Astronomy & Astrophysics manuscript no. echopap c© ESO 2021
November 8, 2021
Line and continuum variability of two intermediate-redshift,
high-luminosity quasars
D.Trevese1, D.Paris1, G. M. Stirpe2, F.Vagnetti3, and V. Zitelli2
1 Dipartimento di Fisica, Università di Roma “La Sapienza”, P.le A.Moro 2, I-00185 Roma(Italy)
e-mail: dario.trevese@roma1.infn.it
2 INAF - Osservatorio Astronomico di Bologna, via Ranzani, 1 - 40127 Bologna (Italy)
3 Dipartimento di Fisica, Università di Roma “Tor Vergata”, Via della Ricerca Scientifica, 1, I-00133 Roma (Italy)
ABSTRACT
Context. It has been shown that the luminosity of active galactic nuclei and the size of their broad line region obey a simple relation of
the type RBLR = aLγ, from faint Seyfert nuclei to bright quasars, allowing single-epoch determination of the central black hole mass
MBH = bLγ∆2Hβ from their luminosity L and width of Hβ emission line. Adopting this mass determination for cosmological studies
requires the extrapolation to high redshift and luminosity of a relation whose calibration, relies so far on reverberation mapping mea-
surements performed for L <∼1046erg s−1 and redshift z <∼ 0.4.
Aims. We initiated a campaign for the spectrophotometric monitoring of a few luminous, intermediate redshift quasars whose apparent
magnitude, V < 15.7, allows observations with a 1.8 m telescope, aimed at proving that emission lines vary and respond to continuum
variations even for luminosities >∼1047erg s−1, and determining eventually their MBH from reverberation mapping.
Methods. We have repeatedly performed simultaneous spectrophotometric observations of quasars and reference stars to determine
relative variability of continuum and emission lines. We describe the observations and methods of analysis.
Results. For the quasars PG 1634+706 and PG 1247+268 we obtain light-curves respectively for CIII](λλ1909Å), MgII(λλ2798Å)
and for CIV(λλ1549Å), CIII](λλ1909Å) emission lines with the relevant continua. During 3.2 years of observation, in the former case
no continuum variability has been detected and the evidence for line variability is marginal, while in the latter case both continuum
and line variability are detected with high significance and the line variations appear correlated with continuum variations.
Conclusions. The detection of the emission line variability in a quasar with L ∼1047erg s−1 encourages the prosecution of the moni-
toring campaign which should provide a black hole mass estimate in other 5-6 years, constraining the mass-luminosity relation in a
poorly explored range of luminosity.
Key words. galaxies:active – quasars: emission lines – quasars: general – quasar:individual:PG 1634+706, PG 1247+268
1. Introduction
Supermassive black holes (SMBHs) are believed to inhabit most,
if not all, the bulges of present-epoch galaxies (Kormendy &
Richstone, 1995), and strong evidences exist of a correlation be-
tween the black hole mass and either the mass Mbulge and lu-
minosity (Marconi et al., 2003, and refs. therein) or the veloc-
ity dispersion σ∗ of the host bulge (Ferrarese & Merrit, 2000;
Tremaine et al., 2002). This strongly suggests that the formation
and growth of SMBHs and galaxies are physically related pro-
cesses and provides a basis for a theory of cosmic structure for-
mation, incorporating the feedback from Active Galactic Nuclei
(AGNs) (Silk & Rees, 1998; Vittorini, Shankar, & Cavaliere,
2005, and refs.therein). Black hole masses determinations based
on stellar or gas kinematics are intrinsically limited, by angu-
lar resolution, to relatively nearby objects and cannot be applied
to bright AGNs where the nuclear light prevails over the galac-
tic component, just in the central region where the galactic gas
or star motion is dominated by the black hole gravitational field.
The reverberation-mapping technique does not suffer of this lim-
itation and represents the only mean to measure the mass of
SMBH in bright AGNs.
Emission lines, in the optical-UV region, are interpreted as
recombination of a gas which is photoionized by the continuum
radiation emitted by the inner region of the nucleus, presumably
an accretion disk surrounding the black hole. Emission lines re-
spond to variation of the ionizing continuum. Although the phys-
ical origin of these variations is poorly known (e.g. Trevese &
Vagnetti, 2002; Vanden Berk et al., 2004; De Vries et al., 2005)
it is possible to use the response of lines to continuum varia-
tions to investigate the structure of the line emitting region. This
requires long campaigns of accurate spectrophotomeric monitor-
ing of AGNs, which have led in the past to major progresses to-
wards understanding the physics of the “atmosphere” of Seyfert
1 galaxies. A summary of these results is given in Peterson
(1993).
Line widths, e.g. the FWHM ∆Hβ of the Hβ emission line,
correspond to r.m.s. velocities of the emitting gas clouds. A
cross-correlation analysis of continuum and emission-line light-
curves, evidencing a time delay τ of line respect to continuum
variations, allows to estimate the size R = τ/c of the region
where the line photons are generated. If the gas motion in the
Broad Line Region (BLR) is dominated by gravitation (Peterson
& Wandel, 2000), the size estimate RBLR can be combined with
the line width to yield a primary estimate of the virial mass of
the black hole MBH ∝ ∆2Hβ/GRBLR and the relevant Eddington
ratio. For Seyfert 1 galaxies typical BLR sizes are of the order of
light-days to light-weeks. Similar studies are more difficult for
quasars (QSOs) which require a longer monitoring.
2 D. Trevese et al.: Line and continuum variability
A long term campaign for a subsample of 28 QSOs was
started in 1991 with the Wise 1.0m and the Steward 2.3m tele-
scopes (Maoz et al., 1994). As a result, nine years later Kaspi
et al. (2000) provided mass estimates for the entire sample. The
new data, combined with previous results on Seyfert 1 galax-
ies, thus spanning a much wider range of intrinsic luminosity,
allowed to establish an average relation between the intrinsic lu-
minosity and the size, RBLR = aLγ, with γ ' 0.7, which allows a
secondary estimate of the black hole mass based on single-epoch
observations of luminosity and line width: MBH = bLγ∆2Hβ ,
where both the constant b and γ are determined statistically from
the available echo-mapping data. Recent studies show that γ is
in the range 0.5-0.7, depending on how luminosity is defined,
which lines are selected for the echo-mapping and the fitting pro-
cedure adopted (Kaspi et al., 2005; Bentz et al., 2006).
The extreme importance of secondary mass estimates relies
on the fact that on the basis of single epoch observations it is
possible to study the evolution in cosmic time of the mass distri-
bution of QSOs/AGNs, and to extend the studies of the relation
existing between QSOs and host bulges properties. However,
the above correlations with primary masses, based on echo-
mapping, were established for relatively close and faint AGNs
with z ≤ 0.4 and [λLλ(5100Å) <∼ 1046 erg s−1], thus it is
presently unknown whether they can be extrapolated to higher
luminosities and/or redshifts. For instance, the extrapolation of
the MBH − L relation (Kaspi et al., 2000), together with the
assumption that the known MBH − Mbulge − σ∗ relations holds
(Tremaine et al., 2002) , leads to predict the existence of galax-
ies with Mbulge ∼ 1013.1 − 1013.4M� and σ∗ exceeding 800 km
s−1. Such galaxies have never been observed, and their exis-
tence would put important constraints on galaxy formation mod-
els (Netzer, 2003). Therefore it is essential to extend the primary
mass measures to higher redshifts and luminosities. On the other
hand, for high QSO luminosities a large size of the broad line
region is expected. This would cause both a smoothing of the
line light-curve and larger time delays with respect to contin-
uum variations (Wilhite et al., 2005), thus the very detectability
and the amplitude of line variations are open questions.
A sample of objects with redshifts in the range 2 < z < 3.4
and apparent magnitude as faint as mV ∼18 is being monitored
by Kaspi et al. (2004) with the 9m Hobby-Eberly Telescope
(HET; Ramsey et al., 1998) and new results have been presented
recently (Kaspi et al., 2006). During their 6-year monitoring of 6
QSOs, significant continuum and emission-line variations were
detected in all targets and a preliminary black hole mass estima-
tion is given for one of them.
In the present paper we describe a new monitoring campaign
limited to objects with V < 15.7 and 1 < z < 4 which, thanks
to their apparent brightness, can be observed with the medium-
small 1.82 m Copernicus Telescope at Cima Ekar (Asiago, Italy),
through a service mode scheduling of a long term monitoring,
and allow to investigate whether: i) echo-mapping is feasible
for objects as bright as λLλ(5100Å) ∼ 1047 erg s−1 and ii) the
RBLR-luminosity correlation can be extrapolated to such bright-
ness. The paper is organized as follows. Section 2 describes the
sample, observations and the data reduction procedure. Section
3 describes the results for two quasars of the sample. Section 4
summarizes the results and discuss future prospects. In the fol-
lowing we derive λLλ(5100Å) from the flux in the Johnson V
band, extrapolating the flux density to the rest-frame λ = 5100Å
with a power law fν ∼ ν−0.5, and assuming a standard cosmology
Ho = 70 km s−1 Mpc−1, ΩM = 0.3, and ΩΛ = 0.7.
Table 1. The quasars monitored
Object z V log[λLλ(5100Å)]
[erg s−1]
APM 08279+5255 3.911 15.20 47.7
PG 1247+268 2.042 15.60 47.0
PG 1634+706 1.337 15.27 46.7
HS 2154+2228 1.290 15.30 46.7
4000 5000 6000 7000
Fig. 1. Average spectra of PG 1634+706 (upper panel) and PG
1247+268 (lower panel). Spectral ranges for continuum deter-
mination (short ticks) and line flux evaluation (long ticks), as
reported in Table 2, are shown.
2. Observations
2.1. Object selection
The sample has been extracted from the Veron-Cetty & Veron
(2003) (11th ed.) catalog with the condition δ > 0, V < 15.7
mag and z > 1 in order to select objects of bright enough intrinsic
luminosity to investigate the bright end extension of the RBLR vs.
λLλ(5100Å) relation (Kaspi et al., 2000, 2005). These conditions
identify 12 objects, only four of which were monitored, owing
to the limits on observing time. These four objects are listed in
Table 1.
Observations at intermediate redshift allow to sample the
variability of MgII λ2798, CIII] λ1909, CIV λ1559 lines, instead
of Hα, Hβ, Hγ observed in the low redshift study of Kaspi et al.
(2000). This allows to study either BLR at smaller distance from
the center, or regions of the same sizes of those producing the
Balmer lines, but using lines which respond to different part of
the continuum spectrum. The main emission lines falling within
the observed wavelength interval are indicated for two objects
in Table 2. Figure 1 shows the average spectra of PG 1634+706
and PG 1247+268.
With respect to other QSO monitoring programmes, ours
is the first which includes the MgII line in part of the ob-
served sources. This line is particularly significant, because i)
past Seyfert 1 monitoring campaigns conducted with IUE have
D. Trevese et al.: Line and continuum variability 3
shown its lag to be similar to that of Hβ (Clavel et al., 1991;
Reichert et al., 1994); ii) this line is most often used to derive es-
timates of black-hole mass from single-epoch spectra of high-z
QSOs (McLure & Jarvis, 2002), because its width is tightly cor-
related to that of Hβ. Deriving a lag for the MgII line would
therefore allow to estimate the black-hole mass most consis-
tently with respect to the results of Seyferts, which are mostly
based on the monitoring of Hβ.
APM 08279+5255 is one of the most luminous known QSOs
if its emission is considered isotropic. However it has been
shown to be lensed by a foreground galaxy (Irwin et al., 1998).
Three components, separated by a few tenths of an arcsec have
been detected in near-infrared images obtained with Keck tele-
scope and different models of the lensing field predict a few
days delay between photometric variations of the components
(Ledoux et al., 1998; Ibata et al., 1999; Egami et al., 2000):
a short time compared with the expected time scale of intrin-
sic variations. APM 08279+5255 is also a Broad Absorption
Line QSO (Irwin et al., 1998), which makes more difficult to
define regions free from either emission or absorption features
to measure continuum variations. The analysis of this object is
deferred to a forthcoming paper. HS 2154+2228 has been ob-
served so far only 5 times and the analysis requires further mon-
itoring. The other two objects PG 1247+268 and PG1634+706
are analyzed in this paper with the aim of : i) verifying the ad-
equacy of the observational data and reduction procedures, un-
der the assumption that variability amplitudes and characteristic
time scales can be extrapolated from the properties of fainter
objects; ii) possibly detecting line variations in objects as lumi-
nous as λLλ(5100Å) ∼ 1047 erg s−1, and compare their amplitude
with continuum variations. For both objects a star of comparable
magnitude, as close as possible to the QSO, has been selected
for the relative spectrophotometric calibration, described in the
next section.
2.2. Spectrophotometric Observations and Data Reduction
Observations were carried out at Asiago 1.82 m telescope
equipped with the Faint Object Spectrograph & Camera AFOSC
which is a focal reducer with reduction factor of 0.58, designed
to allow a quick switching between spectroscopic and imaging
modes. The scale at the focal plane is 21.7 ′′/mm. The detector
is a 1024x1024 thinned CCD array TEK1024 with 22x22 µm2
pixels corresponding to a scale of 0.473 arcsec pixel−1 and a
FOV of 8.14x8.14 arcmin2. We adopted a 8”.44-wide slit and
a grism with a dispersion of 4.99 Å pixel−1, providing a typ-
ical resolution of ∼ 15 Å in the spectral range 3500-8450 Å.
Spectrophotometric exposures are performed after orienting the
slit to include both the QSO and the reference star of comparable
magnitude, located at (12:50:11.5 +26:33:32) and (16:34:57.4
+70:32:49) (J2000) for PG 1247+268 and PG 1647+706 respec-
tively. The reference stars are included as internal calibrators for
QSO spectra, as described by Maoz et al. (1990) and Netzer et
al. (1990). The wide slit is necessary to avoid different fractional
losses of the QSO and star light due to possible non perfect slit
alignment and differential refraction, which could cause spuri-
ous variation of the flux ratios. Lamp flats and Hg-Cd ark spectra
were also taken for wavelength calibration. At each epoch, typi-
cal science observations consist of two consecutive exposures of
1800 s.
QSO and star spectra,Q(λ) and S (λ) were extracted with the
standard IRAF procedures. The QSO/star ratio as a function of
wavelength is computed for each exposure k = 1, 2
µ(k)(λ) = Q(k)(λ)/S (k)(λ). (1)
This quantity is independent of extinction changes during the
night. This allows a check of consistency between the two expo-
sures and the rejection of the data if inconsistencies occur (under
the assumption that QSO variations are negligible on 1 h time
scale). In fact, typical spectra of two consecutive exposures have
a ratio µ(1)(λ)/µ(2)(λ) of order unity, with deviations smaller than
0.02 when averaged over 500 Å, at least in the 4000Å − 7000Å
range. When discrepancies are larger than 0.04 both exposures
are rejected.
It is important to notice that, because of the large slit width,
there are small changes in the λ scale due to changes of the ob-
ject position within the slit (which are in general negligible in
the case of pairs of consecutive exposures). This suggested to
proceed as follows. As a first step we register the zero point of
the λ scale. However, since the changes of position within the slit
do not correspond exactly to rigid shifts of the λ scale, we select
portions of the spectra in wavelength intervals |∆λ| <∼ 1000Å
around each of the QSO emission lines considered (see Table 2),
and we determine the shifts for each portion.
Once the wavelength scales are registered, the QSO and star
spectra taken in the two exposures are co-added and the ratio
µi(λ) = (Q
i + Q
i )/(S
i + S
i ) (2)
is computed, at each epoch i. The reference star is flux calibrated
at a reference epoch. Since our aim is to compute relative flux
variations, we reduce all quasar spectra to this reference epoch,
by multiplying all the µi(λ) by the flux calibrated star spectrum
f S (λ): f Qi (λ) ≡ µi(λ) f
S (λ). We stress that we are not interested
in the absolute flux calibration, whose accuracy is of the order
of 20% and which is applied for the sole purpose reporting the
QSO spectra in physical units. The star spectrum f S (λ) adopted
is the same for all epochs, thus it does not affect relative flux
changes we want to measure.
The QSO spectrum f Q(λ) = c(λ) + l(λ) can be decom-
posed in a line l(λ) and continuum c(λ) spectra. Two values
of the continuum, cshort and clong at shorter and longer wave-
length with respect to the most prominent QSO emission lines,
are evaluated in regions as free as possible of other emission
features, defined by the wavelengths ranges (λshort,1,λshort,2) and
(λlong,1,λlong,2) respectively. Table 2 reports the observer-frame
wavelengths defining the continuum regions and the intervals for
line integration. These were chosen on the basis of the analysis
of Clavel et al. (1991); Reichert et al. (1994) of UV spectra of
low redshift AGNs observed with the International Ultraviolet
Explorer (IUE), with slight modifications to maximize the S/N
ratio (clong of MgII(λ2798 Å) falls outside the wavelength range
covered by the IUE spectrograph).
Line fluxes are computed as:
[F(Q)(λ) − cint(λ)]dλ, (3)
where cint(λ) is the linear interpolation through cshort and clong.
The extremes λ1 and λ2, also listed in Table 2, not necessarily
coincide with λshort ≡ (λshort,1 + λshort,2)/2 and λlong ≡ (λlong,1 +
λlong,2)/2, and are chosen to optimize the fl signal to noise ratio.
Direct B and R images in 15x15 arcmin fields centered on the
QSOs were taken to check possible variability of the reference
4 D. Trevese et al.: Line and continuum variability
Table 2. Wavelength intervals for lines and continua [Å]
Object z line λshort,1 λshort,2 λ1 λ2 λlong,1 λlong,2
PG1634+706 1.337 CIII]λ1909 4252 4298 4382 4590 4630 4674
PG1634+706 1.337 MgIIλ2798 6215 6285 6251 6777 6755 6815
PG1247+268 2.042 CIVλ1549 4380 4502 4598 4801 5202 5262
PG1247+268 2.042 CIII]λ1909 5535 5595 5770 5900 6025 6085
CIII]
2800 3200 3600 4000
t (JD-2450000)
2800 3200 3600 4000
t (JD-2450000)
Fig. 2. Light-curves of PG 1634+706 in the observer’s frame,
as relative flux variations ∆ f / f . Upper panels: emission lines,
middle panel: shorter wavelength continuum cshort, lower panel:
longer wavelength continuum clong, left: CIII]/λ1909 Å, right:
MgII λ2798 Å.
star. Comparison of the reference star with the brightest star in
the field showed a r.m.s. fractional flux variation of 0.02 in both
cases. Thus we can assume that the reference star is not variable,
at this level of accuracy.
Relative-flux calibration 1-σ errors, reported in figures 2 and
3, are estimated by adding in quadrature the r.m.s. uncertainty
on the flux of the reference star to the r.m.s. fractional uncertain-
ties of line or continuum fluxes, which, in turn, are computed by
comparing pairs of exposures which are added to form quasar
spectra of individual epochs.
3. Results and discussion
Figures 2 and 3 show the light-curves in the observer’s frame, for
emission lines and continua listed in Table 2, for PG1634+706
and PG1247+268 respectively. The light curves are expressed
in terms of the relative flux variations ∆ f / f , with respect to the
flux f at the reference epoch where the absolute calibration was
performed (see Section 2.2). The total time base is 3.3 years in
the observer’s frame, corresponding in the rest-frame to 1.4 yr
for the former QSOs and 1.1 yr for the latter. In the case of
PG1634+706 no significant variations are detected: neither in
CIII] line, which has the lowest S/N ratio due to its intrinsic
faintness, nor in the MgII line. The relevant continua are also
2800 3200 3600 4000
t (JD-2450000)
CIII]
2800 3200 3600 4000
t (JD-2450000)
Fig. 3. Light-curves of PG 1247+268 in the observer’s frame,
as relative flux variations ∆ f / f . Upper panels: emission lines,
middle panel: shorter wavelength continuum cshort, lower panel:
longer wavelength continuum clong, left: CIV/λ1549 Å, right:
CIII] λ1909 Å.
consistent with no variability at a level of ∆ f / f ∼ 0.02-0.03
r.m.s..
The result is different in the case of PG 1247+268. The
faintest line CIII] λ1909 Å shows a marginal evidence of vari-
ability and the stronger line CIV λ1549 Å seems to decrease
steadily during the observing period.
An almost steady decrease does not allow to derive possi-
ble line-continuum time delays from cross-correlation analysis.
However we can establish quantitatively the evidence of contin-
uum and, most important, of line variability. For this purpose we
define the unbinned discrete structure function
UDS F(τi j) =
|y(ti) − y(t j)|, (4)
where y(t) represents any of the line or continuum light-curves
considered, ti and t j are two observation epochs and τi j = ti − t j
is the time delay. The (binned) structure function can be defined
in bins of time delay, centered at τ:
S (τ) =
∑
UDS F(τi j)
 , (5)
where the sum is extended to all the M values of UDSF belong-
ing to a given bin of τ. We adopt the average of UDSF, instead
D. Trevese et al.: Line and continuum variability 5
0 200 400 600 800 1000
0 200 400 600 800 1000
Fig. 4. Structure functions for PG 1634+706 in the observer’s
frame. Upper panels: emission lines, middle panel: shorter wave-
length continuum cshort, lower panel: longer wavelength contin-
uum clong, left: CIII]/λ1909 Å, right: MgII λ2798 Å.
of UDSF2, since it is more stable, i.e. less sensitive to deviant
points; the factor
2 in equation (4) makes S (τ) equal to the
standard deviation in the case of a Gaussian distribution (see Di
Clemente et al., 1996). Structure functions for lines and continua
0 200 400 600 800 1000
0 200 400 600 800 1000
Fig. 5. Structure functions for PG 1247+268.Upper panels:
emission lines, middle panel: shorter wavelength continuum
cshort, lower panel: longer wavelength continuum clong, left:
CIV/λ1549 Å, right: CIII] λ1909 Å.
of the two quasars PG 1634+706 and PG 1247+268 are reported
in Figures 4 and 5 respectively, as a function of the time delay
in the observer’s frame. The error bars reported in the figures
represent simply the standard deviation of the UDSF in each bin
and not the uncertainty on their average value. The latter would
be 1/(M − 1) times the standard deviation, for M uncorrelated
UDSF values.
In the case of PG 1634+706 the structure function analy-
sis simply confirms what already appears from the light-curves,
namely no significant variations of lines and continua are de-
tected. The r.m.s. noise level is about 0.02-0.03 for continua and
0.05 for lines.
On the contrary, in the case of PG 1247+268, there is a clear
increase of variability for long time lags. To establish the signifi-
cance of this variability it is necessary to evaluate the probability
of the null hypothesis that a value of S (τ) is produced by pure
noise. For this purpose we generated mock noise light-curves ,
n(ti), i=1,Nepo, where Nepo is the number of observing epochs,
by extracting random numbers with a Gaussian distribution of
standard deviation σ, which represents pure noise, under the as-
sumption that the photometric noise is not correlated on the time
scale of our minimum sampling interval (namely about 10 days).
Then, the relevant structure function S mock(τ) has been computed
with the same binning adopted for real data. The simulation was
iterated for 107 times to generate, for each bin, the statistical dis-
tribution of the S mock(τ) values.
The value adopted for σ has been estimated, conservatively,
from the observed structure functions of figure 4 and 5 them-
selves, as the value of S (τ) in the first bin. This is, in fact, an
overestimate of noise, since it includes short-time-scale varia-
tions of the QSO, which however should be negligible on the
basis of previous studies of (fainter) QSOs (Giveon et al., 1999;
Kaspi et al., 2000). The almost constant value of the structure
functions in the case of PG 1634+706 confirms the adopted hy-
pothesis that noise is not correlated on long time scale.
For the CIII] line the r.m.s. noise (∼ 0.05) is larger than in
the case of the CIV line (∼ 0.02). This depends on the intrinsic
faintness of the CIII] line. According to the above simulations
the probability of the null hypothesis, for the second and fourth
bins, is P(> S ) ∼ 4 × 10−4. The local maximum of S (τ), for
CIII] at τ ∼400 days, is due to an “oscillation” of the light-curve
on a time scale comparable with the total time base of the ob-
servations. This indicates that a longer time base is needed to
properly average individual oscillations and to obtain a “station-
ary” structure function. Further sampling, trying to avoid yearly
periodicity due to periods of optimal observability, is needed to
obtain the resolution required for time delay measurements.
The structure function of CIV line shows a much more sig-
nificant variability: the probability of the null hypothesis in the
case of the last two bins, where S (τ) is 0.13 and 0.21 respec-
tively, is less than 10−6. The systematic decrease of CIV line
luminosity and the relevant continua might seem to rely on the
value of the light-curves at the first epoch. Thus we have recom-
puted the S (τ) after removing the first epoch. Indeed, in the case
of continua the systematic decrease of S (τ) is no longer evident.
On the contrary, for the line it remains practically unchanged
(except that the last bin is no longer present). Thus the fact that
the probability of the null hypothesis is smaller than 10−6 does
not rely on that particular point in the light-curve.
The main goal of new echo-mapping campaigns is to obtain
a direct measure of the central black hole mass and to establish
whether the scaling of QSO broad line region size and masses re-
main the same at higher quasar luminosity. The very detectabil-
ity of line variation in the most luminous QSOs was, until re-
6 D. Trevese et al.: Line and continuum variability
cently, an open question, since 3C273 with (λLλ(5100Å) <∼ 1046
erg s−1) was the brightest QSO with detected line variability
(Kaspi et al., 2005). An evidence of line variability in brighter
objects has been reported more recently by Kaspi et al. (2006).
The present detection of line variability in one of the two quasars
considered strengthens the notion that quasar emission lines re-
spond to continuum changes in high luminosity QSOs like in
lower luminosity ones and indicates that echo-mapping cam-
paigns are feasible for λLλ(5100Å) >∼ 1047 erg s−1.
Our results are consistent with those of Kaspi et al. (2006),
who have monitored sources of comparable luminosities and
redshifts. We have also detected variability, of comparable am-
plitude, in the continuum and CIII] and CIV emission lines in
one of our sources. Considering the shorter time baseline of
our observations, the fact that variability has not been found in
PG 1634+706 is consistent with the segments without variations
which can be seen in the Kaspi et al. (2006) light curves on this
time scale.
To estimate the further monitoring needed to measure line-
continuum delays, we should take into account that, on aver-
age, the amplitude of variability increases with the rest-frame
time lag. However this dependency is weak: log S (τrest) ∼
0.15 log τrest, according to previous statistical studies on large
quasar samples (Vanden Berk et al., 2004; De Vries et al., 2005).
In the case of our quasars PG 1634+706 and PG 1247+268, the
observer’s frame τobs = τrest(1+z) is dilated by a factor 2.3 and 3
respectively. The size of the broad line region, based on Balmer
line reverberation-mapping, is expected to scale with luminosity
as RBLR ∝ Lγ with γ=0.5-0.7, depending on how L is measured
and the RBLR vs. L fit is obtained (Kaspi et al., 2000, 2005; Bentz
et al., 2006; Kaspi et al., 2006). A Balmer line BLR of about 3
light years in the rest frame is expected for the relevant luminosi-
ties. On the other the hand CIV line corresponds to a BLR size
2 times smaller than Balmer lines (Peterson & Wandel, 2000).
Considering that in general a time baseline of about twice the
light-crossing time is required in order to obtain a reliable lag
from the light curves, we estimate that reverberation mapping
should be feasible after 5-6 more years of monitoring.
4. Summary
– We have initiated a campaign for the monitoring of 4 high
luminosity QSOs and we present the result for two of them
on the basis of 3.3 years of observation with the 1.82 m tele-
scope of the Asiago Observatory.
– We discuss the data reduction, the procedures adopted and
the level of accuracy attained in relative spectrophotometric
variability measurements.
– We perform a structure function analysis of the light-curve
and numerical simulations to establish the confidence level
of line variability detection.
– We detect line and continuum variability in one of the two
QSOs, PG 1247+268 of λLλ(5100Å) = 1047 erg s−1, with a
probability of less than 10−6 of the null hypothesis that the
observed structure function is produced by pure noise.
– This detection supports the notion that emission lines re-
spond to continuum variations as in substantially less lumi-
nous QSOs.
The results encourage the prosecution of the campaign which
should provide time delay, and black hole mass, estimates in 5-6
years.
Acknowledgements. We acknowledge support from the Asiago Observatory
team, in particular Hripsime Navasardian, for observations.
References
Clavel,J., Reichert, G. A., Alloin,D. et al., 1991, ApJ, 366, 64
Bentz, M. C., Peterson, B. M., Pogge, R. W., Vestergaard, M., & Onken, C. A.,
2006, ApJ, 644,133
De Vries, W.H., Becker, R. H., White, R.L., & Loomis, C., 2005, ApJ, 129, 615
di Clemente, A., Giallongo, E., Natali, G., Trevese, D., & Vagnetti, F., 1996,
ApJ, 463, 466
Egami,E., Matthews, K., Ressler, M., et al., 2000, ApJ, 535, 561
Ferrarese, L., & Merrit, D. 2000, ApJ, 539, L9
Giveon, Uriel, Maoz, Dan, Kaspi, Shai, Netzer, Hagai, Smith, Paul S., 1999,
MNRAS, 306, 637
Ibata, R. A., Lewis, M.J., Irwin, M. J., Lehar, J., & Totten, E. J., 1999, ApJ,
118,1922
Irwin, M. J., Ibata, R. A., Lewis, G. F., & Totten, E. J., 1998, ApJ, 505, 529
Kaspi, S., Smith, P.S., Meters, H., Maoz, D., Jannuzi, B.T. &Giveon, U., 2000,
ApJ, 533, 631
Kaspi, S., Netzer, H., Maoz, D., Shemmer, O., Brandt, W. N.,& Schneider, D. P.,
2004, in Coevolution of Black Holes and Galaxies, Carnegie Obs. Astroph.
Ser. Vol. 1, L. C. Ho ed., http://www.ociw.edu/ociw/symposia/series/
Kaspi, S., Maoz, D., Netzer, H., Peterson, B. M., Vestergaard, M., & Jannuzi, B.
T., 2005, ApJ, 629, 61
Kaspi, S., Bradt, W. N., Maoz, D., Netzer, H., Schneider, D. P., & Shemmer, O.,
2006, ApJ (in press), astro-ph/0612722
Kormendy, J., & Richstone, D. 1995, ARA&A,33, 581
Ledoux, C., Theodore, B., Petitjean, P., et al., 1998, A&A,339,L77
Maoz, D., Netzer, H., Leibowitz, E., et al. 1990, ApJ, 351, 75
Maoz, D., Smith, P. S., Jannuzi, B. T., Kaspi, S., & Netzer, H., 1994, ApJ, 421,
Marconi, A., Hunt, L. K., 2003, ApJ, 589L, 21
McLure, R. J., & Jarvis, M. J., 2002, MNRAS, 337, 109
Netzer, H., Maoz, D., Laor, A. et al., 1990, ApJ, 353, 108
Netzer, H., 2003, ApJ, 583, L5
Peterson, B. M., 1993, PASP, 105, 247
Peterson, B. M., & Wandel, A., 2000, ApJ, 540, L13
Silk, J. & Rees, M.J. 1998, A&A, 331, L1
Ramsey, L. W., et al. 1998, SPIE, 3352, 34
Reichert, G. A., Rodriguez-Pascual, P. M., Alloin, D. et al., 1994, ApJ, 425, 582
Tremaine, S. et al., 2002, ApJ, 574, 740
Trevese, D., & Vagnetti, F., 2002, ApJ, 564, 624
Vanden Berk, D, 2004, ApJ, 601, 692
Veron-Cetty M.P., Veron P., 2003, A&A, 412, 399
Vittorini, V., Shankar, F., & Cavaliere, A., 2005, MNRAS, 363, 1376
Wilhite, B. C.,Vanden Berk, D. E., Kron, R. G. et al., 2005, ApJ, 633, 638
http://www.ociw.edu/ociw/symposia/series/
http://arxiv.org/abs/astro-ph/0612722
	Introduction
	Observations
	Object selection
	Spectrophotometric Observations and Data Reduction
	Results and discussion
	Summary

  It has been shown that the luminosity of AGNs and the size of their broad
line region obey a simple relation of the type R=a L^g, from faint Seyfert
nuclei to bright quasars, allowing single-epoch determination of the central
black hole mass M=b L^g D^2 from their luminosity L and width of H_beta
emission line. Adopting this mass determination for cosmological studies
requires the extrapolation to high z and L of a relation whose calibration
relies so far on reverberation mapping measurements performed for L<10^46 erg/s
and z<0.4. We initiated a campaign for the monitoring of a few luminous,
intermediate z quasars whose apparent magnitude V<15.7 allows observations with
a 1.8m telescope, aimed at proving that emission lines vary and respond to
continuum variations even for luminosities >10^47 erg/s, and determining
eventually their M_BH from reverberation mapping. We have repeatedly performed
simultaneous observations of quasars and reference stars to determine relative
variability of continuum and emission lines. We describe the observations and
methods of analysis. For the quasars PG1634+706 and PG1247+268 we obtain
light-curves respectively for CIII], MgII and for CIV, CIII] emission lines
with the relevant continua. During 3.2 years of observation, in the former case
no continuum variability has been detected and the evidence for line
variability is marginal, while in the latter case both continuum and line
variability are detected with high significance and the line variations appear
correlated with continuum variations. The detection of the emission line
variability in a quasar with L~10^47 erg/s encourages the prosecution of the
campaign which should provide a black hole mass estimate in other 5-6 years,
constraining the M_BH-L relation in a poorly explored range of luminosity.

Introduction
Supermassive black holes (SMBHs) are believed to inhabit most,
if not all, the bulges of present-epoch galaxies (Kormendy &
Richstone, 1995), and strong evidences exist of a correlation be-
tween the black hole mass and either the mass Mbulge and lu-
minosity (Marconi et al., 2003, and refs. therein) or the veloc-
ity dispersion σ∗ of the host bulge (Ferrarese & Merrit, 2000;
Tremaine et al., 2002). This strongly suggests that the formation
and growth of SMBHs and galaxies are physically related pro-
cesses and provides a basis for a theory of cosmic structure for-
mation, incorporating the feedback from Active Galactic Nuclei
(AGNs) (Silk & Rees, 1998; Vittorini, Shankar, & Cavaliere,
2005, and refs.therein). Black hole masses determinations based
on stellar or gas kinematics are intrinsically limited, by angu-
lar resolution, to relatively nearby objects and cannot be applied
to bright AGNs where the nuclear light prevails over the galac-
tic component, just in the central region where the galactic gas
or star motion is dominated by the black hole gravitational field.
The reverberation-mapping technique does not suffer of this lim-
itation and represents the only mean to measure the mass of
SMBH in bright AGNs.
Emission lines, in the optical-UV region, are interpreted as
recombination of a gas which is photoionized by the continuum
radiation emitted by the inner region of the nucleus, presumably
an accretion disk surrounding the black hole. Emission lines re-
spond to variation of the ionizing continuum. Although the phys-
ical origin of these variations is poorly known (e.g. Trevese &
Vagnetti, 2002; Vanden Berk et al., 2004; De Vries et al., 2005)
it is possible to use the response of lines to continuum varia-
tions to investigate the structure of the line emitting region. This
requires long campaigns of accurate spectrophotomeric monitor-
ing of AGNs, which have led in the past to major progresses to-
wards understanding the physics of the “atmosphere” of Seyfert
1 galaxies. A summary of these results is given in Peterson
(1993).
Line widths, e.g. the FWHM ∆Hβ of the Hβ emission line,
correspond to r.m.s. velocities of the emitting gas clouds. A
cross-correlation analysis of continuum and emission-line light-
curves, evidencing a time delay τ of line respect to continuum
variations, allows to estimate the size R = τ/c of the region
where the line photons are generated. If the gas motion in the
Broad Line Region (BLR) is dominated by gravitation (Peterson
& Wandel, 2000), the size estimate RBLR can be combined with
the line width to yield a primary estimate of the virial mass of
the black hole MBH ∝ ∆2Hβ/GRBLR and the relevant Eddington
ratio. For Seyfert 1 galaxies typical BLR sizes are of the order of
light-days to light-weeks. Similar studies are more difficult for
quasars (QSOs) which require a longer monitoring.
2 D. Trevese et al.: Line and continuum variability
A long term campaign for a subsample of 28 QSOs was
started in 1991 with the Wise 1.0m and the Steward 2.3m tele-
scopes (Maoz et al., 1994). As a result, nine years later Kaspi
et al. (2000) provided mass estimates for the entire sample. The
new data, combined with previous results on Seyfert 1 galax-
ies, thus spanning a much wider range of intrinsic luminosity,
allowed to establish an average relation between the intrinsic lu-
minosity and the size, RBLR = aLγ, with γ ' 0.7, which allows a
secondary estimate of the black hole mass based on single-epoch
observations of luminosity and line width: MBH = bLγ∆2Hβ ,
where both the constant b and γ are determined statistically from
the available echo-mapping data. Recent studies show that γ is
in the range 0.5-0.7, depending on how luminosity is defined,
which lines are selected for the echo-mapping and the fitting pro-
cedure adopted (Kaspi et al., 2005; Bentz et al., 2006).
The extreme importance of secondary mass estimates relies
on the fact that on the basis of single epoch observations it is
possible to study the evolution in cosmic time of the mass distri-
bution of QSOs/AGNs, and to extend the studies of the relation
existing between QSOs and host bulges properties. However,
the above correlations with primary masses, based on echo-
mapping, were established for relatively close and faint AGNs
with z ≤ 0.4 and [λLλ(5100Å) <∼ 1046 erg s−1], thus it is
presently unknown whether they can be extrapolated to higher
luminosities and/or redshifts. For instance, the extrapolation of
the MBH − L relation (Kaspi et al., 2000), together with the
assumption that the known MBH − Mbulge − σ∗ relations holds
(Tremaine et al., 2002) , leads to predict the existence of galax-
ies with Mbulge ∼ 1013.1 − 1013.4M� and σ∗ exceeding 800 km
s−1. Such galaxies have never been observed, and their exis-
tence would put important constraints on galaxy formation mod-
els (Netzer, 2003). Therefore it is essential to extend the primary
mass measures to higher redshifts and luminosities. On the other
hand, for high QSO luminosities a large size of the broad line
region is expected. This would cause both a smoothing of the
line light-curve and larger time delays with respect to contin-
uum variations (Wilhite et al., 2005), thus the very detectability
and the amplitude of line variations are open questions.
A sample of objects with redshifts in the range 2 < z < 3.4
and apparent magnitude as faint as mV ∼18 is being monitored
by Kaspi et al. (2004) with the 9m Hobby-Eberly Telescope
(HET; Ramsey et al., 1998) and new results have been presented
recently (Kaspi et al., 2006). During their 6-year monitoring of 6
QSOs, significant continuum and emission-line variations were
detected in all targets and a preliminary black hole mass estima-
tion is given for one of them.
In the present paper we describe a new monitoring campaign
limited to objects with V < 15.7 and 1 < z < 4 which, thanks
to their apparent brightness, can be observed with the medium-
small 1.82 m Copernicus Telescope at Cima Ekar (Asiago, Italy),
through a service mode scheduling of a long term monitoring,
and allow to investigate whether: i) echo-mapping is feasible
for objects as bright as λLλ(5100Å) ∼ 1047 erg s−1 and ii) the
RBLR-luminosity correlation can be extrapolated to such bright-
ness. The paper is organized as follows. Section 2 describes the
sample, observations and the data reduction procedure. Section
3 describes the results for two quasars of the sample. Section 4
summarizes the results and discuss future prospects. In the fol-
lowing we derive λLλ(5100Å) from the flux in the Johnson V
band, extrapolating the flux density to the rest-frame λ = 5100Å
with a power law fν ∼ ν−0.5, and assuming a standard cosmology
Ho = 70 km s−1 Mpc−1, ΩM = 0.3, and ΩΛ = 0.7.
Table 1. The quasars monitored
Object z V log[λLλ(5100Å)]
[erg s−1]
APM 08279+5255 3.911 15.20 47.7
PG 1247+268 2.042 15.60 47.0
PG 1634+706 1.337 15.27 46.7
HS 2154+2228 1.290 15.30 46.7
4000 5000 6000 7000
Fig. 1. Average spectra of PG 1634+706 (upper panel) and PG
1247+268 (lower panel). Spectral ranges for continuum deter-
mination (short ticks) and line flux evaluation (long ticks), as
reported in Table 2, are shown.
2. Observations
2.1. Object selection
The sample has been extracted from the Veron-Cetty & Veron
(2003) (11th ed.) catalog with the condition δ > 0, V < 15.7
mag and z > 1 in order to select objects of bright enough intrinsic
luminosity to investigate the bright end extension of the RBLR vs.
λLλ(5100Å) relation (Kaspi et al., 2000, 2005). These conditions
identify 12 objects, only four of which were monitored, owing
to the limits on observing time. These four objects are listed in
Table 1.
Observations at intermediate redshift allow to sample the
variability of MgII λ2798, CIII] λ1909, CIV λ1559 lines, instead
of Hα, Hβ, Hγ observed in the low redshift study of Kaspi et al.
(2000). This allows to study either BLR at smaller distance from
the center, or regions of the same sizes of those producing the
Balmer lines, but using lines which respond to different part of
the continuum spectrum. The main emission lines falling within
the observed wavelength interval are indicated for two objects
in Table 2. Figure 1 shows the average spectra of PG 1634+706
and PG 1247+268.
With respect to other QSO monitoring programmes, ours
is the first which includes the MgII line in part of the ob-
served sources. This line is particularly significant, because i)
past Seyfert 1 monitoring campaigns conducted with IUE have
D. Trevese et al.: Line and continuum variability 3
shown its lag to be similar to that of Hβ (Clavel et al., 1991;
Reichert et al., 1994); ii) this line is most often used to derive es-
timates of black-hole mass from single-epoch spectra of high-z
QSOs (McLure & Jarvis, 2002), because its width is tightly cor-
related to that of Hβ. Deriving a lag for the MgII line would
therefore allow to estimate the black-hole mass most consis-
tently with respect to the results of Seyferts, which are mostly
based on the monitoring of Hβ.
APM 08279+5255 is one of the most luminous known QSOs
if its emission is considered isotropic. However it has been
shown to be lensed by a foreground galaxy (Irwin et al., 1998).
Three components, separated by a few tenths of an arcsec have
been detected in near-infrared images obtained with Keck tele-
scope and different models of the lensing field predict a few
days delay between photometric variations of the components
(Ledoux et al., 1998; Ibata et al., 1999; Egami et al., 2000):
a short time compared with the expected time scale of intrin-
sic variations. APM 08279+5255 is also a Broad Absorption
Line QSO (Irwin et al., 1998), which makes more difficult to
define regions free from either emission or absorption features
to measure continuum variations. The analysis of this object is
deferred to a forthcoming paper. HS 2154+2228 has been ob-
served so far only 5 times and the analysis requires further mon-
itoring. The other two objects PG 1247+268 and PG1634+706
are analyzed in this paper with the aim of : i) verifying the ad-
equacy of the observational data and reduction procedures, un-
der the assumption that variability amplitudes and characteristic
time scales can be extrapolated from the properties of fainter
objects; ii) possibly detecting line variations in objects as lumi-
nous as λLλ(5100Å) ∼ 1047 erg s−1, and compare their amplitude
with continuum variations. For both objects a star of comparable
magnitude, as close as possible to the QSO, has been selected
for the relative spectrophotometric calibration, described in the
next section.
2.2. Spectrophotometric Observations and Data Reduction
Observations were carried out at Asiago 1.82 m telescope
equipped with the Faint Object Spectrograph & Camera AFOSC
which is a focal reducer with reduction factor of 0.58, designed
to allow a quick switching between spectroscopic and imaging
modes. The scale at the focal plane is 21.7 ′′/mm. The detector
is a 1024x1024 thinned CCD array TEK1024 with 22x22 µm2
pixels corresponding to a scale of 0.473 arcsec pixel−1 and a
FOV of 8.14x8.14 arcmin2. We adopted a 8”.44-wide slit and
a grism with a dispersion of 4.99 Å pixel−1, providing a typ-
ical resolution of ∼ 15 Å in the spectral range 3500-8450 Å.
Spectrophotometric exposures are performed after orienting the
slit to include both the QSO and the reference star of comparable
magnitude, located at (12:50:11.5 +26:33:32) and (16:34:57.4
+70:32:49) (J2000) for PG 1247+268 and PG 1647+706 respec-
tively. The reference stars are included as internal calibrators for
QSO spectra, as described by Maoz et al. (1990) and Netzer et
al. (1990). The wide slit is necessary to avoid different fractional
losses of the QSO and star light due to possible non perfect slit
alignment and differential refraction, which could cause spuri-
ous variation of the flux ratios. Lamp flats and Hg-Cd ark spectra
were also taken for wavelength calibration. At each epoch, typi-
cal science observations consist of two consecutive exposures of
1800 s.
QSO and star spectra,Q(λ) and S (λ) were extracted with the
standard IRAF procedures. The QSO/star ratio as a function of
wavelength is computed for each exposure k = 1, 2
µ(k)(λ) = Q(k)(λ)/S (k)(λ). (1)
This quantity is independent of extinction changes during the
night. This allows a check of consistency between the two expo-
sures and the rejection of the data if inconsistencies occur (under
the assumption that QSO variations are negligible on 1 h time
scale). In fact, typical spectra of two consecutive exposures have
a ratio µ(1)(λ)/µ(2)(λ) of order unity, with deviations smaller than
0.02 when averaged over 500 Å, at least in the 4000Å − 7000Å
range. When discrepancies are larger than 0.04 both exposures
are rejected.
It is important to notice that, because of the large slit width,
there are small changes in the λ scale due to changes of the ob-
ject position within the slit (which are in general negligible in
the case of pairs of consecutive exposures). This suggested to
proceed as follows. As a first step we register the zero point of
the λ scale. However, since the changes of position within the slit
do not correspond exactly to rigid shifts of the λ scale, we select
portions of the spectra in wavelength intervals |∆λ| <∼ 1000Å
around each of the QSO emission lines considered (see Table 2),
and we determine the shifts for each portion.
Once the wavelength scales are registered, the QSO and star
spectra taken in the two exposures are co-added and the ratio
µi(λ) = (Q
i + Q
i )/(S
i + S
i ) (2)
is computed, at each epoch i. The reference star is flux calibrated
at a reference epoch. Since our aim is to compute relative flux
variations, we reduce all quasar spectra to this reference epoch,
by multiplying all the µi(λ) by the flux calibrated star spectrum
f S (λ): f Qi (λ) ≡ µi(λ) f
S (λ). We stress that we are not interested
in the absolute flux calibration, whose accuracy is of the order
of 20% and which is applied for the sole purpose reporting the
QSO spectra in physical units. The star spectrum f S (λ) adopted
is the same for all epochs, thus it does not affect relative flux
changes we want to measure.
The QSO spectrum f Q(λ) = c(λ) + l(λ) can be decom-
posed in a line l(λ) and continuum c(λ) spectra. Two values
of the continuum, cshort and clong at shorter and longer wave-
length with respect to the most prominent QSO emission lines,
are evaluated in regions as free as possible of other emission
features, defined by the wavelengths ranges (λshort,1,λshort,2) and
(λlong,1,λlong,2) respectively. Table 2 reports the observer-frame
wavelengths defining the continuum regions and the intervals for
line integration. These were chosen on the basis of the analysis
of Clavel et al. (1991); Reichert et al. (1994) of UV spectra of
low redshift AGNs observed with the International Ultraviolet
Explorer (IUE), with slight modifications to maximize the S/N
ratio (clong of MgII(λ2798 Å) falls outside the wavelength range
covered by the IUE spectrograph).
Line fluxes are computed as:
[F(Q)(λ) − cint(λ)]dλ, (3)
where cint(λ) is the linear interpolation through cshort and clong.
The extremes λ1 and λ2, also listed in Table 2, not necessarily
coincide with λshort ≡ (λshort,1 + λshort,2)/2 and λlong ≡ (λlong,1 +
λlong,2)/2, and are chosen to optimize the fl signal to noise ratio.
Direct B and R images in 15x15 arcmin fields centered on the
QSOs were taken to check possible variability of the reference
4 D. Trevese et al.: Line and continuum variability
Table 2. Wavelength intervals for lines and continua [Å]
Object z line λshort,1 λshort,2 λ1 λ2 λlong,1 λlong,2
PG1634+706 1.337 CIII]λ1909 4252 4298 4382 4590 4630 4674
PG1634+706 1.337 MgIIλ2798 6215 6285 6251 6777 6755 6815
PG1247+268 2.042 CIVλ1549 4380 4502 4598 4801 5202 5262
PG1247+268 2.042 CIII]λ1909 5535 5595 5770 5900 6025 6085
CIII]
2800 3200 3600 4000
t (JD-2450000)
2800 3200 3600 4000
t (JD-2450000)
Fig. 2. Light-curves of PG 1634+706 in the observer’s frame,
as relative flux variations ∆ f / f . Upper panels: emission lines,
middle panel: shorter wavelength continuum cshort, lower panel:
longer wavelength continuum clong, left: CIII]/λ1909 Å, right:
MgII λ2798 Å.
star. Comparison of the reference star with the brightest star in
the field showed a r.m.s. fractional flux variation of 0.02 in both
cases. Thus we can assume that the reference star is not variable,
at this level of accuracy.
Relative-flux calibration 1-σ errors, reported in figures 2 and
3, are estimated by adding in quadrature the r.m.s. uncertainty
on the flux of the reference star to the r.m.s. fractional uncertain-
ties of line or continuum fluxes, which, in turn, are computed by
comparing pairs of exposures which are added to form quasar
spectra of individual epochs.
3. Results and discussion
Figures 2 and 3 show the light-curves in the observer’s frame, for
emission lines and continua listed in Table 2, for PG1634+706
and PG1247+268 respectively. The light curves are expressed
in terms of the relative flux variations ∆ f / f , with respect to the
flux f at the reference epoch where the absolute calibration was
performed (see Section 2.2). The total time base is 3.3 years in
the observer’s frame, corresponding in the rest-frame to 1.4 yr
for the former QSOs and 1.1 yr for the latter. In the case of
PG1634+706 no significant variations are detected: neither in
CIII] line, which has the lowest S/N ratio due to its intrinsic
faintness, nor in the MgII line. The relevant continua are also
2800 3200 3600 4000
t (JD-2450000)
CIII]
2800 3200 3600 4000
t (JD-2450000)
Fig. 3. Light-curves of PG 1247+268 in the observer’s frame,
as relative flux variations ∆ f / f . Upper panels: emission lines,
middle panel: shorter wavelength continuum cshort, lower panel:
longer wavelength continuum clong, left: CIV/λ1549 Å, right:
CIII] λ1909 Å.
consistent with no variability at a level of ∆ f / f ∼ 0.02-0.03
r.m.s..
The result is different in the case of PG 1247+268. The
faintest line CIII] λ1909 Å shows a marginal evidence of vari-
ability and the stronger line CIV λ1549 Å seems to decrease
steadily during the observing period.
An almost steady decrease does not allow to derive possi-
ble line-continuum time delays from cross-correlation analysis.
However we can establish quantitatively the evidence of contin-
uum and, most important, of line variability. For this purpose we
define the unbinned discrete structure function
UDS F(τi j) =
|y(ti) − y(t j)|, (4)
where y(t) represents any of the line or continuum light-curves
considered, ti and t j are two observation epochs and τi j = ti − t j
is the time delay. The (binned) structure function can be defined
in bins of time delay, centered at τ:
S (τ) =
∑
UDS F(τi j)
 , (5)
where the sum is extended to all the M values of UDSF belong-
ing to a given bin of τ. We adopt the average of UDSF, instead
D. Trevese et al.: Line and continuum variability 5
0 200 400 600 800 1000
0 200 400 600 800 1000
Fig. 4. Structure functions for PG 1634+706 in the observer’s
frame. Upper panels: emission lines, middle panel: shorter wave-
length continuum cshort, lower panel: longer wavelength contin-
uum clong, left: CIII]/λ1909 Å, right: MgII λ2798 Å.
of UDSF2, since it is more stable, i.e. less sensitive to deviant
points; the factor
2 in equation (4) makes S (τ) equal to the
standard deviation in the case of a Gaussian distribution (see Di
Clemente et al., 1996). Structure functions for lines and continua
0 200 400 600 800 1000
0 200 400 600 800 1000
Fig. 5. Structure functions for PG 1247+268.Upper panels:
emission lines, middle panel: shorter wavelength continuum
cshort, lower panel: longer wavelength continuum clong, left:
CIV/λ1549 Å, right: CIII] λ1909 Å.
of the two quasars PG 1634+706 and PG 1247+268 are reported
in Figures 4 and 5 respectively, as a function of the time delay
in the observer’s frame. The error bars reported in the figures
represent simply the standard deviation of the UDSF in each bin
and not the uncertainty on their average value. The latter would
be 1/(M − 1) times the standard deviation, for M uncorrelated
UDSF values.
In the case of PG 1634+706 the structure function analy-
sis simply confirms what already appears from the light-curves,
namely no significant variations of lines and continua are de-
tected. The r.m.s. noise level is about 0.02-0.03 for continua and
0.05 for lines.
On the contrary, in the case of PG 1247+268, there is a clear
increase of variability for long time lags. To establish the signifi-
cance of this variability it is necessary to evaluate the probability
of the null hypothesis that a value of S (τ) is produced by pure
noise. For this purpose we generated mock noise light-curves ,
n(ti), i=1,Nepo, where Nepo is the number of observing epochs,
by extracting random numbers with a Gaussian distribution of
standard deviation σ, which represents pure noise, under the as-
sumption that the photometric noise is not correlated on the time
scale of our minimum sampling interval (namely about 10 days).
Then, the relevant structure function S mock(τ) has been computed
with the same binning adopted for real data. The simulation was
iterated for 107 times to generate, for each bin, the statistical dis-
tribution of the S mock(τ) values.
The value adopted for σ has been estimated, conservatively,
from the observed structure functions of figure 4 and 5 them-
selves, as the value of S (τ) in the first bin. This is, in fact, an
overestimate of noise, since it includes short-time-scale varia-
tions of the QSO, which however should be negligible on the
basis of previous studies of (fainter) QSOs (Giveon et al., 1999;
Kaspi et al., 2000). The almost constant value of the structure
functions in the case of PG 1634+706 confirms the adopted hy-
pothesis that noise is not correlated on long time scale.
For the CIII] line the r.m.s. noise (∼ 0.05) is larger than in
the case of the CIV line (∼ 0.02). This depends on the intrinsic
faintness of the CIII] line. According to the above simulations
the probability of the null hypothesis, for the second and fourth
bins, is P(> S ) ∼ 4 × 10−4. The local maximum of S (τ), for
CIII] at τ ∼400 days, is due to an “oscillation” of the light-curve
on a time scale comparable with the total time base of the ob-
servations. This indicates that a longer time base is needed to
properly average individual oscillations and to obtain a “station-
ary” structure function. Further sampling, trying to avoid yearly
periodicity due to periods of optimal observability, is needed to
obtain the resolution required for time delay measurements.
The structure function of CIV line shows a much more sig-
nificant variability: the probability of the null hypothesis in the
case of the last two bins, where S (τ) is 0.13 and 0.21 respec-
tively, is less than 10−6. The systematic decrease of CIV line
luminosity and the relevant continua might seem to rely on the
value of the light-curves at the first epoch. Thus we have recom-
puted the S (τ) after removing the first epoch. Indeed, in the case
of continua the systematic decrease of S (τ) is no longer evident.
On the contrary, for the line it remains practically unchanged
(except that the last bin is no longer present). Thus the fact that
the probability of the null hypothesis is smaller than 10−6 does
not rely on that particular point in the light-curve.
The main goal of new echo-mapping campaigns is to obtain
a direct measure of the central black hole mass and to establish
whether the scaling of QSO broad line region size and masses re-
main the same at higher quasar luminosity. The very detectabil-
ity of line variation in the most luminous QSOs was, until re-
6 D. Trevese et al.: Line and continuum variability
cently, an open question, since 3C273 with (λLλ(5100Å) <∼ 1046
erg s−1) was the brightest QSO with detected line variability
(Kaspi et al., 2005). An evidence of line variability in brighter
objects has been reported more recently by Kaspi et al. (2006).
The present detection of line variability in one of the two quasars
considered strengthens the notion that quasar emission lines re-
spond to continuum changes in high luminosity QSOs like in
lower luminosity ones and indicates that echo-mapping cam-
paigns are feasible for λLλ(5100Å) >∼ 1047 erg s−1.
Our results are consistent with those of Kaspi et al. (2006),
who have monitored sources of comparable luminosities and
redshifts. We have also detected variability, of comparable am-
plitude, in the continuum and CIII] and CIV emission lines in
one of our sources. Considering the shorter time baseline of
our observations, the fact that variability has not been found in
PG 1634+706 is consistent with the segments without variations
which can be seen in the Kaspi et al. (2006) light curves on this
time scale.
To estimate the further monitoring needed to measure line-
continuum delays, we should take into account that, on aver-
age, the amplitude of variability increases with the rest-frame
time lag. However this dependency is weak: log S (τrest) ∼
0.15 log τrest, according to previous statistical studies on large
quasar samples (Vanden Berk et al., 2004; De Vries et al., 2005).
In the case of our quasars PG 1634+706 and PG 1247+268, the
observer’s frame τobs = τrest(1+z) is dilated by a factor 2.3 and 3
respectively. The size of the broad line region, based on Balmer
line reverberation-mapping, is expected to scale with luminosity
as RBLR ∝ Lγ with γ=0.5-0.7, depending on how L is measured
and the RBLR vs. L fit is obtained (Kaspi et al., 2000, 2005; Bentz
et al., 2006; Kaspi et al., 2006). A Balmer line BLR of about 3
light years in the rest frame is expected for the relevant luminosi-
ties. On the other the hand CIV line corresponds to a BLR size
2 times smaller than Balmer lines (Peterson & Wandel, 2000).
Considering that in general a time baseline of about twice the
light-crossing time is required in order to obtain a reliable lag
from the light curves, we estimate that reverberation mapping
should be feasible after 5-6 more years of monitoring.
4. Summary
– We have initiated a campaign for the monitoring of 4 high
luminosity QSOs and we present the result for two of them
on the basis of 3.3 years of observation with the 1.82 m tele-
scope of the Asiago Observatory.
– We discuss the data reduction, the procedures adopted and
the level of accuracy attained in relative spectrophotometric
variability measurements.
– We perform a structure function analysis of the light-curve
and numerical simulations to establish the confidence level
of line variability detection.
– We detect line and continuum variability in one of the two
QSOs, PG 1247+268 of λLλ(5100Å) = 1047 erg s−1, with a
probability of less than 10−6 of the null hypothesis that the
observed structure function is produced by pure noise.
– This detection supports the notion that emission lines re-
spond to continuum variations as in substantially less lumi-
nous QSOs.
The results encourage the prosecution of the campaign which
should provide time delay, and black hole mass, estimates in 5-6
years.
Acknowledgements. We acknowledge support from the Asiago Observatory
team, in particular Hripsime Navasardian, for observations.
References
Clavel,J., Reichert, G. A., Alloin,D. et al., 1991, ApJ, 366, 64
Bentz, M. C., Peterson, B. M., Pogge, R. W., Vestergaard, M., & Onken, C. A.,
2006, ApJ, 644,133
De Vries, W.H., Becker, R. H., White, R.L., & Loomis, C., 2005, ApJ, 129, 615
di Clemente, A., Giallongo, E., Natali, G., Trevese, D., & Vagnetti, F., 1996,
ApJ, 463, 466
Egami,E., Matthews, K., Ressler, M., et al., 2000, ApJ, 535, 561
Ferrarese, L., & Merrit, D. 2000, ApJ, 539, L9
Giveon, Uriel, Maoz, Dan, Kaspi, Shai, Netzer, Hagai, Smith, Paul S., 1999,
MNRAS, 306, 637
Ibata, R. A., Lewis, M.J., Irwin, M. J., Lehar, J., & Totten, E. J., 1999, ApJ,
118,1922
Irwin, M. J., Ibata, R. A., Lewis, G. F., & Totten, E. J., 1998, ApJ, 505, 529
Kaspi, S., Smith, P.S., Meters, H., Maoz, D., Jannuzi, B.T. &Giveon, U., 2000,
ApJ, 533, 631
Kaspi, S., Netzer, H., Maoz, D., Shemmer, O., Brandt, W. N.,& Schneider, D. P.,
2004, in Coevolution of Black Holes and Galaxies, Carnegie Obs. Astroph.
Ser. Vol. 1, L. C. Ho ed., http://www.ociw.edu/ociw/symposia/series/
Kaspi, S., Maoz, D., Netzer, H., Peterson, B. M., Vestergaard, M., & Jannuzi, B.
T., 2005, ApJ, 629, 61
Kaspi, S., Bradt, W. N., Maoz, D., Netzer, H., Schneider, D. P., & Shemmer, O.,
2006, ApJ (in press), astro-ph/0612722
Kormendy, J., & Richstone, D. 1995, ARA&A,33, 581
Ledoux, C., Theodore, B., Petitjean, P., et al., 1998, A&A,339,L77
Maoz, D., Netzer, H., Leibowitz, E., et al. 1990, ApJ, 351, 75
Maoz, D., Smith, P. S., Jannuzi, B. T., Kaspi, S., & Netzer, H., 1994, ApJ, 421,
Marconi, A., Hunt, L. K., 2003, ApJ, 589L, 21
McLure, R. J., & Jarvis, M. J., 2002, MNRAS, 337, 109
Netzer, H., Maoz, D., Laor, A. et al., 1990, ApJ, 353, 108
Netzer, H., 2003, ApJ, 583, L5
Peterson, B. M., 1993, PASP, 105, 247
Peterson, B. M., & Wandel, A., 2000, ApJ, 540, L13
Silk, J. & Rees, M.J. 1998, A&A, 331, L1
Ramsey, L. W., et al. 1998, SPIE, 3352, 34
Reichert, G. A., Rodriguez-Pascual, P. M., Alloin, D. et al., 1994, ApJ, 425, 582
Tremaine, S. et al., 2002, ApJ, 574, 740
Trevese, D., & Vagnetti, F., 2002, ApJ, 564, 624
Vanden Berk, D, 2004, ApJ, 601, 692
Veron-Cetty M.P., Veron P., 2003, A&A, 412, 399
Vittorini, V., Shankar, F., & Cavaliere, A., 2005, MNRAS, 363, 1376
Wilhite, B. C.,Vanden Berk, D. E., Kron, R. G. et al., 2005, ApJ, 633, 638
http://www.ociw.edu/ociw/symposia/series/
http://arxiv.org/abs/astro-ph/0612722
	Introduction
	Observations
	Object selection
	Spectrophotometric Observations and Data Reduction
	Results and discussion
	Summary
