OPTICAL SPECTROPHOTOMETRY OF A COMPLETE SAMPLE OF 3CR LOBE-DOMINATED QUASARS
A 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; caars@angelo.edu, dhough@trinity.edu R. C. Vermeulen Netherlands Foundation for Research in Astronomy, Postbus 2, 7990 AA Dwingeloo, Netherlands; rvermeulen@astron.nl and A. C. S. Readhead California Institute of Technology, 105-24, Pasadena, CA 91125; acr@astro.caltech.edu 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. 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