=1, where f is the oscillator strength. It can be shown that most Voigt profiles are well represented by the purely velocity broadened Doppler core. The size of the fit region depends on the complexity and extent of the absorption line ensembles. Indeed, the number of free parameters should be less than 100 to preserve the numerical efficiency. One specific characteristic of our technique is the simultaneous continuum normalization which can reconstruct the true continuum level even in cases, where the background is hidden by numerous lines. The multi component profile is the convolution of the intrinsic spectrum and the instrumental spread function $P(\Delta \lambda)$:

\begin{displaymath}F(\lambda) = P(\Delta\lambda) \otimes \left\{ F_{\rm c}(\lamb...
\exp[-\tau_i(\lambda,\lambda_{\rm c},N,b)] \right\}\cdot
\end{displaymath} (2)

If the program fails to converge on a reasonable model, the parameters can be adjusted by hand. In this way the fit can be modified to be acceptable by eye and then re-minimized. In some exceptional cases this procedure is the only chance to free a converged solution from a local chi-square minimum.

After line identification the parameters of the test lines can be transformed to the actual redshift and oscillator strength. However, the contribution of unknown profile components can still be considered using the test line results. A final Voigt profile fit with all identified components includes the simultaneous multiplet treatment, keeping the redshift, column density, and line width the same during the chi-square minimization. The upper limit of the column densities of non-detected lines is estimated assuming a $5\,\sigma$significance level for the equivalent width.

3 Absorption systems with OVI lines

We searched for O VI lines associated with known Ly $\alpha$ and Ly $\beta$absorbers. Therefore, as a starting point, we tried to identify all Ly $\alpha$ lines. Line identification and the analysis of the Ly $\alpha$ forest will be presented in some detail in a later paper. At the resolution of $\sim$30000 (STIS) and $\sim$50000 (UVES), narrow metal lines can usually be distinguished easily from hydrogen lines. In all Ly $\alpha$ absorption systems with column densities log $N_{\rm H} \geq$ 13.5 we searched for metal lines, in particular for O VI, C IV, N V, Si IV, C III, N III. In a first step, all lines within $\pm 200\, {\rm km\,s^{-1}}$ were considered to be plausibly associated with the Ly $\alpha$/Ly $\beta$ systems. Within this selection criterium we have found 6 systems with probable O VI absorption, listed in Table 1. Due to the moderate S/N ratio of the STIS spectra (between 10 and 20 per resolution element) the detection limit of O VI is estimated to lie between $\log N = 13.3$ at the lower limit of the z range and $\log N = 13.1$ near the quasar.

\includegraphics*{ms1253f6r.eps} \end{figure} Figure 1: Selected absorption line profiles of systems with O VI detection. The normalized flux is plotted vs. rest-frame velocity of the hydrogen main component. Long tick marks indicate the position of the primary lines, while short tick marks indicate additional absorption components. It should be noted that some profile ensembles contain lines which do not belong to the same absorption system. The dotted curves represent our fit models.
Open with DEXTER

Table 1: Absorption line systems with O VI detections.

\hline {\rule[-2mm]{-1mm}{6mm}...
...m 0.04$ & $20.2 \pm 2.0$ & 42 & $-1$&2 \\

Table 1: continued.

\hline {\rule[-2mm]{-1mm}{6mm}...
... 0.05$ & $22.0 \pm 2.8$ & 32 & $-12$&2 \\

The Doppler parameter has been fixed to improve the goodness-of-fit.

4 The cosmological mass-density of the O VI phase

With the present STIS spectra of HE 0515-4414 the redshift range $1.21 \leq z \leq 1.73$ has been covered for the first time at sufficiently high resolution to undertake a sensitive search for O VI absorbers. We have detected 6 O VI systems. Two of them (z = 1.697, 1.736) are either associated with the QSO or in the proximity zone of the extremely luminous QSO. The system z = 1.416 is marginal, since only the 1031Å line is detected. Including the latter, we have 4 detections in the range z = 1.21to 1.67 which yield a number density of O VI absorbers with $W_{\rm rest} \geq$ 25mÅ of ${\rm d}N/{\rm d}z
\leq 10$. Compared with the findings by Tripp et al. (2000) of ${\rm
d}N/{\rm d}z = 48$ at $\bar{z} \simeq 0.21$, the number density at $\bar{z} \simeq 1.44$ is roughly a factor of 5 lower. Tripp et al. (2000) compared their finding of a high number density of weak O VI absorbers ( $W_{\rm rest} \geq$ 30 mÅ) in H 1821+643 and PG 0953+415 with other classes of absorbers and found that the weak O VI number density is more comparable to that of the low z weak Ly $\alpha$ absorbers - which have ${\rm d}N/{\rm d}z \approx 100$ for $W_{\rm rest} \geq$ 50mÅ - than to other types of metal absorbers like Mg II. In HE 0515-4414 we have at least 42 Ly $\alpha$ systems (the exact number being unknown due to the line blending problem) with $W_{\rm rest} \geq$ 50mÅ in the range $1.21 \leq z \leq
1.67$ which yields roughly ${\rm d}N/{\rm d}z = 90$[*]. Among these, roughly half of them are strong, saturated Ly $\alpha$ lines with a detected Ly $\beta$ line. Again, while our STIS spectrum of HE 0515-4414 confirms the number density of Ly $\alpha$ absorbers found previously (see Weymann et al. 1998), the number of O VI absorbers with $W_{\rm rest}
\geq 25$mÅ is lower than the number of Ly $\alpha$ absorbers with $W_{\rm rest} \geq 50$mÅ by a factor of 10. It is noteworthy that, except the z = 1.674 system, O VI is detected in lower column density Ly $\alpha$ absorbers (log $N_{\rm H} \leq 14$). Following the calculations by Tripp et al. (2000) and earlier work by Storrie-Lombardi et al. (1996) and Burles & Tytler (1996), the mean cosmological mass-density of O VI absorbers can be written in units of the critical density $\rho_{\rm c}$ as

\begin{displaymath}\Omega_{\rm b} (O{\sc vi}) = \frac{\mu\,m_{\rm
...}_{O{\sc vi}} \frac{\sum_i N_i ({\rm
O{\sc vi}})} {\Delta X},
\end{displaymath} (3)

where [O/H] is the assumed oxygen abundance in the O VI absorbing gas, f(O VI) is the fraction of oxygen in O VI, $\sum_{i} N_i(O{\sc vi})$ is the total O VI column density from all absorbers, and $\Delta X$ is the absorption distance (Bahcall & Peebles 1969).

Over the redshift interval z = 1.21 to z = 1.67 we have $\Delta X = 0.72$ for q0 = 1/2. $\sum_i N_i ({\rm
O{\sc vi}})$ is $2.1 \times 10^{14}\,{\rm cm^{-2}}$ (Table 1). Assuming $f(O{\sc vi}) = 0.2$, following Tripp et al. (2000) and Tripp & Savage (2000), which is close to the maximum for both collisional ionization and photoionization, we obtain a lower limit $\Omega_{\rm b}$(O VI) $\geq 3 \times 10^{-5}$ $[{\rm (O/H)/(O/H)}_{\odot}]^{-1}\,h^{-1}_{75}$. The only reliably measured heavy element abundances at $\bar{z} = 1.4$ are from DLAs. Typically the metal abundance (e.g. from Zn) is 1/10 solar (Pettini et al. 1999; Vladilo et al. 2000). There is, however, no guarantee that these abundances apply also to the O VI absorbers among the low column density systems. Assuming 1/10 solar for the oxygen abundance, we have $\Omega_{\rm b} ({\rm
O{\sc vi}}) \geq 3\times 10^{-4}\, h^{-1}_{75}$. With the same assumptions Tripp et al. (2000) derived a value $\geq 4\times
10^{-3}\, h^{-1}_{75}$. Using a somewhat different formalism for the derivation of $\Omega_{b}(O{\sc vi})$, namely Eq. (6) from Tripp & Savage (2000), we get with the same assumptions $\Omega_{\rm b}$(O VI) $\geq 1.5 \times 10^{-4}\,
h^{-1}_{75}$. Both from the number counts of the O VI systems and the estimate of the mean O VI density the unavoidable conclusion seems to be that at $\bar{z} = 1.5$, the baryon content of the O VI phase contains a factor of $\geq
10$ less material than at $\bar{z} = 0.21$.

5 Conclusions

Our results on O VI absorbing clouds in HE 0515-4414 can be summarized as follows:

The occurrence of an extremely broad component superimposed on the "normal" Doppler profile as observed in the z = 1.674 system is a rare Ly $\alpha$ profile type. In fact, we have never seen such a profile combination. In the context of modern interpretations of the Ly $\alpha$ forest as caused by a gradually varying density field characterized by a network of filaments and sheets (e.g. Bi & Davidsen 1997), a multi-component Voigt profile fitting is artificial and without a physical meaning anyway. The observed Ly $\alpha$ profile at z = 1.674 could be easily modelled by an overdense structure with inflow or outflow velocities of the order of $100\,{\rm km\,s^{-1}}$. We abstain from such an exercise, since the line profile decomposition would not lead to a unique solution. The coincidence with a strong O VI doublet at the same velocity is remarkable.

Our finding, that the O VI phase at $\bar{z} = 1.5$contains a factor of $\geq
10$ less material than at z = 0.21, provided the O VI/O ratio and the oxygen abundance are similar, appears to be inconsistent with the simulations of Davé et al. (2001) who predict an increase of the mass-fraction of baryons in the warm-hot phase of the IGM by at most a factor of 4 between z = 1.5 and 0.2. An increase in the mean oxygen abundance in the low density IGM by a factor of $\sim$3 over the same redshift range would restore consistency with the theoretical predictions. However, at present we do not see a possibility to test this hypothesis. Furthermore, as long as we do not understand the ionization to O VI quantitatively, the fractional ionization O VI/O might vary between z = 1.5 and 0.2. Finally, our result is still debatable due to small number statistics. More lines of sight, both at low and intermediate redshift, have to be probed.

This work has been supported by the Verbundforschung of the BMBF/DLR under Grant No. 50 OR 99111. S.L. acknowledges financial support by FONDECYT grant N $^{\rm o}
3\,000\,001$ and by the Deutsche Zentralstelle für Arbeitsvermittlung.



Copyright ESO 2001