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The Astronomical Journal, 130:23–46, 2005 July
# 2005. The American Astronomical Society. All rights reserved. Printed in U.S.A.
OPTICAL SPECTROPHOTOMETRY OF A COMPLETE SAMPLE OF 3CR
LOBE-DOMINATED QUASARS
C. E. Aars,1 D. H. Hough, L. H. Yu,2 J. P. Linick, and P. J. Beyer
Department of Physics and Astronomy, Trinity University, San Antonio, TX 78212; [email protected], [email protected]
R. C. Vermeulen
Netherlands Foundation for Research in Astronomy, Postbus 2, 7990 AA Dwingeloo, Netherlands; [email protected]
and
A. C. S. Readhead
California Institute of Technology, 105-24, Pasadena, CA 91125; [email protected]
Receivved 2004 December 16; accepted 2005 March 22
ABSTRACT
We present results from optical spectrophotometry of 19 of the 25 lobe-dominated quasars in the 3CR complete
sample. The optical spectra were obtained with the Hale 5 m telescope at Palomar Observatory, using the blue and red
CCDs of the Double Spectrograph, between 1984 and 1992. Additional data from the literature allow us to analyze
broad UV/optical emission lines in all 25 objects (100% completeness), with a total of 191 broad-line measurements
(far more than the 68 points in the preliminary results presented in the 2002 work of Hough et al.). We examine
correlations between the widths of the broad emission lines and several radio-based orientation indicators. These
include three measures involving the beamed radio emission on the parsec scale: (1) the prominence of the radio
nucleus R, defined as the rest-frame ratio of the nuclear to extended flux density at 5 GHz; (2) an alternate measure of
relative strength of the radio nucleus RV , defined as the rest-frame ratio of the 5 GHz radio core luminosity to the
optical V-band luminosity; and (3) a pseudoangle pseudo , derived from a rank ordering of sources based on both R and
projected linear size L. An additional orientation indicator based on kiloparsec-scale radio emission was also used: the
prominence of the straight, inner Very Large Array jets (Fj ), defined as the rest-frame ratio of the jet to jet-lobe flux
density at 5 GHz. We confirm earlier studies demonstrating a strong anticorrelation between R and the FWHM of the
Mg ii k2798 line; the FWHM also anticorrelates with RV and pseudo . However, the anticorrelation between R and the
FWHM of the C iii] k1909 line originally reported for only 12 objects in the 2002 work of Hough et al. is not seen in
the complete sample. To make separate investigations of lines that may originate in two physically distinct regions—
the intermediate-line region (ILR) and the very broad line region (VBLR)—we scaled the FWHM measurements to a
common mean and standard deviation for each region. The scaled FWHM of ILR lines (C iii] k1909, Mg ii k2798, and
H k4861) anticorrelates with R, RV , and pseudo , but the FWHM of the VBLR lines (Ly k1216, N v k1240, C iv
k1549, and He ii k1640) shows no evidence of any correlations with these three orientation indicators. These results
are consistent with models that divide the broad-line region surrounding the central supermassive black hole into an
inner, spherical, high-ionization VBLR and an outer, disklike, low-ionization ILR. These results are also consistent
with the unification of core- and lobe-dominated quasars. Simple models fit the log R versus Mg ii k2798 FWHM
relationship without unbeamed radio emission or turbulent broad-line cloud velocities; small unbeamed and
turbulent components offer slight improvements to the fits, but large contributions are ruled out and the best-fitting
range of orientation angles is only mildly restricted (10 –80 ). Some surprising results are anticorrelations
between the C iv k1549 and scaled VBLR line widths and log Fj . The explanation for these anticorrelations—and
the lack thereof for the main ILR lines—is not obvious but may be related to weaker beaming in the kiloparsec-scale
jets.
Key words: galaxies: jets — galaxies: nuclei — quasars: emission lines
Online material: machine-readable table
1. INTRODUCTION
observational classification of an object with a given black hole
mass, angular momentum, and accretion rate (e.g., Blandford
1987; Blandford & Gehrels 1999).
The morphology and kinematics of the radio jets have been
interpreted in terms of relativistic jet models (e.g., Blandford &
Ko¨nigl 1979; Scheuer & Readhead 1979; Bridle et al. 1994;
Vermeulen & Cohen 1994; Wardle & Aaron 1997). Most observational tests of models of parsec-scale jets have concentrated on
VLBI surveys of strong, compact radio sources—core-dominated
quasars and BL Lac objects—whose Doppler-boosted jets are
oriented at small angles to our line of sight (e.g., Pearson &
Readhead 1981, 1988; Taylor et al. 1994; Kellermann et al. 1998;
Orientation effects play a major role in the ‘‘unification’’ of
quasars and active galactic nuclei (AGNs; e.g., Scheuer &
Readhead 1979; Orr & Browne 1982; Barthel 1989; Padovani
& Urry 1992; Antonucci 1993; Ghisellini et al. 1993). The perspective from which we view the relativistic radio jets, accretion disk, and torus could be the main factor that determines the
1
Department of Physics, Angelo State University, San Angelo, TX 76909.
Department of Physics and Astronomy, Rice University, Houston, TX
77005.
2
23
24
AARS ET AL.
TABLE 1
Target List
Source
(1)
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
4C
3C
9 .............
14 ...........
47 ...........
68.1 ........
175 .........
181 .........
190 .........
191 .........
204 .........
205 .........
207 .........
208 .........
212 .........
215 .........
245 .........
249.1 ......
263 .........
268.4 ......
270.1 ......
275.1 ......
334 .........
336 .........
351 .........
16.49 ......
432 .........
z
(2)
mV
(mag)
(3)
MV
(mag)
(4)
2.018
1.469
0.425
1.238
0.768
1.387
1.195
1.956
1.112
1.534
0.684
1.11
1.048
0.411
1.029
0.311
0.646
1.396
1.519
0.557
0.555
0.927
0.371
1.296
1.805
18.21
20.
18.1
19.5
16.6
18.92
20.32
18.4
18.21
17.62
18.15
17.42
19.06
18.27
17.29
15.72
16.32
18.42
18.6
19.0
16.41
17.47
15.28
18.4
17.96
26.39
23.94
23.20
24.08
26.0
24.90
23.18
26.14
25.14
26.41
24.17
25.92
24.15
22.95
25.90
24.89
25.87
25.42
25.4
22.9
25.46
25.49
25.72
25.28
26.41
(J2000.0)
(5)
00
00
01
02
07
07
08
08
08
08
08
08
08
09
10
11
11
12
12
12
16
16
17
17
21
20
36
36
32
13
28
01
04
37
39
40
53
58
06
42
04
39
09
20
43
20
24
04
34
22
25.226
06.445
24.432
28.880
02.427
10.220
33.568
47.983
45.009
06.528
47.587
08.606
41.450
31.874
44.618
13.876
57.021
13.617
33.880
57.657
21.817
39.088
41.380
42.607
46.327
(J2000.0)
(6)
15 40 54.79
18 37 58.90
20 57 27.38
34 23 46.73
11 46 16.30
14 37 36.62
14 14 42.80
10 15 23.69
65 13 35.21
57 54 17.11
13 12 23.54
13 52 54.80
14 09 44.79
16 46 11.79
12 03 31.07
76 58 58.20
65 47 49.46
43 39 20.90
33 43 12.04
16 22 53.44
17 36 23.89
23 45 12.24
60 44 30.52
16 00 32.05
17 04 37.92
Notes.—Col. (1): Source name. Col. (2): Redshift of source. All z values are
from the compilations in HR89 or Paper I. Col. (3): Apparent V magnitude of
source. All mV values are from the compilation in HR89, except for 3C 68.1,
which is from Spinrad et al. (1985). Col. (4): Absolute V magnitude of source.
All MV values are K-corrected, as discussed in x 3. Col. (5): Right ascension of
source ( Paper I ). Col. (6): Declination of source ( Paper I ). Units of right ascension are hours, minutes, and seconds, and units of declination are degrees,
arcminutes, and arcseconds.
Zensus et al. 2002). In order to observe objects covering a wide
range in orientation, we have undertaken a VLBI study of the relatively weak nuclei in a complete sample of 25 lobe-dominated
quasars (LDQs) from the revised 3CR survey (Laing et al.1983);
see Table 1. An LDQ is defined to have a ratio of nuclear to extended flux density at an emitted frequency of 5 GHz, R, less than
1. These sources were selected on the basis of their steep spectrum, extended emission at low frequency (178 MHz), which
should minimize orientation bias due to the beaming of compact structures and thus permit statistical tests of orientationdependent radio jet properties. The sample for the VLBI study
was defined in Hough & Readhead (1989, hereafter HR89) and
has also been the subject of a deep Very Large Array (VLA)
imaging survey to probe the detailed kiloparsec-scale morphology in these objects (Bridle et al. 1994).
The principal results from the VLBI and VLA surveys
(based on the work of Hough et al. 2002, hereafter Paper I;
Bridle et al. 1994 and references therein; and some as yet unpublished data from D. H. Hough et al. 2005, in preparation) include
the following: (1) all 24 objects imaged by VLBI exhibit onesided parsec-scale jets (counting a couple of objects that show
only marginal, stubby extensions of low contours around the
core); (2) all 24 objects with high-quality VLA images display
one-sided kiloparsec-scale jets, with counterjet candidates in at
least seven sources, in stark contrast to the 15% jet detection
rate and total absence of counterjet candidates in Fanaroff &
Vol. 130
Riley (1974) type II (FR II) 3CR radio galaxies (Fernini et al.
1993, 1997); (3) all 23 objects with both VLBI and the VLA
images have the same ‘‘sidedness,’’ i.e., the small- and largescale jets lie on the same side of the VLBI core; (4) the mean
position angle of the VLBI jets lies within 6 of the position
angle of the innermost segment of the VLA jets; (5) estimates of
VLBI jet speeds for 14 sources are in the range of 0c to 7c
(H0 ¼ 70 km s1 Mpc1, q0 ¼ 0:5), with a distribution that is
most easily accommodated by a wide, but restricted, range of
orientations; (6) the prominence (measure of relative strength)
of the inner, straight VLA jet segments correlates well with the
prominence of the VLA radio nucleus (similar to log R but
normalized by the jet-side lobe only, rather than both lobes), but
with a slope less than unity; and (7) the prominence of terminal
‘‘hot spots’’ for the VLA jets anticorrelates with the amount of jet
bending. These results are generally consistent with relativistic
beaming models and unification scenarios for core-dominated
quasars, LDQs, and FR II radio galaxies, but with jet deceleration on the kiloparsec scale and with an additional phenomenon
( perhaps associated with jet bending) being primarily responsible for the appearance of counterjets.
Of course, the wide range of orientation in the 3CR LDQ
sample is valuable for testing any features of AGN models that
depend on viewing angle—at any wavelength—whether they
are due to relativistic beaming, aspect-dependent obscuration,
different projections of optically thick disks, or other physical
effects. Several observed properties of AGNs have been postulated to be useful orientation indicators. The radio parameter
R was introduced early on by Readhead et al. (1978) and Orr
& Browne (1982). The radio spectral index has been used in
several studies (e.g., Browne & Murphy 1987; Corbin 1991;
Ganguly et al. 2001). A measure that uses a combined rank
ordering of R and projected linear size of the kiloparsec-scale
radio source was introduced in Paper I. Wills & Brotherton
(1995) developed an improvement on R that they designate RV ,
which normalizes the radio core luminosity by the optical restframe V-band luminosity; Barthel et al. (2000) confirm the utility
of RV. Optical luminosity has been considered, since it could depend on viewing angle because of obscuration or optical beaming (Hutchings & Gower 1985; Browne & Wright 1985; Baker
1997). The slope of the optical continuum, thought to be related to
reddening due to obscuration, has also been suggested as a possible orientation indicator (Baker & Hunstead 1995; Baker 1997;
Richards et al. 2003). Wills et al. (1992) report that the percentage
polarization of the optical continuum may depend on orientation.
Several studies have employed the number distribution of broad
optical line widths (Brotherton 1996; Rudge & Raine 1998, 1999;
Grupe et al. 1999) to investigate the range of viewing angles. The
equivalent widths of narrow-line features have been postulated to
depend on orientation (e.g., Jackson & Browne 1991; Boroson &
Green 1992). X-ray properties, including both luminosity and
spectra, could be good indicators of viewing angle as well (e.g.,
Browne & Murphy 1987; Baker et al. 1995).
Of particular interest is the optical/UV broad-line region
(BLR), which is on a comparable physical scale to the VLBI
radio jets. The velocity field of the BLR has been the subject of
numerous studies, including those of Wills & Browne (1986)
and Wills & Wills (1986), who showed a strong inverse correlation between the relative strength of the radio nucleus (log R) and
the FWHM of the H Balmer line for a combined sample of
core- and lobe-dominated quasars (confirmed by, e.g., Brotherton
1996). In addition, Baker & Hunstead (1995) demonstrated that
the broad Balmer H and H lines and the Mg ii k2798 line
broaden systematically in composite Molonglo sample quasar
No. 1, 2005
OPTICAL SPECTROPHOTOMETRY OF 3CR QUASARS
spectra (Kapahi et al. 1998), increasing in width from coredominated (R > 1) to intermediate lobe-dominated (0:1 < R < 1)
to strongly lobe-dominated (R < 0:1) objects. These results are
intriguing, for as these authors point out, they are consistent with
broad-line emission from a flattened distribution, e.g., a disk. But
other studies have found a lack of correlations with other optical/
UV lines, or even positive correlations (e.g., for the C iv k1549
line and log R, in Wills et al. 1993). Building on a host of additional results, ‘‘stratified’’ models of the BLR have been constructed that include a spherical high-ionization, very broad line
region (VBLR) and a disklike low-ionization, intermediate broad
line region (e.g., Osterbrock 1993; Wills et al. 1993; Brotherton
et al. 1994; Corbin 1997; Puchnarewicz et al. 1997; Gaskell
2000). Controversy remains, however, over the structure of the
BLR, with contentions that there is no clear evidence yet as to
whether its geometry is spherical, disklike, cylindrical, bipolar, or
some combination of these (e.g., Corbett et al. 1998).
With the main goal of searching for correlations of broad
optical/UV emission-line widths with orientation indicators in
these objects, we performed optical spectrophotometry of 19
objects in the complete sample of 25 3CR LDQs. Together with
UV and additional optical data from the literature, we have sufficient data to include all 25 objects in the sample in various
correlation studies, attaining 100% completeness. In this paper,
we describe the observations, data reduction, and methods of
analysis (x 2), present notes on previous work/line measurements for selected objects in the sample (x 3), present detailed
results for all the individual objects, examine correlations of broad
optical/UV emission-line widths with several radio-based orientation indicators, and discuss the implications of the results for
models of the velocity field in the BLR of quasars and active
galaxies (x 4). These results are summarized in x 5.
2. OBSERVATIONS AND DATA REDUCTION
2.1. The Target Sample
The basic properties of the 25 3CR LDQs in our sample are
listed in Table 1. Of these 25 objects, a subsample of 13 were
observed between 1992 and 1993 (Table 2). Of the remaining
12 objects, six have spectra available in hard-copy form from
previous work (Paper I), although these spectra are not converted to rest frame or corrected for Galactic extinction. These
spectra are reproduced in this paper, and direct measurements of
line widths have been made for them. The line widths reported
for the other six objects use only data compiled from literature
searches.
The 13 objects with spectra in electronic form are 3C 9, 14,
47, 68.1, 181, 191, 204, 205, 268.4, 336, 351, 4C 16.49, and
3C 432 (Table 2). The six objects with spectra in hard-copy form
are 3C 175, 205, 208, 245, 263, and 334. The six objects whose
reported line widths are solely from the literature are 3C 190,
212, 215, 249.1, 270.1, and 275.1.
2.2. Observations
The quasars listed in Table 2 were observed at the 5 m Hale
telescope at Palomar Observatory, using the two CCD cameras
of the Double Spectrograph (Oke & Gunn 1982), between 1992
September and 1993 February (Table 2). The observations used
the D48 dichroic, and both CCDs are 800 ; 800 pixel arrays.
The blue channel used a diffraction grating with 300 lines mm1
and a tilt of 23 080 , giving an approximate wavelength coverage
from 3400 to 4700 8. In the spatial direction, only columns
150–340 were read out and the rows used a binning factor of
2, resulting in an approximate spectral resolution of 9 8 for a
25
TABLE 2
Observation Log
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
4C
3C
Source
(1)
V
(mag)
(2)
9 .............
14 ...........
47 ...........
68.1 ........
181 .........
191 .........
204 .........
207 .........
268.4 ......
336 .........
351 .........
16.49 ......
432 .........
18.21
20.0
18.1
19.5
18.92
18.4
18.21
18.15
18.42
17.47
15.28
18.4
17.96
Date
(3)
1992
1992
1992
1992
1993
1993
1993
1993
1993
1992
1992
1992
1992
Sep
Sep
Sep
Sep
Feb
Feb
Feb
Feb
Feb
Sep
Sep
Sep
Sep
21
20
20
20
12
12
12
12
10
21
20
21
20
Integration Time
(s)
(4)
S/N
( pixel1)
(5)
S/N
(6)
3600
3600
3600
2700
3000
3000
2700
3000
2700
2700
1800
3000
2700
57.0
7.8
51.7
14.1
11.3
16.1
5.1
17.4
18.1
25.2
133.1
9.5
38.3
25.2
6.9
22.3
11.3
6.0
7.8
2.8
8.0
12.1
13.7
54.5
5.4
16.6
Notes.—Col. (1): Source name. Col. (2): V magnitude from Table 1. Col. (3):
Date of observation. Col. (4): Integration time. Col. (5): Mean S/ N per pixel.
Col. (6): Standard deviation of S/ N per pixel across entire spectrum.
slit width of 200 . The red channel used a diffraction grating with
158 lines mm1 at a tilt angle of 20 370 . The resulting spectral
coverage ranged from 5000 to 9000 8. Only columns 310–566
were read out in the spatial direction, and there was no binning
in the dispersion direction, resulting in a spectral resolution of
approximately 18 8 with a slit width of 200 . Both channels used
a clear filter.
Dome flats were obtained during the night, along with arcs
for dispersion solutions. These arcs were taken immediately before or after an object’s spectrum was acquired to account for
flexion as the telescope moved between different pointings.
Standard stars (Oke & Gunn 1983) were followed in air mass to
provide absolute flux calibration.
2.3. Data Reduction
There were no bias frames taken during the run; the intent
was to use the dark strips to either side of the slit image to calculate a bias. However, light leakage into these areas of the CCD,
particularly when brighter objects were observed, became an unanticipated problem that made the dark areas of the CCD unsuitable for direct use as a template for the bias level. In order to
counter this problem, the following method was used.
The 2700 and 3600 s exposures near the zenith (air mass <
1:1) were used for a starting estimate of the bias level; the fields
around the quasars are relatively dark, and the low zenith angle
minimizes sky glow accumulating in the shadowed areas of the
chip. The first six columns in the dark strip to the left of the slit
image were averaged to form a one-dimensional estimate of the
bias as a function of wavelength, k. (We note that although bias
is clearly not a function of wavelength, contamination from sky
glow leaking into the dark strips does vary with k.) A linear fit to
the pixel values as a function of k and integration time provided
an estimate of the count rate due to dark current and sky leakage
for each row on the chip. This estimated sky and dark current
contamination was then subtracted from the template bias to leave
only pure bias signal. The bias template was then restored to two
dimensions and smoothed in the dispersion direction to eliminate
artificial striation.
The flats were summed in the spatial direction and a spline was
fitted to the resulting one-dimensional flat, which was then expanded back into two dimensions. The result was then divided
26
AARS ET AL.
into the flat to effectively remove the blackbody spectrum of the
dome flat lamp. The spectra were bias subtracted using the biases
described above and then flat fielded.
A dispersion solution was obtained for each object spectrum
individually, using HeHgAr arc lamps to provide reference
spectra (Hg for blue, Ar for red, and He for both channels). An
artificial sky was generated for and subtracted from each object
spectrum using software written by one of the authors (C. E. A.)
in IDL. The sky was sampled in two windows to either side of
the object spectrum. For each row of pixels, a quadratic and cubic
polynomial fit was made to the pixel values. The function that
produced the best fit was retained for that row and used to calculate pixel values for the corresponding row in the artificial sky.
The result was an accurate reproduction of the sky that could then
be subtracted from the original image.
The absolute flux calibration, correction for atmospheric extinction, and removal of sky absorption features was accomplished
using the method outlined in Lawrence et al. (1996). The spectra
were corrected for Galactic extinction and reddening with IDL
software using the Cardelli et al. (1989) reddening law and color
excesses from Schlegel et al. (1998). Finally, the spectra were
de-redshifted, and the fluxes scaled to the rest frame.
2.4. Observations and Data Reduction for the Six Spectra
in Hard-Copy Format
The quasars 3C 175, 205, 208, 245, 263, and 334 were observed using the instrumental setup discussed in x 2.2. The quasar
3C 334 was observed on 1984 September 18 (integration time was
1800 s). The remaining quasars were observed on 1990 January
24. Except for 3C 263 (1200 s), these objects were exposed for
3000 s. All six of these objects were reduced using the techniques
outlined in Lawrence et al. (1996).
2.5. Data Analysis
The spectra for our complete sample of 25 quasars come
in different forms, which require different analyses. The 13
spectra in electronic form (Table 2) were put through a standard
computer-based analysis to determine line equivalent widths
(EWs) and FWHMs. In addition, manual, or ‘‘direct,’’ estimates
of FWHM were made as a check on the accuracy of such measurements. For the six previously unpublished spectra in hard copy
form, only manual FWHM estimates could be made. Finally, for
the data from the literature—including spectra for the six sources
we did not observe—we used published EW and FWHM values
or, if no FWHM was quoted, we made manual estimates using
the published spectra.
The detailed procedure for the analysis of the 13 spectra in
electronic format was as follows. Major emission lines were
identified and EWs and FWHMs were measured for all identified lines. In the case of emission features with superposed absorption, the reported EW for the emission feature incorporates
the flux lost to absorption features. In all cases, the continuum
for the EW calculation was defined by the best of a linear or
power-law fit to a continuum on either side of the feature. The
FWHM were obtained in IRAF, using the ONEDSPEC package
and SPLOT. Single Gaussian and Lorentzian fits were used
except in cases in which (1) the feature was known or suspected
to be a blend of other features, or (2) the feature appeared to have
both a broad and narrow component. In those cases, SPLOT’s
line-deblending algorithm was used, incorporating a user-defined
number of Gaussian and Lorentzian components into the fit. For
single fits, the function (Gaussian or Lorentzian) that produced the
best fit to the line profile was retained. For deblends, an automated
routine was used that determined the best fit to the blend using
user-defined numbers and types of fitting functions. In practice,
for each blend the input parameters were varied to find the
combination of functions that best fit the blend. Gaussians were
most often the best fit for broad features in a blend, or the broad
component of a single line. Lorentzians tended to provide better
fits for narrow features or the narrow component of a single line.
When deblending multiple features, several variants (combinations of numbers and types of fitting functions) were tried, and
the combination that produced the best overall fit to the blended
feature and continuum was used. The type of function and the
number of components used to fit a particular line or blended
feature reported in Table 3 is available from the authors on
request.
The detailed procedure for the manual estimates of FWHM
was quite straightforward: fit a local continuum for the line and
then measure its FWHM directly on the hard copy. For our six
Hale 5 m spectra in hard-copy form, and for literature values
as needed, this ‘‘direct’’ FWHM was measured in the observed
frame and de-redshifted to the rest frame.
In general, FWHM values obtained by profile fitting and by
‘‘direct’’ measurement are in very good agreement. This is true
both for our own Hale 5 m spectra and for literature values. For
our 13 spectra in electronic format, there are discrepancies of a
factor of 2 or more in about 10% of the cases. These can be
attributed to various complications associated with Mg ii: absorption; a poorly defined continuum; a sharp spike in the line
profile (3C 9, 68.1, 268.4, and 432); noisy lines, sometimes on
the spectrum edge (C iii] in 3C 207 and 4C 16.49; C iv in
3C 204); or strong line asymmetry (C iii] in 3C 432). In all these
cases, the direct measurement is systematically less than the fitted
value, because the fitted profiles do not reproduce the line peaks
well. Thus, the direct measurement provides a single, well-defined
measure of FWHM that can be applied to all 25 sources in our
sample (see x 4).
Table 3 lists every detected line for all 25 objects in our sample.
For each line, a FWHM is measured, along with an EW when a
digitized spectrum was available. None of the FWHMs listed have
been deconvolved with the instrumental resolution. The reported
FWHMs are designated with the letters s, d, n, and w to clarify the
source of the line measurement. The s refers to a single-line fit
done using the IRAF software described previously. The d refers
to the direct measurements mentioned in the preceding paragraph. The n and the w are used for lines that appear to have both
a narrow and a broad component. The n marks the width of the
narrow component (determined using IRAF’s SPLOT routine).
The w marks the width of the broad (wide) component.
Table 3 includes line measurements both for objects the authors do not have spectra for (taken from literature searches)
and for lines beyond the wavelength range for our spectra (also
taken from literature searches). These lines are referenced in the
footnotes for Table 3 and are classified based on the measurement method detailed in the listed reference: (1) a single line
measurement is marked with a d and (2) a narrow/wide component deblended using a software algorithm has the narrow component designated with an n and the wide component designated
with a w.
3. NOTES ON INDIVIDUAL OBJECTS
For all the objects below, absolute magnitudes, MV, were
calculated assuming H0 ¼ 70 km s1 Mpc1 and a simple
K-correction given by
K(z) ¼ 2:5( 1) log (1 þ z);
ð1Þ
TABLE 3
Line Widths For UV and Visible Emission Features
Line Widths (8)
Source
(1)
Line
(2)
EW
(8)
(3)
3C 9 .................................
Ly k1216
Si/O iv k1400
C iv k1549
C iii] k1909
Mg ii k2798
Mg ii k2798
[O ii] k3727
Ly k1216
N v k1240b
O i k1305b
Si/O iv k1400b
C iv k1549
He ii k1640b
Al iii k1859b
Si iii] k1892b
C iii] k1909
Mg ii k2798
O iii k3133
O iii kk3429,3444
[O ii] k3727
[Ne iii] k3869
H k3970
H k4101
H k4340c
[O iii] k4363c
He ii k4686
H k4861
[O iii] k4959
[O iii] k5007
Fe iii k5720
H k6563
C iv k1549
C iii] k1909
Mg ii k2798
[O ii] k3727
[Ne iii] k3869
Ly k1026 + O vi k1035b
Ly k1216
N v k1240b
O i k1305b
Si/O iv k1400b
C iv k1549
He ii k1640b
Al iii k1859b
C iii] k1909
Mg ii k2798
[O ii] k3727
[Ne iii] k3869
H k4340
H k4861
[O iii] k4959
[O iii] k5007
C iv k1549
He ii k1640
C iii] k1909
Mg ii k2798
O iii k3429
[O ii] k3727
[Ne iii] k3869
76.1
7.2
21.1
4.3
34.1
37.8
15.4
...
...
...
...
...
...
...
...
...
109
1.8
11.7
4.9
18.2
5.4
7.0
7.4
11.5
5.8
49
35
93
3.0
461
63
5.9
33.5
8.2
3.2
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
138
2.5
19
40
3.9
11
5.3
3C 14 ...............................
3C 47 ...............................
3C 68.1 ............................
3C 175 .............................
3C 181 .............................
d
(4)
s
(5)
n
(6)
w
(7)
10
34
19
22
61
58
12
15a
27
11
13
22a
56
2
25
34a
91
19
27
16
27
16
53
...
...
60
43
26
22
37
146
6
7
76
9
9
12
26a
17
20
96
37a
85
43
20
91
16
27
96
242
30
15
20
5
14
83
8
14
9
41
39
22
24
128
61
11
...
...
...
...
...
...
...
...
...
82
22
25
12
25
12
54
31
30
94
30d
18
17
25
163
7
7
145
9
8
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
23
3
25
76
7
12
6
10
...
...
...
...
11
...
11b
...
...
...
22b
...
...
...
10b
18
...
...
...
...
...
...
...
...
...
24
...
...
...
64
...
...
...
...
...
...
17b
...
...
...
27b
...
...
...
...
...
...
...
...
...
...
20
...
4
10
...
...
...
66
...
...
...
...
101
...
34b
...
...
...
89b
...
...
...
44b
95
...
...
...
...
...
...
...
...
...
155
...
...
...
142
...
...
...
...
...
...
83b
...
...
...
97b
...
...
...
...
...
...
...
...
...
...
92
...
41
104
...
...
...
TABLE 3— Continued
Line Widths (8)
Source
(1)
Line
(2)
EW
(8)
(3)
3C 190e .......................
C iii] k1909
Mg ii k2798
O iii k3133
[ Ne v] k3426
[O ii] k3727
[ Ne iii] k3869
Ly k1216
C iv k1549
He ii k1640
C iii] k1909
Mg ii k2798
C iv k1549
C iii] k1909
Mg ii k2798
O iii k3133
C iv k1549
He ii k1640
C iii] k1909
Mg ii k2798
Ly k1026 + O vi k1035b
Ly k1216
N v k1240b
O i k1305b
Si /O k1400b
C iv k1549
He ii k1640b
Al iii k1859b
C iii] k1909
[ Ne iv] k2425
[O ii] k2470
Mg ii k2798
O iii k3133
[O ii] k3727
[ Ne iii] k3869
[ Ne iii] k3968 + H k3970
H k4101
H k4340c
[O iii] k4363c
H k4861
[O iii] k4959
[O iii] k5007
C iv k1549
He ii k1640
C iii] k1909
Mg ii k2798
[ Ne v] k3426
[O ii] k3727
[ Ne iii] k3869
H k3970
H k4101
H k4340
C iii] k1909
C ii k2326
Mg ii k2798
Ly k1216
N v k1240b
O i k1305b
Si /O iv k1400b
C iv k1549
...
...
...
...
...
...
168
14
4.2
4.0
43
42
8.8
108
4.0
...
...
...
...
...
...
...
...
...
...
...
...
19.7
1.3
1.9
27f
3.4
1.8
1.1
2.7
10.5
17.6
6.7
43
6.0
19.5
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
3C 191 ........................
3C 204 ........................
3C 205 ........................
3C 207 ........................
3C 208 ........................
3C 212g .......................
3C 215 ........................
28
d
(4)
s
(5)
n
(6)
w
(7)
13
59
16
5
7
7
6
15
7
13
62
16
6
49
17
67
9
56
84
19
14a
35
13
13
20a
73
22
12
3
1.5
44
12
11
13
14
31
...
...
45
14
16
25
40
27
73
18
13
21
20
53
53
47
85
80
15a
29
15
22
29a
...
...
...
...
...
...
12
17
9
10
79
34
4
53
22
...
...
...
...
...
...
...
...
...
...
...
...
27
4
...
33f
17
15
9
19
37
32
23
52
12
14
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
5
...
...
...
...
...
...
...
...
...
...
...
...
...
10b
...
...
...
19b
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
7b
...
...
...
23b
...
...
...
...
...
...
105
...
...
...
...
...
...
...
...
...
...
...
...
...
24b
...
...
...
50b
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
27b
...
...
...
116b
TABLE 3—Continued
Line Widths (8)
Source
(1)
3C 245 .............................
3C 249.1b .........................
3C 263 .............................
3C 268.4 ..........................
3C 270.1i ..........................
Line
(2)
EW
(8)
(3)
d
(4)
s
(5)
n
(6)
w
(7)
He ii k1640b
Al iii k1859b
Si iii] k1892b
C iii] k1909
Mg ii k2798
H k4861h
Ly k1216b
N v k1240b
C iv k1549
C iii] k1909
Mg ii k2798
[Ne v] k3426
He i k3587
[O ii] k3727
[Fe vii] k3760
[Ne iii] k3869
H k3970
H k4101
H k4340
Ly k1026 + O vi k1035
Ly k1216
N v k1240
O i k1305
Si/O iv k1400
C iv k1549
He ii k1640
Al iii k1859
Si iii] k1892
C iii] k1909
H k4861h
Ly k1026 + O vi k1035b
Ly k1216
N v k1240b
O i k1305b
Si/O iv k1400b
C iv k1549
He ii k1640b
Al iii k1859b
Si iii] k1892b
C iii] k1909
Mg ii k2798
[Ne v] k3426
[O ii] k3727
[Ne iii] k3869
H k3970
H k4101
H k4340
H k4861
[O iii] k4959
[O iii] k5007
Si/O iv k1400
C iv k1549
He ii k1640
C iii] k1909
Mg ii k2798
[O ii] k3727
[Ne iii] k3869
Si/O iv k1400
C iv k1549
Mg ii k2798
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
41
39
3.5
10
24
6.8
5.1
...
...
...
42
22
25
21a
60a
77
...
24
14
20
34
15
22
20
9
50
34
68
51
27
...
25
14
36
...
72
53
39
...
100
17
11a
33
17
22
17a
66
74
38
16
40
16
15
21
24
46
80
79
29
18
17
10
6
7
23
14
17
29
29
42
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
26
12
18
6
55
17
15
...
...
...
...
...
...
7b
39b
...
10
...
...
...
...
...
...
...
...
...
...
...
...
...
5
...
...
...
9
...
...
...
5
...
...
8b
...
...
...
13b
...
...
...
9b
43b
...
...
...
...
...
...
...
...
...
...
8
...
5
19
...
...
...
...
...
...
...
...
29b
134b
...
45
...
...
...
...
...
...
...
...
...
...
...
...
...
42
...
...
...
54
...
...
...
34
...
...
35b
...
...
...
52b
...
...
...
37b
159b
...
...
...
...
...
...
...
...
...
...
44
...
29
132
...
...
...
...
...
29
TABLE 3— Continued
Line Widths (8)
Source
(1)
Line
(2)
EW
(8)
(3)
3C 275.1j ...........................
Mg ii k2798
[ Ne v] k3426
[O ii] k3727
[ Ne iii] k3869
H k4340
H k4861
[O iii] k5007
Ly k1026 + O vi k1035b
Ly k1216
N v k1240b
O i k1305b
Si /O iv k1400b
C iv k1549
He ii k1640b
Al iii k1859b
Si iii] k1892b
C iii] k1909
Mg ii k2798
[ Ne v] k3346
[ Ne v] k3426
[O ii] k3727
[ Ne iii] k3869
H k3970
H k4101
H k4340
He ii k4686
H k4861
[O iii] k4959
[O iii] k5007
C iii] k1909
Mg ii k2798
O iii k3133
O iii k3429
[O ii] k3727
[ Ne iii] k3869
H k4340 + [O iii] k4363
H k4861
Ly k1026 + O vi k1035b
Ly k1216b
N v k1240b
O i k1305b
Si /O iv k1400b
C iv k1549b
He ii k1640b
Al iii k1859b
C iii] k1909b
Mg ii k2798
[O ii] k3727
[ Ne iii] k3869
[ Ne iii] k3968 + H k3970
H k4101
H k4340
[O iii] k4363
He ii k4686
H k4861
[O iii] k4959
[O iii] k5007
He i k5876
H k6563
C iv k1549
C iii] k1909
Mg ii k2798
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
70
61
1.8
12.4
3.1
2.6
12.9
11.5
...
...
...
...
...
...
...
...
...
6.2
1.2
1.7
1.2
...
11l
11l
0.43
3
7
20
0.9
271
214
71
73
3C 334 ..............................
3C 336 ..............................
3C 351 ..............................
4C 16.49 ...........................
d
(4)
s
(5)
n
(6)
w
(7)
41
15
18
13
49
45
13
34
19a
4
17
49
26a
55
37
24
48a
69
20
20
16
16
15
41
74
20
108
24
21
10
54
29
13
12
10
36k
44k
10
...
19
11
21
...
75
40
40
52
12
16
24
17
...
...
16
34
15
14
24
109
26
28
120
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
11
59
27
90
11
9
71
33
...
...
...
...
...
...
...
...
...
43
11
17
26
11
...
9c
14
51
14
14
25
108
32
91
154
...
...
...
...
...
...
...
...
10b
...
...
...
18b
...
...
...
10b
...
...
...
...
...
...
...
...
...
...
...
...
10
42
...
...
...
...
...
...
...
19
...
...
...
11
...
...
...
...
...
...
...
...
11c
...
...
19
...
...
...
33
...
...
...
...
...
...
...
...
...
...
...
54b
...
...
...
77b
...
...
...
42b
...
...
...
...
...
...
...
...
...
...
...
...
33
138
...
...
...
...
...
...
...
28
...
...
...
60
...
...
...
...
...
...
...
...
98c
...
...
192
...
...
...
272
...
...
...
OPTICAL SPECTROPHOTOMETRY OF 3CR QUASARS
31
TABLE 3—Continued
Line Widths (8)
Source
(1)
Line
(2)
EW
(8)
(3)
3C 432 ...................................
Ly k1216
N v k1240
Si /O iv k1400m
C iv k1549
He ii k1640
C iii] k1909
Mg ii k2798
101
1.1
4.2
63
2.1
63
27
d
(4)
s
(5)
n
(6)
w
(7)
15
5
...
41
6
44
68
25
5
13
49
9
80
106
...
...
...
...
...
...
...
...
...
...
...
...
...
...
Notes.—Col. (1): Source name. Col. (2): Line identity and wavelength. Col. (3): Equivalent width, in
angstroms. Col. (4): FWHM (8) for direct fit. Col. (5): FWHM (8) for single-component fit in IRAF.
Col. (6): FWHM (8) for narrow component in two-component fit to line in IRAF. Col. (7): FWHM (8) for
broad component in two-component fit to line in IRAF. For Cols. (4)–(7), an ellipsis indicates that a fit of
that type was not made to the line. Table 3 is also available in machine-readable form in the electronic edition
of the Astronomical Journal.
a
From Wills et al. (1995). The value is averaged with the value reported in Netzer et al. (1995) when
available.
b
From Kuraszkiewicz et al. (2002).
c
The FWHM values presented here are obtained from a successful deblend of the H and [O iii] lines
using the IRAF software package SPLOT.
d
The single-fit value for this line is probably a significant underestimate because of SPLOT’s inability to
properly locate the continuum when fitting blended lines consecutively (as opposed to simultaneously).
e
From Stockton & Ridgway (2001).
f
The Mg ii line in this spectrum is located at the blue-red interface of a spectrum taken on a nonphotometric night. Thus, the error associated with locating the continuum is higher than normal, which should
be considered when judging the accuracy of the EW and FWHM presented here.
g
From Smith & Spinrad (1980). Please note that in Paper I these line widths were incorrectly reported as
observed frame rather than rest frame.
h
An average of the values reported in Netzer et al. (1995) and Brotherton (1996) (when available). The
Mg ii k2798 and H features reported in Paper I were rough measurements taken from the discovery paper for
this object (Schmidt 1966) and are omitted /retracted in this paper.
i
Measured from spectra published in Aldcroft et al. (1994).
j
Measured from spectra published in Hintzen & Stocke (1986).
k
The spectrum is noisy at the location of the H + [O iii] blend and the H line, rendering an accurate
direct measurement of the FWHM difficult. The values quoted here should be assumed to have a high (100%)
error.
l
The EW reported for these two lines is an estimate for the blend and assumes that the H line possesses a
broad component.
m
The Si /O line is too noisy for an effective direct measurement of the FWHM to be made.
where is the spectral index and is taken to be 12 (Peterson
1997). The resulting values for K(z) are consistent with empirical results published by Cristiani & Vio (1990), Peterson
(1997), and Natali et al. (1998).
In x 3.1 we discuss the 13 quasars with spectra for which the
original data were reduced by the techniques described in x 2.
In x 3.2 we discuss the six quasars with spectra where the
original data were not available and measurements were made
using hard copies provided by two of the authors (R. C. V. and
A. C. S. R.). These spectra are presented in the observed frame
and have not been corrected for Galactic extinction. In x 3.3, we
comment on data taken from the literature for six additional
sources.
3.1. Quasars with Spectra in Electronic Format
The reduced spectra of the quasars discussed in this section
are presented in Figures 1–13. Emission-line measurements for
all the quasars in this sample, including those discussed in this
section, are listed in Table 3.
3C 9.—This quasar has already been established as an
absorption-line system, although not a broad absorption line
QSO, by previously published spectra (e.g., Anderson et al.
1987; Barthel et al. 1990, hereafter BTT90; Junkkarinen et al.
1991, hereafter JHB91). Major absorption is not apparent in the
Fig. 1.—Near-UV and optical rest-frame spectrum of quasar 3C 9 (z ¼ 2:018)
obtained with the Double Spectrograph on the 5 m Hale telescope at Palomar
Observatory. Broad and narrow emission lines are labeled. The break in the spectrum (right of C iv line for 3C 9) is caused by the gap in calibratable spectral coverage between the blue and red cameras.
Fig. 5.—Near-UV and optical spectrum for 3C 181 (z ¼ 1:382). This quasar
was imaged on an unphotometric night, limiting the accuracy of the absolute
flux calibration to 0.3 mag, resulting in the discontinuity between the blue and
red channels at around 1900 8. While this figure is corrected for the previously
stated redshift, the locations of the major lines suggest a slightly revised redshift
of z ¼ 1:393 0:003. We note that in this object’s spectrum we find no systematic offset between the redshifts of the broad and narrow lines.
Fig. 2.—Same as Fig. 1, but for 3C 14 (z ¼ 1:469). This object was too faint
for the blue camera to detect continuum emission; the flat region of the spectrum
blueward of the break represents the noise level for the spectrum, not the actual
continuum.
Fig. 3.—Same as Fig. 1, but for 3C 47 (z ¼ 0:425).
Fig. 6.—Same as Fig. 5, but for 3C 191 (z ¼ 1:956). The discontinuity is around
1600 8. The break around 1400 8 is the result of a masked cosmic-ray strike.
Fig. 4.—Same as Fig. 2 (region with flat continuum is at the noise level), but for
3C 68.1 (z ¼ 1:238). The break around 1800 8 is the result of a masked cosmic-ray
strike. While this figure is corrected for the previously stated redshift, the locations
of the major lines suggest a revised redshift of z ¼ 1:227 0:003. We note that in
this object’s spectrum we find no systematic offset between the redshifts of the
broad and narrow lines.
Fig. 7.—Same as Fig. 5, but for 3C 204 (z ¼ 1:112). The discontinuity is around
2500 8. The break around 2550 8 is the result of a masked cosmic-ray strike.
32
Fig. 11.—Same as Fig. 1, but for 3C 351 (z ¼ 0:371).
Fig. 8.—Same as Fig. 5, but for 3C 207 (z ¼ 0:684). The discontinuity is
around 2750 8. The break at 2400 8 is the result of a masked cosmic-ray strike.
Fig. 12.—Same as Fig. 1, but for 4C 16.49 (z ¼ 1:296).
Fig. 9.—Same as Fig. 5, but for 3C 268.4 (z ¼ 1:396). The discontinuity is
around 1900 8.
Fig. 13.—Same as Fig. 1, but for 3C 432 (z ¼ 1:805). While this figure is
corrected for the previously stated redshift, the locations of the major lines
suggest a revised redshift of z ¼ 1:826 0:002. No narrow lines are identified
in this spectrum, so we are unable to determine if a systematic offset exists
between the narrow and the broad lines.
Fig. 10.—Same as Fig. 1 (spectrum taken on a photometric night), but for 3C
336 (z ¼ 0:927).
33
34
AARS ET AL.
resonance lines (such as C iv k1549 or Mg ii k2798; Fig. 1),
indicating that the absorption is not in-system (a result supported
by the previous references). The object’s variability is rather
sedate; Netzer et al. (1996) show a median B ¼ 0:28 with an
intrinsic timescale, 0 ¼ 1:13 yr, in the observed frame. This
level of optical variability suggests that there should be a relatively small amount of variability in the broad emission lines and
implies that reasonable agreement should be expected between
EWs and FWHMs previously published and those published in
this paper. We note that the results reported in Table 3 are generally within 25% of those published in previous analyses of
this object’s spectrum (e.g., BTT90; Tytler & Fan 1992; Corbin
& Francis 1994), although there are individual disagreements
between certain FWHM measurements in the literature. We retract the reported line width for O iii k3133 in Paper I, as this
feature failed to meet our new, more rigorous line identification
criteria.
3C 14.—The low flux from this object results in a spectrum
with rather poor signal-to-noise ratio (S/N; Table 2). Because of
the low S/ N for this object, the published spectrum has been
smoothed by 4 pixels (in the observed frame, 8.8 8 in blue and
24 8 in red; both values are on par with the spectral resolutions
of 9 and 18 8, respectively). 3C 14’s spectrum (Fig. 2) seems to
show a flat continuum in the blue channel and a steeper continuum in the red. Because the transition from flat to steep is so
abrupt and is channel specific, it is likely that the ‘‘blue continuum’’ is below the sky noise, calling into significant doubt
the veracity of the measurements for both the C iv k1549 and
He ii k1640 lines. A NASA/IPAC Extragalactic Database (NED)
reference search on this object did not turn up any previously
published spectra.
3C 47.—This quasar showed significant variability during
the times when many published spectra were taken (including
those presented here; Netzer et al. 1996). Both the intrinsic variability and the difficulty of establishing the level for the blue channel continuum in this object make it likely that this object will
show significant line variability between spectra published on different dates. This line variability is, in fact, observed: the EWs
reported in this paper are markedly different from those reported in
Wilkes et al. (1999) but similar to EWs reported in Corbin &
Boroson (1996). Note that in the latter case, the spectra were taken
within 1 month of those presented here. Wills & Browne (1986)
report a FWHM for H that agrees with our measurement, but
Netzer et al. (1995) report a FWHM for this line that is significantly higher than ours (by a factor of 5).
3C 68.1.—The low-S/N spectrum for this object has been
smoothed by 4 pixels. As with 3C 14, it is unlikely that we are
picking up the blue continuum over the sky background at k <
1800 8, but the continuum for k > 1800 8 is steep (see Fig. 4).
Because the C iv k1549 and C iii] k1909 lines are in a region of
our spectrum where it is unclear whether the continuum is detected above the noise level, our estimates of the EWs and
FWHMs for these lines have large errors. We note in-system
absorption in the Mg ii k2798 line ( previously observed in a
spectrum published by Aldcroft et al. 1994). The quasar 3C
68.1 is a rather unusual object; the optical-UV continuum is
extremely red (F
/ 6 ; Boksenberg et al. 1976; Smith &
Spinrad 1980). It is one of the most lobe-dominated quasars
known (R 0:0007; Bridle et al. 1994; Paper I), and the unified
scheme would interpret it as the most highly inclined of the
3CR quasars (e.g., Orr & Browne 1982; Hoekstra et al. 1997;
Brotherton et al. 1998). The most recent in-depth spectroscopic
treatment of 3C 68.1 is Brotherton et al. (1998), who conclude
that the line of sight to the object is most likely skimming the
Vol. 130
edge of an obscuring torus. We note that a spectrum of 3C 68.1
taken by C. Lawrence and W. Xu on 1990 January 1 (1992,
private communication) was very similar to the one presented
in this paper.
3C 181.—This quasar is classed as an optically weak variable
(OWV) in Sirola et al. (1998), having a m < 0:15 mag yr1.
Previously published spectra for the object show very narrow
C iv k1549 absorption (Anderson et al. 1987; JHB91) that is also
present in our spectrum at full resolution (rather than as depicted
in Fig. 5). Aldcroft et al. (1994) also report Mg ii absorption (not
in-system) and EW measurements for C iii] k1909 and Mg ii k2798
that are in good agreement with those reported here. Furthermore, the EW measurement for C iv k1549 reported in Baldwin
(1977) is in good agreement with our value, although his reported FWHM matches only the narrow-component FWHM for
our spectra. The EW similarities are consistent with this object’s
status as an OWV.
3C 191.—The spectrum presented has relatively low S/N
(16.1) and is smoothed over 3 pixels (6.6 8 in the blue channel,
18 8 in the red channel). Immediately obvious in the spectrum
(Fig. 6) are the numerous and strong absorption features, which
are already quite well documented and studied (see, for example, Burbidge et al. 1966; Stockton & Lynds 1966; Bahcall
et al. 1967; Williams et al. 1975; Anderson et al. 1987; JHB91;
and Hamann et al. 2001). The presence of strong absorption in
the resonance lines implies a significant amount of in-system
absorption and may be responsible for the anomalously low
EW reported for the C iv k1549 feature. We do note that this
feature happened to fall on the interface between the red and
blue channels in our spectra and was imaged on a marginally
unphotometric night. Furthermore, the EWs reported in BTT90
and Williams et al. (1975) do not match those reported here,
although 3C 191 is an optically strong variable (OSV) according to Sirola et al. (1998) and so we do not expect the emissionline EWs taken from different spectra to match particularly
well. The FWHM measurements for the UV resonance lines
also show poor agreement between spectra of different epochs
(e.g., compare our results to those of Turnshek 1984). Interestingly, the variability study of Netzer et al. (1996) shows
relatively little variability in the (observed) B band compared to
the (observed) R band, which shows rapid and strong variation
(indicating that this quasar’s color fluctuates significantly with
time).
3C 204.—This object was exposed twice for 2700 s but was
only detected in one of the exposures, indicating that the conditions were clearly unphotometric. The exposure in which the
spectrum was detected produced a very low S/N (5.1), indicating
that intervening clouds were severely attenuating the light from
the object. Thus, the reported absolute flux scale for this spectrum is dubious, as is the relative shape of the spectrum. Because
of the low S/N, the spectrum was smoothed over 4 pixels. Furthermore, the continuum was not adequately detected in either
channel, as evidenced by its apparent flatness (Fig. 7), in spite
of the source having a relatively steep optical spectrum ( ¼
1:09; Yu 1987). This object shows a moderate variability amplitude, but the timescale for the variablility appears to be rather
long (Netzer et al. 1996). We were not able to find previously
published spectra for this quasar.
3C 207.—The spectrum taken is probably marginally photometric, and the S/ N is 17.4; the published spectrum (Fig. 8)
is smoothed over 3 pixels. This quasar is a strong variable
(Marziani et al. 1996) with a short timescale (Netzer et al.
1996). The EW of the H line reported in Marziani et al. (1996)
is similar to what we report (within 25%; see Table 3) but the
No. 1, 2005
OPTICAL SPECTROPHOTOMETRY OF 3CR QUASARS
FWHM of the [O iii] k5007 line differs by a factor of 4, a discrepancy that can be accounted for by the lower resolution of
the spectra presented in this paper.
3C 268.4.—Much of the previous interest in this quasar
has centered around the absorption features present in its spectrum (Anderson et al. 1987; JHB91; Aldcroft et al. 1993, 1994).
The spectrum we present here has been smoothed over 3 pixels
(approximately equal to the spectral resolution). We find that
this object is not part of any long-term variability monitoring program with published results, although NED does note that the
object is optically variable. Aldcroft et al. (1994) report an EW for
the Mg ii k2798 feature that is in very good agreement with the
value reported in this paper, although without an amplitude and
timescale for the variability we cannot comment on whether this
result is expected. Our spectrum (Fig. 9) possesses what appear to
be two absorption features in the 1900–2100 8 range that do not
appear to be observed in the aforementioned papers. The quasar
3C 268.4, like 3C 207, was not imaged on a photometric night, so
it is unclear how accurate the absolute flux scale is. Furthermore,
the absorption just redward of the C iii] k1909 feature makes it
difficult to properly match the blue channel’s continuum to that of
the red channel.
3C 336.—Netzer et al. (1996) suggest that this object is a
fairly strong variable (V > 0:5 over a 200 day interval), but
the long-term coverage of the object’s luminosity is rather poor
(e.g., we found no papers that adequately monitor 3C 336’s variability for >1 yr; the quasar is included in the Barbieri & Romano
[1981] study of quasar variability but was reported as too faint on
the blue photographic plates to analyze). In spite of the object’s
variability, however, our line measurements agree reasonably well
with similar measurements published by Brotherton et al. (1994)
and Steidel & Sargent (1991). Aldcroft et al. (1994), however,
report a markedly different value for the Mg ii k2798 EW than do
we or the previously mentioned papers (factor of 2 difference; and
see Fig. 10 and Table 3).
3C 351.—The data on this object’s variability in Netzer
et al. (1996) suggest small amplitude (B 0:1) brightness
changes over relatively short timescales (100 days) and somewhat larger brightness changes (B 0:3) over much longer
timescales (i.e., year-long scales). This would be sufficient
to catalog the object as an OSV using Sirola et al.’s (1998) criteria. There is a large amount of already published data on the
emission-line properties of this quasar, and a search through the
literature (e.g., Wills & Browne 1986; Boroson & Green 1992;
Marziani et al. 1996; Wilkes et al. 1999; Corbin 1997) indicates
that the object shows a fair amount of line variability (made
all the more certain by the high-S/N ratios that can be obtained
for spectra of this object with a minimum of effort), although
in many cases identical line strengths are reported because of
several papers in a series referencing the same spectrum (e.g.,
Boroson & Green 1992 and Corbin 1997). The line width for
He ii k3203 reported in Paper I is not included in Table 3, as this
feature failed to meet our new, more rigorous line identification
criteria.
4C 16.49.—This object is not as well studied as the other
quasars in our sample—the current literature focuses primarily
on its radio and bulk optical properties. Our literature search did
not turn up any previously published spectra. Because of the low
S/N, the spectrum published here (Fig. 12) has been smoothed
over 4 pixels.
3C 432.—Our reported EWs are similar to those published
by Osmer et al. (1994; converted to rest frame) and to those in
Corbin & Francis (1994), although the line widths reported in
the latter paper are not very good matches to ours (Table 3 and
35
Fig. 13). Our reported emission-line measurements are also somewhat consistent with those in Corbin (1991), Anderson et al.
(1987), and Foltz et al. (1986). Sirola et al. (1998) identify this
quasar as a weak variable.
3.2. Quasars with Spectra in Hard-copy Format
The reduced spectra of the quasars discussed in this section are
presented in Figures 14–19. These spectra have not been converted to the rest frame or corrected for Galactic extinction.
Emission-line measurements for these quasars are listed in Table 3.
Since only a limited analysis is possible for these quasars, we
present only brief comments on previous FWHM data from the
literature.
3C 175.—Our measurement of the H FWHM is in good
agreement with previous work (Wills & Browne 1986; Netzer
et al. 1995; Brotherton 1996). However, our measurement of
the C iii] FWHM is a factor of 3 smaller than that given by
Wills et al. (1995).
3C 205.—Our measurements of the Mg ii and C iii] FWHMs
are in good agreement with those reported by BTT90.
3C 208.—We did not find any previously published spectra
in the literature.
3C 245.—Our measurement of the Mg ii FWHM is in excellent agreement with Foley & Barthel (1990).
3C 263.—Our measurements of the H, Mg ii, and C iii]
FWHMs are in reasonably good agreement with previous work
(Wills & Browne 1986; Wills et al. 1995; Netzer et al. 1995;
Brotherton 1996).
3C 334.—Our measurements of the H and Mg ii FWHMs
are in reasonably good agreement with previous work (Corbin
1991; Wills et al. 1995; Netzer et al. 1995; Brotherton 1996).
3.3. Quasars with Spectra from the Literature
For six additional quasars, we rely solely on published data
as indicated below.
3C 190.—Stockton & Ridgway (2001).
3C 212.—Smith & Spinrad (1980).
3C 215.—Netzer et al. (1995), Wills et al. (1995), Brotherton
(1996), and Kuraszkiewicz et al. (2002).
3C 249.1.—Brotherton (1996) and Kuraszkiewicz et al.
(2002).
3C 270.1.—Aldcroft et al. (1994).
3C 275.1.—Hintzen & Stocke (1986).
For many of the 19 quasars observed with the Hale 5 m, we
augmented our line data with Hubble Space Telescope UV/optical
line measurements from Netzer et al. (1995), Wills et al. (1995),
and Kuraszkiewicz et al. (2002). For all references except
Kuraszkiewicz et al. (2002), spectra were published that allowed
us to confirm tabular FHWM values or to make our own direct
FWHM measurements. For Kuraszkiewicz et al. (2002), the
FWHM of their narrow and broad components nicely bracket
values from single-component fits or ‘‘direct’’ measurements by
other authors in every case but one, giving us confidence that we
can rely on their fitted FWHM values.
4. RESULTS
As detailed previously, the EW and FWHM measurements
obtained from the spectra of the target objects are presented
in Table 3. The reduced spectra, with major emission features
identified, are presented in Figures 1–19, although 3C 68.1 is
not used in any of the correlation tests discussed (see x 3). For
consistency the direct measurement of the line FWHM (d in
Table 3) is used in the correlation tests unless otherwise noted.
Fig. 14.—Same as Fig. 1, but for 3C 175 (z ¼ 0:768). This spectrum is
scanned from a hard-copy version provided by R. C. V. and A. C. S. R., with
labels for the spectral features added by C. E. A. and D. H. H. The ordinate is in
units of Jy. The spectrum is in the observed frame and has not been corrected
for Galactic extinction.
Fig. 16.—Same as Fig. 14, but for 3C 208 (z ¼ 1:11).
Fig. 15.—Same as Fig. 14, but for 3C 205 (z ¼ 1:534).
Fig. 17.—Same as Fig. 14, but for 3C 245 (z ¼ 1:029).
Fig. 18.—Same as Fig. 14, but for 3C 263 (z ¼ 0:646).
OPTICAL SPECTROPHOTOMETRY OF 3CR QUASARS
37
Fig. 19.—Same as Fig. 14, but for 3C 334 (z ¼ 0:555). The blue channel for this spectrum is shown in (a), and the red channel is shown in (b).
If a narrow/broad-line decomposition was the only available
measure of a particular feature’s FWHM, the narrow component was used.
4.1. The Baldwin Effect
(85% confidence for C iv and less than 80% confidence for
Mg ii). In addition, the spectra of 3C 191 and 4C 16.49 show
absorption-line systems that could be contaminating the reported
EW’s, in spite of efforts to identify these features and sum the line
flux estimated as lost to absorption back into the EW of the emission feature. Specifically, the Mg ii k2798 emission lines show
relatively strong absorption features with a slight blueshift relative
to line center that we take to be in-system.
Thus, while Figures 20 and 21 are suggestive of the Baldwin
effect, we do not have enough data to verify its presence at a
statistically significant level.
Baldwin (1977) demonstrated that the EW of the C iv k1549
emission line is inversely related to the core luminosity of the
source. Possible explanations for this effect include (1) decrease
in ionization parameter with luminosity, (2) decrease in covering factor with luminosity, and (3) an accretion disk that appears brighter at lower inclinations (Peterson 1997). The last
case assumes that the emission-line flux comes from an extended region outside the disk that radiates isotropically (and is
consistent with broad-line emission originating from a roughly
‘‘layered-shell’’ region outside the continuum source).
Figure 20 depicts log EW versus MV for quasars where the
equivalent width of the C iv k1549 line is available. Figure 21
depicts the same quantities for the Mg ii k2798 line. While a
visual inspection of both plots suggests that the Baldwin effect
is present (and stronger for C iv than for Mg ii), the low number
statistics definitively place the Pearson coefficients calculated
from the data at significantly less than the minimum required for
even marginal confidence in the significance of the correlations
As discussed in Paper I, Wills & Browne (1986) and Wills &
Wills (1986) showed an inverse correlation between the relative
strength of the radio nucleus (log R) and the FWHM of the H
Balmer line. Baker & Hunstead (1995) extended this anticorrelation to the H and Mg ii k2798 lines. The latter was confirmed in
Paper I, which reported r ¼ 0:55 at the 97.4% confidence level.
Paper I also reported a similar, but weaker, correlation for the
C iii] k1909 line (r ¼ 0:53 at the 92% confidence level). Both
results are interesting because they are consistent with broad-line
emission from a flattened distribution (such as a disk, as suggested
Fig. 20.—Luminosity of the C iv k1549 line (expressed as the log of the line’s
equivalent width) vs. absolute magnitude of the quasars in our sample (MV). The
Baldwin effect is apparent, although not at a statistically significant level (confidence 85%).
Fig. 21.—Luminosity of the Mg ii k2798 line (expressed as the log of the
line’s equivalent width) vs. absolute magnitude of the quasars in our sample
(MV). The Baldwin effect is (barely) apparent, although not at a statistically
significant level (confidence <80%).
4.2. FWHM Anticorrelations with R and RV
38
AARS ET AL.
Fig. 22.—Prominence of nucleus R (5 GHz emitted, log scale) vs. Mg ii k2798
emitted FWHM for 23 3CR LDQs. There is an anticorrelation between log R
and Mg ii FWHM.
Vol. 130
Fig. 24.—Same as Fig. 22, but for C iii] k1909 emitted FWHM for 21 3CR
LDQs. There is no correlation between log R and C iii] FWHM.
by Wills & Browne 1986). Furthermore, Wills & Brotherton
(1995) introduced the RV parameter, citing it as a better indicator
of orientation angle than log R. Correlations between various line
parameters and both RV and R are discussed in this section. We
expect the correlations between the line parameters and R to be
similar to those for the line parameters and RV because of the very
strong correlation between R and RV (r ¼ 0:74, >99.9% confidence). Figures 22–33 plot various FWHM measurements against
log R and log RV , with the correlation coefficient listed in the
upper right.
We first reevaluate the anticorrelations reported in Paper I,
taking into account the addition of nine objects not included in
the optical line-width study in Paper I, compounded with a more
rigorous, systematic technique for the measurement of line
FWHMs for the Mg ii k2798, C iii] k1909, and C iv k1549 lines.
A moderate anticorrelation is still detected for the Mg ii line.
The correlation coefficient is r ¼ 0:43 versus log R and r ¼
0:36 versus log RV (96.0% and 91.0% confidence, respectively; Figs. 22 and 23). These numbers are smaller than those
reported in Paper I. If the nine objects added to the data set in
this paper are removed, the correlation becomes r ¼ 0:35 (provided that Mg ii line widths for 3C 68.1 and 3C 249.1 are used, as
they were in Paper I), indicating that it is the reevaluated values
of the FWHM for the Mg ii line that are responsible for the change
in the correlation (rather than the inclusion of new objects in the
data set). It is clear, however, that the reported anticorrelation is
real, at least for the Mg ii feature.
In contrast to what was reported in Paper I, however, a reevaluation of the line FWHMs for the C iii] k1909 feature results in a reduction of any anticorrelation with log R to below
the marginal level. Whereas Paper I reports r ¼ 0:53, we find
r ¼ 0:15 (48% confidence; Fig. 24) on reevaluation of the data.
The anticorrelation between the width of C iii] and log RV is also
poor (r ¼ 0:19 with 59% confidence; Fig. 25). We must therefore conclude that these newer data do not show any correlation
between the width of the C iii] feature and log R or log RV .
Although the number statistics for the C iv k1549 line are a
bit lower than those for either the Mg ii or C iii] lines—we report
C iv FWHM measurements for 19 of the 25 quasars in the
combined optical sample derived from Paper I and this paper—
we evaluate the linear correlation coefficient for this line as
well. As with the C iii] line, we find no correlation between the
width of the C iv line and log R (r ¼ 0:22; Fig. 26) or log RV
(r ¼ 0:15; Fig. 27).
Paper I reported a significant anticorrelation between the
68 lines in the sample and log R (r ¼ 0:37 at the 99.2% confidence level). In this paper, we have expanded our database to
Fig. 23.—Same as Fig. 22, but with RV (ratio of core flux density at 5 GHz to
optical luminosity, log scale) rather than R. There is an anticorrelation between
log RV and Mg ii FWHM.
Fig. 25.—Same as Fig. 24, but with RV (ratio of core flux density at 5 GHz to
optical luminosity, log scale) rather than R. There is no correlation between
log RV and C iii] FWHM.
No. 1, 2005
OPTICAL SPECTROPHOTOMETRY OF 3CR QUASARS
Fig. 26.—Same as Fig. 22, but for C iv k1549 emitted FWHM for 19 3CR
LDQs. There is no correlation between log R and C iv FWHM.
191 lines and report that the addition of these lines erases the
correlation versus log R (r ¼ 0:018). The correlation between
log RV and the line widths is stronger but still below even a
marginal level of certainty (r ¼ 0:10 with 84% confidence).
As with Paper I, we also calculate an average velocity width
for each source and plot the results versus log R and log RV. In
Paper I, a strong anticorrelation was reported versus log R (r ¼
0:56). Once again, the additional data almost totally erases the
anticorrelation with log R (r ¼ 0:10), and no believable anticorrelation with log RV is seen (r ¼ 0:15). We conclude that
there is no overall correlation between all the lines and either
log R or log RV . This result is not unexpected, as our complete
data set now possesses a preponderance of high-ionization lines
(HILs), which we do not predict will show strong correlation to
orientation indicators because of the expectation that most HIL
emission originates in a non-disklike VBLR (Puchnarewicz et al.
1997; Wills et al. 1993).
Because different broad-line features have different intrinsic
widths, the comparisons presented above are a clear oversimplification. In Paper I, this oversimplification was dealt with by
‘‘scaling’’ the individual line species to a mean that matched the
overall mean of the sample. The result was an overall reduction in
the strength of the correlations but, presumably, a more accurate
representation of the true dependence of broad-line widths on the
Fig. 27.—Same as Fig. 26, but with RV (ratio of core flux density at 5 GHz to
optical luminosity, log scale) rather than R. There is no correlation between
log RV and C iv FWHM.
39
Fig. 28.—Prominence of nucleus R (5 GHz emitted, log scale) vs. weighted
mean of FWHM for broad lines in the spectra of 24 LDQs. No statistically
significant correlation is evident.
orientation indicator, R: the correlations dropped to r ¼ 0:29
(97.8% confidence) for the 68 lines in the sample and r ¼ 0:49
(94% confidence) when the scaled lines are averaged for each
source.
In this paper we have further refined our scaling technique so
that each individual line species is scaled to both the mean and
the standard deviation of the line widths as a whole. Specifically, if
v¯ is the mean line width and is the standard deviation of the
191 lines and v¯ line and line are the same quantities for a specific
species, then for each species a pair of scaling coefficients, a and b,
are calculated, where a ¼ v¯ /¯vline and b ¼ /(line a). Then a vector
of line widths for a given species, X, can be scaled to a new vector,
Y, with mean v¯ and standard deviation , while still preserving the
distribution of elements in X, by using the transformation
Y ¼ a½v¯ line (1 b) þ bX :
ð2Þ
This transformation causes all line species to have the same
mean and standard deviation across the 25 objects in our sample,
while maintaining the relative distribution of velocities for any
given species. Thus, the line species will all be on an ‘‘even
footing’’ prior to averaging. As with Paper I, the result of the
scaling was to actually reduce the strength of the anticorrelation,
Fig. 29.—Same as Fig. 28, but with RV (ratio of core flux density at 5 GHz to
optical luminosity, log scale) rather than R.
40
AARS ET AL.
Fig. 30.—Same as Fig. 28, but with only the VBLR ‘‘dominated’’ species,
Ly k1216, N v k1240, C iv k1549, and He ii k1640, contributing to the
weighted mean FWHM.
although the anticorrelations remained stronger for log RV than
for log R (and were not statistically significant in either case).
Specifically, for all 191 lines r ¼ 0:00 versus log R, and for log RV
remained at r ¼ 0:10. When averaged over sources, r ¼ 0:04
versus log R (Fig. 28) and r ¼ 0:15 versus log RV (Fig. 29).
None is a statistically significant correlation, although additional
line measurements might confirm a marginally weak correlation
for log RV versus the entire, unaveraged data set. The failure of
this scaling technique to highlight a correlation when all the broad
lines are included in the average is expected for the reasons discussed above (preponderance of HIL lines dominated by VBLR
emission).
Even a rudimentary numerical test will show that if correlated data are combined with uncorrelated data that cover the
same range, the value for r will decrease. The scaling algorithm
presented here does not distinguish between individual lines
that show correlations (such as Mg ii) and ones that do not (such
as C iv). Thus, even the scaling algorithm will group uncorrelated data with correlated data and therefore result in a lowered
value for r. From a physical standpoint, it is more informative to
group the various lines based on the environment they are believed to form in, predict the expected correlations, and then see
if the real data support the predictions. Reverberation mapping
Fig. 31.—Same as Fig. 30, but with RV (ratio of core flux density at 5 GHz to
optical luminosity, log scale) rather than R.
Vol. 130
Fig. 32.—Same as Fig. 28, but with only the ILR ‘‘dominated’’ species,
C iii] k1909, Mg ii k2798, and H k4861, contributing to the weighted mean
FWHM. There is an anticorrelation between log R and the weighted mean FWHM.
results have shown that the high-ionization lines form very
close to the central engine (e.g., Peterson 1997 and references
therein). While the exact geometry of the line-emitting clouds
close to the central engine is a matter of debate, there is a general consensus that the emitting regions are not purely disklike
(e.g., Gaskell 2000; Puchnarewicz et al. 1997; Marziani et al.
1996; Brotherton 1996; and Wills et al. 1993). This is in contrast
to low-ionization lines (such as H or Mg ii), which are presumed to come from clouds farther from the central engine,
arranged in an axisymmetric distribution (for example, Wu &
Han 2001; Rudge & Raine 1998, 2000; Gaskell 2000; Grupe
et al. 1999; Srianand 1998; Corbin 1997; Peterson 1997;
Puchnarewicz et al. 1997; Marziani et al. 1996; Wills et al.
1993; and Wills & Browne 1986). We therefore separate the
lines that occur most often in our sample into lines associated
with the so-called VBLR, which is close to the central engine,
and those associated with the intermediate-line region (ILR).
The VBLR is presumed to have a nondisk geometry, while the
ILR is presumed to be disk shaped. Given that R and RV are
indicators of orientation, we would expect to see an anticorrelation between R (or RV) and line FWHM for combined
data for lines dominated by ILR emission but not for lines
dominated by VBLR emission. Our VBLR lines are N v k1240,
Fig. 33.—Same as Fig. 32, but with RV (ratio of core flux density at 5 GHz to
optical luminosity, log scale) rather than R. There is an anticorrelation between
log RV and the weighted mean FWHM.
No. 1, 2005
OPTICAL SPECTROPHOTOMETRY OF 3CR QUASARS
Fig. 34.—Prominence of VLA-scale straight jet component (5 GHz emitted,
log scale) vs. averaged, weighted FWHM for VBLR lines for 18 3CR LDQs.
There is a marginal anticorrelation between log F j and VBLR FWHM.
He ii k1640, Ly k1216, and C iv k1549. To select these lines,
we focused on the C iv resonance line, which is present in nearly
all of our spectra but clearly shows no correlation to R or RV . We
then selected all lines from our sample that occur (1) in the
spectra of at least 10 quasars and (2) have a cross-correlation
time lag (derived from reverberation mapping) roughly equal to
or less than that of C iv (see Peterson 1997). Our ILR lines are
H k4861, C iii] k1909, and Mg ii k2798. Our selection criteria
for these lines were that they (1) be present in the spectra of at
least 10 quasars in our target sample and (2) have a crosscorrelation time lag of roughly twice that of C iv and Ly (the
species with the longest cross-correlation lags from among the
VBLR species or, in the case of Mg ii, a lower ionization energy
than C iii]). Thus, our ILR species are dominated by emission
originating at least twice as far from the central engine as our
VBLR species. If the VBLR emission is from clouds with a
nondisk geometry, then we would predict no correlation between their FWHM and the orientation indicators R and RV .
Furthermore, if the ILR emission is from clouds confined to a
disk, then we would expect to see an anticorrelation between
their FWHM and the orientation indicators. The VBLR lines
were scaled and averaged for each source, but no correlations
to R or RV are seen, as expected (r ¼ 0:10 versus log R, and
r ¼ 0:03 versus log RV ; see Figs. 30 and 31—note that we
excluded Ly as a test because of possible concerns about
absorption and radiative transfer effects, and still found no
correlations). When this procedure is repeated for the ILR
species, the expected anticorrelations are detected. The anticorrelation for log R is marginal, with r ¼ 0:34 at the 90%
confidence level (Fig. 32), but a bit stronger for log RV, where
r ¼ 0:37 at the 92.7% confidence level (Fig. 33). This anticorrelation is, of course, driven by Mg ii, but weak negative
trends in C iii] and H are consistent with an overall ILR anticorrelation. We interpret these results as supporting the hypothesis that the BLR emission in most AGNs comes from two
distinct ‘‘regions,’’ a nondisk VBLR close to the central engine,
with emission dominated by the high-ionization species, and a
disk-shaped ILR farther from the core, with emission dominated by the low-ionization species.
In order to verify that averaging and scaling would not create false correlations in our data set where none existed before,
we conducted numerical tests to confirm that (1) if a truly
random sample consisting of N1 data points is averaged into N2
41
Fig. 35.—Same as Fig. 34, but vs. the C iv k1549 emitted FWHM. There is an
anticorrelation between log F j and C iv FWHM.
data points by grouping the data into bins of size N1 /N2 (where
N2 < N1 ), the resulting data set shows a correlation where none
existed before, with a probability that exactly matches the
confidence levels associated with the Pearson coefficient for N2
data points, and (2) the scaling process used in this paper produces a correlation when none existed before <0.01% of the time.
These tests confirm that the scaling method discussed above does
not create correlations from uncorrelated data, and that the confidence levels quoted for the ‘‘averaged’’ lines are accurate.
4.3. Correlations with Other Orientation Indicators
Although Paper I only explored correlations with log R (to
which this paper has also added correlations with the related
RV parameter), we also test for correlations in the data between
the line widths and other common orientation indicators: the
straight jet prominence (Fj), the ‘‘RL’’ pseudoangle (pseudo),
and VLBI-based measurements of jet speed (app). The straight
jet prominence is defined as the ratio of the rest-frame 5 GHz
VLA straight jet flux to the lobe emission on the jet side (Bridle
et al. 1994). Surprisingly, this parameter appears to show at
least a marginal anticorrelation to the widths of the VBLR lines
(r ¼ 0:40 at 90% confidence; Fig. 34—although this correlation drops below marginal at r ¼ 0:31 if Ly is excluded)
and an anticorrelation with the widths of the C iv k1549 features
(r ¼ 0:55 at 98.3% confidence; Fig. 35). No other correlations are seen. The contrast between the anticorrelations with
log F j and with log R (and log RV) are surprising: the VBLR
lines, and specifically C iv, show anticorrelations with log F j
where none existed for log R, but no anticorrelations are seen
for the ILR lines and log F j , or Mg ii and log F j (unlike the case
with log R and log RV). Some of these differences are likely due
to the relatively weak correlations between the log F j orientation indicator and the log R and log RV indicators (r ¼ 0:545 at
99.1% confidence for log R, and r ¼ 0:387 at 92.6% confidence
for log RV), but it is unclear what these data indicate about the
distribution of material in the VBLR and ILR regions. One possibility, discussed in Bridle et al. (1994), is that Fj is simply
more vulnerable to ‘‘nonorientation’’ related effects (such as
deceleration in the straight part of the jet), making orientationrelated correlations less apparent. An inspection of the plot for
C iv (Fig. 35) shows the anticorrelation dominated by the data
points corresponding to 3C 432 (FWHM of 7940 km s1) and
3C 205 (FWHM of 12,980 km s1). If these two data points
42
AARS ET AL.
Fig. 36.—‘‘RL pseudoangle’’ measurement (pseudo) vs. averaged,
weighted FWHM for ILR lines for 24 3CR LDQs. There is a correlation between pseudo and ILR FWHM.
are removed, the anticorrelation becomes insignificant (r ¼
0:25). Since C iv is among the most prominent of the features
dominated by VBLR emission, it is not surprising that a moderately strong anticorrelation between C iv and log F j would
favor an anticorrelation between the weighted VBLR lines and
log F j (Fig. 34). The removal of the data points for these
quasars also decreases the strength of the anticorrelation between log Fj and the VBLR lines’ weighted FWHM (r ¼ 0:31
at 75% confidence). While we cannot completely rule out an anticorrelation between the dispersion velocity of the VBLR and the
radio prominence of the VLA-scale straight jet, the physical
reason for the anticorrelation (if any) remains unclear.
Both the pseudoangle and app are described in detail in
Paper I. There is not sufficient data to determine the degree to
which app and the various line widths discussed in this section
are correlated. On the other hand, because the pseudoangle is
very strongly anticorrelated to both log R and log RV (r ¼ 0:77
and 0.80, both at >99% confidence), we expect to see similar
relationships between pseudo and the various line FWHMs. As
expected, correlations involving pseudo do closely mirror those
involving log R and log RV . Correlations are detected to the
weighted ILR lines (r ¼ 0:37 at 92.2% confidence) and the
Mg ii k2798 feature (r ¼ 0:41 at 94.9% confidence). These
correlations are depicted in Figures 36 and 37.
The values for log R, log RV , log F j , pseudo , app , and averaged line FWHMs (both weighted and unweighted) are listed in
Table 4.
4.4. Correlations Using Broad-Line Components
As is evident from Table 3, many of the major resonance lines,
particularly Ly, C iv, C iii], and Mg ii, can be decomposed into a
broad and a narrow component. In the correlation tests discussed
thus far, either single-line fits or direct measurements of the line
FWHM were used, even when a broad/narrow decomposition
was available. Furthermore, if a single fit or direct measurement
was not available, the narrow component was selected for the
preceding correlation tests. Puchnarewicz et al. (1997) discuss in
detail the stratification hypothesis for the structure of the BLR,
dividing lines into low-ionization/low-energy species (e.g., H )
and high-ionization/high-energy species (e.g., C iv k1549). While
both types of lines have luminosity generated in both the VBLR
and ILR regions, the low-ionization lines (LILs) are dominated by
luminosity from the ILR, and vice versa for the high-ionization
lines. One consequence of this hypothesis is the prediction that
Vol. 130
Fig. 37.—Same as Fig. 36, but vs. the Mg ii k2798 emitted FWHM. There is a
correlation between pseudo and Mg ii FWHM.
the broad component of a line will be dominated by VBLR flux
and will therefore not be correlated with the various orientation
indicators.
This prediction is upheld in our quasar sample. In cases where
the Mg ii k2798 feature can be decomposed into a narrow and
wide component, we find no correlation between the wide components only and log R (r ¼ 0:04), log RV (r ¼ 0:27 with a
confidence of only 63%), and pseudo (r ¼ 0:10). Our conclusion
here is somewhat hampered by low number statistics (only 13 of
25 objects have a Mg ii line that can be divided into narrow and
wide components, and six of these divisions are estimates and not
SPLOT-based fits). Thus, we cannot make any definitive statement but these data do further support the hypothesis that the BLR
is stratified into two major components, including a non-disklike
VBLR.
4.5. Correlations with Other Quantities
In Paper I, no correlations were found between line width and
redshift, optical luminosity, or radio luminosity. These results
persist when the new quasar data were added to the sample. As
before, no correlation was detected between weighted line width
and redshift, z (r ¼ 0:07). Likewise, no significant correlation
was detected with the weighted line width and absolute magnitude, MV (r ¼ 0:28, <80% confidence). Finally, we find no
correlation between radio luminosity and weighted line width
(r ¼ 0:03).
There is a correlation between log RV and MV (r ¼ 0:51) but
this is expected, as RV is defined using MV. Likewise, we see an
anticorrelation between z and MV (r ¼ 0:49)—we expect such
a result from any sample that is not a complete luminositylimited sample over a specified volume. We do not find any correlation between z and log RV (r ¼ 0:06) and log R versus MV
(r ¼ 0:23, but confidence <80%) for our sample.
4.6. Model Fitting
Paper I presented the results of model fits to the anticorrelation between the ‘‘scaled’’ mean FWHMs and the orientation parameter, R, based on the assumptions that (1) the
BLR has a disklike geometry, (2) there is no unbeamed radio
flux, and (3) the measured line FWHMs are due to the orbital
motion of the disk around the central black hole, with no contribution from disk turbulence. We refined the models presented
in Paper I by attempting to account for both unbeamed flux and
disk turbulence and by fitting specifically to the data acquired
for the Mg ii k2798 feature.
No. 1, 2005
OPTICAL SPECTROPHOTOMETRY OF 3CR QUASARS
43
TABLE 4
Derived Quantities for QSO Sample
Source
(1)
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
3C
4C
3C
9 ...........................
14 .........................
47 .........................
68.1 ......................
175 .......................
181 .......................
190 .......................
191 .......................
204 .......................
205 .......................
207 .......................
208 .......................
212 .......................
215 .......................
245 .......................
249.1 ....................
263 .......................
268.4 ....................
270.1 ....................
275.1 ....................
334 .......................
336 .......................
351 .......................
16.49 ....................
432 .......................
log R
(2)
log RV
(3)
log F j
(4)
pseudo
(deg)
(5)
app(c)
(6)
Mean FWHM
(km s1)
(7)
Mean FWHMweighted
(km s1)
(8)
2.40
2.00
1.30
3.15
1.62
2.30
1.09
1.59
1.36
1.74
0.31
1.32
0.92
1.41
0.00
0.92
1.00
1.34
0.96
0.96
0.64
1.62
2.22
2.00
2.05
1.23
2.31
2.44
1.08
1.34
1.62
3.42
2.15
2.02
1.70
3.29
1.97
3.10
1.95
3.17
1.64
2.08
2.34
2.90
3.02
2.05
1.76
0.53
1.59
1.34
1.59
...
1.89
2.22
2.01
...
...
0.56
1.41
2.86
0.79
0.94
0.54
1.82
0.42
1.12
1.54
2.25
2.34
2.38
0.69
1.56
2.01
2.02
2.22
38
39
41
45
42
29
14
17
30
28
11
22
13
36
10
23
26
19
16
20
25
32
44
33
35
...
0
5
...
...
...
2
0
0
0
7
3
7
...
4
5
3
...
...
...
3
2
...
...
...
4686
6219
3969
3470
8262
3317
3300
2871
2731
8107
4212
4754
8976
4937
3408
5697
5340
2049
5445
3520
4697
2747
3814
7434
4420
4080
3855
4352
2832
7789
3779
3327
2064
2611
9192
3593
5198
8793
4552
4840
4736
5472
972
4031
2761
4963
4115
3523
8269
4644
Notes.— Col. (1): Source name. Col. (2): Log of ratio between radio core luminosity and extended radio luminosity at 5 GHz (R).
Col. (3): Log of ratio between radio core luminosity (at 5 GHz) and optical luminosity (RV). Col. (4): Log of straight VLA / MERLIN radio
jet prominence, from Hough (1994). Col. (5): Pseudoangle, assuming a range in inclination between 10 and 45 . Col. (6): Apparent velocity
of jet material (in terms of c), from Hough et al. (2002) and D. H. Hough et al. (2005, in preparation). Col. (7): Mean FWHM (in km s1)
derived from broad emission lines in quasar’s spectrum. Col. (8): Same as col. (7), but scaled using the procedure discussed in x 4.2.
In our models, we assume that R is related to the Lorentz parameter , the orientation angle , and the unbeamed radio flux
Ru , by the equation
R ¼ (R90 =2)
h
i
; (1 cos )(nþ ) þ (1 þ cos )(nþ ) þ Ru :
ð3Þ
In this equation, R90 is the value of R at ¼ 90 that is beamed
at shallower inclination angles, is the spectral index (assumed to be 0.2), and the index n is assumed to be 2 for a
continuous beamed jet. The expression is simply v /c, and
is 0.98 for ¼ 5, the value assumed for all models presented
here. Paper I assumes that the quasars in the sample have inclination angles 10 < < 45 , and this assumption is used
to estimate the inclination angles for the individual objects
(pseudo). Averaging pseudo for the five sample objects believed
to possess the highest inclination angles results in ¯ ¼ 42N2,
corresponding to a mean R, R¯ ¼ 0:01814. The equation above
¯ and an assumed value
¯ ,
can be solved for R90 in terms of R,
of Ru.
The flux for the Mg ii feature is assumed to come ( predominantly) from a disklike region, assumed to have a rotational
speed given by vr . Assuming a disk turbulence component of vt ,
the FWHM of the Mg ii feature will vary with by the relation
(Wills & Browne 1986)
1=2
:
FWHM observed ¼ v 2t þ v 2r sin2 ð4Þ
The ¯ value can be used, along with the mean FWHM of the
Mg ii feature for the five sample objects mentioned previously,
to arrive at an estimate of what FWHM observed will be, given a
turbulent velocity of vt and a viewing angle of ¼ 90 . We refer
to this parameter as FWHM observed ( ¼ 90 ). We use these parameters to construct models of log R versus FWHM where Ru ,
vt , R90 , and FWHM observed ( ¼ 90 ) are allowed to vary, and
select the models that produce the best fits to the data in log R,
FWHM, and both parameters simultaneously. For a given
¯ The R90
model, we define R¯ to be the model’s value of R at .
parameter is determined by allowing R¯ to vary between 0.0024
and 0.0902, while Ru is varied between 0% and 90% of the
value used for R¯ . Thus, R90 was always selected in such a way
as to ensure that R¯ would remain constant over the range used
for Ru . Likewise, FWHM observed ( ¼ 90 ) was allowed to vary
between 2350 and 23,500 km s1, with vt ranging between 0%
and 90% of the value of FWHM observed ( ¼ 90 ) and vr selected
such that the value of FWHM observed ( ¼ 90 ) would remain
constant over the range used for vt .
Models were evaluated that assumed that the sample had
an inclination angle range of 10 < < 45 , 10 < < 80 ,
20 < < 70 , and 1 < < 89 (i.e., totally random orientations), with ¯ and all other dependent parameters scaled to
the range in used. For each model, the goodness-of-fit estimate was computed, assuming (1) that the radio flux used to
calculate R was accurate to 10% and (2) that the FWHM can be
measured accurately to 15% (a value that assumes an error in
the continuum level equal to 10% of the height of the measured
44
AARS ET AL.
Fig. 38.—Model fits to anticorrelation between Mg ii k2798 emitted FWHM
and 5 GHz core prominence (log R). All three models use ¼ 5, n ¼ 2, and ¼
0:2. Model 1 is a fit only to log R and assumes vt ¼ 2800 km s1, vr ¼ 13;800 km
s1, R90 ¼ 0:0004, and no unbeamed flux. Model 2 (the two-parameter fit to the
data) assumes no unbeamed flux, R90 ¼ 0:0004, vr ¼ 14; 000 km s1, and no turbulent velocity. Model 3 is a fit only to FWHM and assumes vt ¼ 4700 km s1,
vr ¼ 10; 700 km s1, R90 ¼ 0:0006, and no unbeamed flux. The assumed range of
inclination angles is 10 –45 .
feature). Figures 38 and 39 display the results of fits in log R
(labeled model 1), in log R and FWHM simultaneously (labeled
model 2), and FWHM alone (labeled model 3) for 10 < < 45
(Fig. 38) and 10 < < 80 (Fig. 39). In all cases, the model 3
fits, on visual inspection, do not appear to follow the data well,
while the model 1 and 2 fits use very similar parameters.
When a two-parameter fit is used, the best model assumes
10 < < 80 , with Ru ¼ 0:0002, R90 ¼ 0:0018, vt ¼ 3400 km
s1, and FWHMobserved ( ¼ 90 ) ¼ 11;400 km s1 (implying
vr ¼ 10;900 km s1). The best two-parameter fit model for 10 <
< 45 (used in Paper I) has a lower value for R90 (0.00042),
no turbulent velocity or unbeamed flux component, and vr ¼
14;000 km s1. The fit, as measured by the reduced 2 (2N ), is
worse than that for the larger range by 81%. The other ranges tested also produce worse fits; 20 < < 70 has a 2N
that is 99% higher than that of the 10 < < 80 range. The totally random range (1 < < 89 ) produces a fit with 2N only
4% higher than for the 10 < < 80 range (and better than the
10 < < 45 fits). These results strongly suggest that our data
support a larger range in orientation angles among the LDQs
than many previous studies have indicated (e.g., Barthel 1989;
Wardle & Aaron 1997).
5. SUMMARY AND CONCLUSIONS
In this paper we present reduced spectra for the quasars 3C 9,
14, 47, 68.1, 181, 191, 204, 207, 268.4, 336, 351, 4C 16.49, and
3C 432. We also include previously reduced spectra for 3C 175,
205, 208, 245, 263, and 334. In these spectra we identify major
spectral features and measure their EWs (when digital data are
available) and FWHMs. We then examine correlations between
four orientation indicators (log R, log RV , log F j , and pseudo)
and line FWHM. We test for correlations using individual line
data for each quasar and both unscaled and scaled averages of
all identified broad lines in a particular quasar’s spectrum. The
unscaled average includes all lines in a particular quasar’s spectrum. The scaled averages include (1) all lines in a particular
quasar’s spectrum, (2) only well-represented lines ( present in at
least 10 quasar spectra) believed to have a significant lumi-
Vol. 130
Fig. 39.—Same as Fig. 38, but assuming a larger range in inclination angles
(10 –80 ). Model 1 is a fit only to log R and assumes vt ¼ 2900 km s1, vr ¼
9300 km s1, R90 ¼ 0:0019, and no unbeamed flux. Model 2 (the two-parameter
fit to the data) assumes Ru ¼ 0:0002, R90 ¼ 0:0018, vr ¼ 10;900 km s1, and a
turbulent velocity of vt ¼ 3400 km s1. Model 3 is a fit only to FWHM and
assumes vt ¼ 4100 km s1, vr ¼ 7100 km s1, R90 ¼ 0:0065, and no unbeamed
flux. Overall, the fits to the data are better than for the smaller inclination angle
range depicted in Fig. 38.
nosity component from the ILR, and (3) only well-represented
lines believed to be dominated by luminosity from the VBLR.
We find the following correlations at 90% or greater confidence
(where a marginal correlation indicates a confidence below
95%):
1. log R versus FWHM for the Mg ii k2798 line
(r ¼ 0:43).
2. log RV versus FWHM for the Mg ii k2798 line (marginal, r ¼ 0:36).
3. pseudo versus FWHM for the Mg ii k2798 line
(r ¼ 0:41).
4. log R versus scaled, mean FWHM for the ILR lines
(marginal, r ¼ 0:34).
5. log RV versus scaled, mean FWHM for the ILR lines
(marginal, r ¼ 0:37).
6. pseudo versus scaled, mean FWHM for the ILR lines
(marginal, r ¼ 0:37).
7. log F j versus FWHM for the C iv k1549 line
(r ¼ 0:55).
8. log F j versus scaled, mean FWHM for the VBLR lines
(marginal, r ¼ 0:40).
The correlations between the ILR lines and log R, log RV , and
pseudo , coupled with the lack of correlation between these orientation indicators and the VBLR lines, supports the hypothesis
that the BLR can be divided into two components, as per Wills
& Browne (1986), an inner region with a non-disklike geometry
(the VBLR region) and an outer region with a disklike geometry
(the ILR region). Furthermore, the Mg ii correlations to log R,
log RV, and pseudo vanish when only broad-line components are
used, a result predicted by a model in which low-ionization
lines are dominated by flux from the disklike ILR but also have
a broad-line contribution from the non-disklike VBLR. Thus,
we conclude that these correlations support the hypothesis that
the BLR clouds in AGNs are ‘‘stratified’’ into an inner VBLR
dominated by emission from HIL species and an outer ILR
dominated by emission from LIL species.
The lack of any correlation detected between the orientation indicators and all 191 line measurements or the mean line
No. 1, 2005
OPTICAL SPECTROPHOTOMETRY OF 3CR QUASARS
widths for each quasar (scaled and unscaled) is most likely due
to the preponderance of HIL species in our data set and is also
in reasonable agreement with the ‘‘stratification hypothesis’’ for
BLR cloud geometry in AGNs. We are unable to explain the
correlations detected between log F j and the C iv resonance line
or the VBLR lines as defined in our sample. We note that there
is only a weak correlation between log F j and the other three
orientation indicators used in this paper, and that this may explain the discrepancy, but we also acknowledge that this answer
is far from definitive.
We have fitted models to our data set for log R versus the
FWHM of the Mg ii k2798 feature. The model that best fits our
data assumes 10 < < 80 , with Ru ¼ 0:0002, R90 ¼ 0:0018,
vt ¼ 3400 km s1, and FWHMobserved ( ¼ 90 ) ¼ 11;400 km
s1 (implying that vr ¼ 10;900 km s1). The values for R90 , vt ,
and vr are consistent with values reported in Paper I and in Wills
& Browne (1986), although our data appear to imply that these
quasars have a larger range in viewing angles than was reported
earlier or is commonly reported in the literature.
We note that the average FWHM of H k4861 (82 8) and
Mg ii k2798 (63 8) for our LDQs are significantly wider (by
25 8) than those in the core-dominated quasars (57 8 for H
and 38 8 for Mg ii) studied by Baker & Hunstead (1995). These
results are consistent with unification of core- and lobe-dominated
quasars.
45
We thank the staff at Palomar Observatory for their strong
support during the observations. C. E. A., D. H. H., and J. P. L.
were supported by National Science Foundation (NSF ) grant
AST 00-98253 to Trinity University. D. H. H., L. H. Y., and
P. J. B. received funding from NSF grant AST 94-22075 to Trinity
University. D. H. H. was also funded by a Research Corporation
Cottrell College Science Award to Trinity University. The work of
D. H. H., R. C. V., and A. C. S. R. at the California Institute
of Technology was supported by NSF grants AST 82-10259,
88-14554, 91-17100, and 94-20018. We thank Charles Lawrence
and Wenge Xu for their efforts at observing and reducing optical
spectra that they provided to us for this paper. The late Larry
Gindler, Neal Pape, Glenn Kroeger, Tim Pearson, and Dave Meier
offered extremely generous and highly expert assistance in the
establishment and maintenance of the Astronomical Computing
Facility at Trinity University. Trinity student Lara Cross contributed to software installation, data reduction, and analysis in the
early portions of this work. Trinity students Eric Danielson and
Alex Webb contributed to recent analyses of Very Long Baseline
Array data that yielded some of the new parsec-scale jet speeds
reported in Table 4. We also thank Bev Wills for her generous assistance with early stages of the optical data reduction and valuable
advice regarding interpretation of the spectra. Finally, we thank the
referee for an extremely careful reading of the manuscript and
suggestions that have led to significant improvements in this paper.
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