PHENOMENOLOGY OF BROAD EMISSION LINES

Transcription

1 Annu. Rev. Astron. Astrophys : Copyright c 2000 by Annual Reviews. All rights reserved PHENOMENOLOGY OF BROAD EMISSION LINES IN ACTIVE GALACTIC NUCLEI J. W. Sulentic, 1 P. Marziani, 2 and D. Dultzin-Hacyan 3 1 Department of Physics & Astronomy, University of Alabama, Tuscaloosa, AL Osservatorio Astronomico di Padova, Vicolo dell Osservatorio 5, I-35122, Padova, Italy 3 Instituto de Astronomia, Universidad Nacional Autonoma de Mexico, Apdo Postal , Mexico D. F , Mexico Key Words Seyfert galaxies, quasars, accretion disks, spectroscopy, emission lines, line formation Abstract Broad emission lines hold fundamental clues about the kinematics and structure of the central regions in AGN. In this article we review the most robust line profile properties and correlations emerging from the best data available. We identify fundamental differences between the profiles of radio-quiet and radio-loud sources as well as differences between the high- and low-ionization lines, especially in the radioquiet majority of AGN. An Eigenvector 1 correlation space involving FWHM Hβ, W(FeII opt )/W(Hβ), and the soft X-ray spectral index provides optimal discrimination between all principal AGN types (from narrow-line Seyfert 1 to radio galaxies). Both optical and radio continuum luminosities appear to be uncorrelated with the E1 parameters. We identify two populations of radio-quiet AGN: Population A sources (with FWHM(Hβ) 4000 km s 1, generally strong FeII emission and a soft X-ray excess) show almost no parameter space overlap with radio-loud sources. Population B shows optical properties largely indistinguishable from radio-loud sources, including usually weak FeII emission, FWHM(Hβ) 4000 km s 1 and lack of a soft X-ray excess. There is growing evidence that a fundamental parameter underlying Eigenvector 1 may be the luminosity-to-mass ratio of the active nucleus (L/M), with source orientation playing a concomitant role. 1. INTRODUCTION Active Galactic Nuclei (AGN) are usually characterized by a strong nuclear emission-line spectrum at UV and optical wavelengths. AGN phenomenology encompasses sources that show: (1) no lines except when a highly variable continuum is in a low phase (Blazars), (2) only narrow lines (Seyfert 2, possibly, LINERS and narrow-line radio galaxies), and (3) both broad (FWHM ,000 km s 1 ) and narrow (FWHM km s 1 ) lines (Seyfert 1 galaxies, broad-line radio galaxies, (BLRGs), QSOs, and quasars). The narrow emission spectrum /00/ $

2 522 SULENTIC ET AL is dominated by forbidden lines, whereas the broad emission lines are produced mainly by permitted transitions and often shows superposed narrow components. The relative strengths of the broad and narrow components (hereafter BC and NC) formed the basis for the subclassification of Seyfert galaxies into intermediate types (Osterbrock & Koski 1976; Osterbrock 1977, 1981). See also Winkler (1992) for a revised version of the Osterbrock classification that was adopted for the AGN catalog of Véron-Cetty & Véron (1998). The AGN commonality in the context of this review involves sources that show broad emission lines most of the time. We do not consider sources with broad lines that are: (1) infrequently seen in BL Lacs (e.g. Corbett et al 1996), (2) only detected in polarized light (Seyfert 2s; Antonucci & Miller 1985, Kay & Moran 1998), and (3) only present at very low intensity levels in the nuclei of nearby galaxies (mini-seyferts; Peimbert & Torres-Peimbert 1981, Filippenko & Sargent 1989, Ho et al 1997). We present the AGN broad-line discussion in a slightly unconventional way by reviewing the subject in three somewhat orthogonal parts: (1) broad-line phenomenology, (2) physical interpretation of the phenomenology, and (3) confrontation with models for the structure and kinematics of the broad-line region (BLR). This division has certain advantages that will lead to a clearer picture of what we really know about the broad lines independently of the favored models. Underlying this approach is the assumption that statistical studies of lines (coupled with reverberation studies) offer the best hope for resolving BLR structure and kinematics. Since the last ARAA review (Osterbrock & Mathews 1986), large samples of high s/n optical spectra have been obtained using linear detectors. Much improved UV data have also come from the Hubble Space Telescope. AGN show an enormous range of redshift (0.006 z 6.0, with a peak in the redshift distribution near z 2.0; Hewitt & Burbidge 1993, Véron-Cetty & Véron 1998). The principal broad lines can be studied in the optical domain over the following source redshift ranges: HI Hα, ; Hβ, ; MgIIλ2800, ; CIII] λ1909, ; CIVλ1549, and HI Lyα A more extensive list of lines is given by Burbidge & Burbidge 1967, Wilkes 1999, and Kriss et al 1999 for EUV (see also the NIST physical reference database for line identification). Many of the largest (and highest s/n) studies have focused on two broad lines (CIVλ1549 and Hβ) that are representative of the high (HIL) and low (LIL) ionization lines (i.e. lines produced by ions with ionization potential 50 ev and 20 ev, respectively). Hβ and CIVλ1549 are also particularly useful because they are less contaminated by nearby lines and because they permit us to compare LIL and HIL properties in the same sources out to z 1.0. Improvements in near infrared detectors offer the promise of obtaining high s/n and resolution LIL spectra at higher redshift in the near future. As will become clear later, the HIL LIL distinction is quite important (Section 3.2). Table 1 summarizes the principal line profile data sources with resolution equal to or of at least 10 Å FWHM. The tabulation includes both original surveys (NEW) and statistical studies that used these data sources (LIT).

3 TABLE 1 Reference Data a NOBJ b RQ c RL d RES (Å) e Lines f Zrange g Gaskell 1982 NEW HIL/LIL Young et al NEW HIL Wilkes 1984 et al NEW 214 Few Most 5 15 HIL/LIL de Robertis 1985 NEW LIL Sargent et al NEW ? 2 HIL Sulentic 1989 LIT LIL Baldwin et al NEW HIL Sargent et al NEW HIL Espey et al NEW HIL/LIL Barthel et al NEW HIL Steidel & Sargent 1991 NEW HIL Corbin 1991 LIT HIL Stirpe 1991 NEW LIL Kinney et al IUE HIL Boroson & Green 1992 BQS LIL Tytler & Fan 1992 LIT HIL Steidel & Sargent 1992 NEW % 50% 4 6 HIL Eracleous & NEW LIL Halpern 1994 Wills et al LIT HIL Corbin 1993 LIT LIL Corbin & Francis 1994 LIT HIL Brotherton et al LIT HIL Wills et al FOS HIL Laor et al FOS HIL Hewett et al refs LBQS % 90% 6 10 HIL/LIL Marziani et al NEW HIL/LIL Brotherton 1996 NEW LIL Corbin & Boroson 1996 NEW HIL/LIL Corbin 1997 NEW LIL Grupe et al 1999 NEW LIL a Reference type, NEW a sample of more than 10 new spectra. LIT a paper that is primarily a statistical analysis of data from other references. b Total number of AGN. c Number of radio quiet sources. d Number of radio loud sources. e Approximate spectral resolution of the data. f Indicates a reference presenting primarily HIL, LIL or mixed sample of source. g The approximate redshift range of the spectra data presented or analyzed. A lower limit less than 0.1 is rounded to zero.

4 524 SULENTIC ET AL 2. LINE PROFILE PHENOMENOLOGY In our phenomenological discussion we identify the most significant similarities and differences in the broad lines within individual sources and between different AGN populations. This is attempted without the encumbrance of uncertain emission line physics or currently popular models, both of which are considered in later sections. The principal independent and robust parameters that exist for significant samples of AGN include: (1) line width; (2) line centroid velocity shift, if possible with respect to the AGN rest frame; and (3) line shape/asymmetry (defined in slightly different ways; de Robertis 1985; Stirpe 1991; BG92 1 ; Brotherton 1996; M96 2 ; Corbin 1997). Our focus on two lines LIL = Hβ and HIL = CIVλ1549 is sufficient to compare the principal broad-line properties and differences that depend on ionization level. Accurate measures depend on our ability to extract pure line profiles, and this requires several data-processing steps: (1) continuum estimation and subtraction, (2) modeling and subtraction of optical (FeII opt )oruv (FeII UV ) singly-ionized iron emission, (3) removal of superposed narrow line components, and (4) subtraction of overlapping narrow lines of other atomic species. A detailed discussion of these procedures is beyond the scope of this review. The references in Table 1 can be consulted for details. Subtraction of the blended UV and, especially, optical (broad-line) FeII emission has taken on a dual importance both to decontaminate the appropriate line and to parameterize the FeII emission strength that is emerging as an important AGN diagnostic (Section 3). It is important to point out that the FeII and other corrections require spectra with resolution of at least 10 Å (Table 1) and continuum (if we are ever really measuring it) s/n Line Profile Widths Full width at half maximum (FWHM) is one of the most robust line profile measures. Measures made at lower levels in the profile will be increasingly affected by uncertainties associated with the processing steps mentioned above. Figure 1 shows a histogram of the most accurate FWHM measures for Hβ BC and CIVλ1549 BC in low redshift radio-loud (RL) and radio-quiet (RQ) samples. The data for Hβ BC is a combination of measures from BG92, M96, Brotherton 1996, Corbin & Boroson 1996, and Corbin This combination provides measures for 180 sources with a strong RL overrepresentation (n = 96). Comparison of measures in common for individual sources indicates that, aside from a few spurious measurements, a 2σ uncertainty of 10% is a reasonable estimate of the error associated with FWHM Hβ BC measures. Figure 2 illustrates the range of Hβ BC width (and FeII strength) for two prototypical Seyfert (RQ) galaxies (indicated by arrows in Figure 1). The matching CIVλ1549 BC sample is limited to 76 objects of which n = 38 are RL (M96, Laor et al 1994, Sulentic et al 2000; a 2σ uncertainty of 10 15% or more is typical; see references for specific values). Other studies do not subtract 1 Abbr. for Boroson & Green Abbr. for Marziani et al 1996.

5 PHENOMENOLOGY OF AGN EMISSION LINES 525 Figure 1 The distributions of rest frame FWHM for upper: Hβ BC (z 0.5; n = 180), and lower: CIVλ1549 BC (z 0.8; n = 76). Radio-quiet (RQ) sources are hatched in each histogram. acivλ1549 NC component, which can yield quite different results (see Sulentic & Marziani 1999 for a recent discussion of this contentious issue). These data provide strong evidence that both HIL and LIL profiles are broader in RL sources. Sample mean and rms values are FWHM RL (Hβ BC ) 5900 ± 3700 km s 1 versus 3600 ± 1900 km s 1 for RQ. The corresponding CIVλ1549 values are FWHM RL (CIVλ1549 BC ) = 7440 ± 2000 km s 1 versus 6200 ± 2150 km s 1 for RQ. A Kolmogorov-Smirnov (K S) test indicates that the RQ RL difference is marginally significant (probability that the samples come from the same parent population). The comparison also indicates that CIVλ1549 BC is systematically broader than Hβ BC in both RL and RQ sources. The HIL-LIL differences show a larger amplitude, and at least for RQ, are statistically more significant than the RL-RQ

6 526 SULENTIC ET AL Figure 2 Spectra of typical AGN in the rest frame range λλ Å. Upper: I Zw 1, a FeII strong narrow line Seyfert 1. Lower: NGC5548, a FeII weak broad line Seyfert 1. Ordinate is flux in units of ergs s 1 cm 2 Å 1 [NGC 5548 was taken with a 2 arcsec slit so absolute flux scale is below the normalized flux scale of international monitoring campaign spectra (Korista et al 1995)]. ones (M96 found K-S probabilities of and for the HIL-LIL differences in RQ and RL samples respectively). These differences contradict results of studies that did not subtract CIVλ1549 NC where mean FWHM(CIVλ1549 BC ) FWHM (Hβ BC ) [e.g. FWHM(Hβ) = FWHM(CIVλ1549 BC ); Corbin 1991 with IUE data, Laor et al 1995, Corbin & Boroson 1996]. There is also evidence for a difference between RL core [FWHM(Hβ BC ) 4700 ± 400 km s 1 (error of the mean): n = 26] and lobe [FWHM(Hβ BC ) = 6000 ± 600 km s 1 n = 32] dominated sources (Brotherton 1996). This result was confirmed in an analogous comparison of FWHM(Hβ BC ) in steep (SS) and flat (FS) spectrum RL sources

7 PHENOMENOLOGY OF AGN EMISSION LINES 527 (9 sources in common) where FWHM FS (Hβ BC ) 4740 ± 2300 km s 1 (n = 32) and FWHM SS (Hβ BC ) 6130 ± 3060 km s 1 (n = 30; Corbin 1997). Sources that have been monitored for Hβ variability usually show an rms FWHM(Hβ BC ) that is within 10% of the mean FWHM(Hβ BC ) values. The 19 most reverberated sources (Wandel et al 1999) show FWHM mean /FWHM rms 0.95 ± 0.14; 74% of the sources show a larger value for FWHM rms (Hβ BC ), which reflects significant variation in the wings of the line. Some extreme exceptions usually involve the appearance of a new broad-line component. These include the transition from Seyfert 2 (or Liner) to Seyfert 1 (e.g. NGC 4151; Ulrich et al 1985) or from a single-peaked to a double-peaked BLR (e.g. Pictor A; Sulentic et al 1995, Halpern & Eracleous 1994). These results suggest that FWHM is a reasonable line diagnostic even in the presence of significant variations. All of these results should be viewed with caution because available data samples are often heterogeneous and subject to hidden selection effects (e.g. as we will see BG92 probably favors sources with narrow RQ profiles) Line Width Dependence on Redshift Albeit an early HIL-LIL comparison used IUE CIVλ1549 spectra (Corbin 1991), reliable direct HIL-LIL comparisons became possible only with the advent of the HST FOS database. High-redshift HIL-LIL FWHM comparisons are difficult for two reasons: (1) published infrared spectra for the Hβ region still suffer from low resolution and s/n (e.g. Nishihara et al 1997); and (2) although large samples of high s/n CIVλ1549 spectra exist for high-redshift sources (Table 1), they have not been corrected for the CIVλ1549 NC contribution. The largest available samples of RQ (BG92; M96) and RL (Brotherton 1996, Corbin 1997) data show no significant correlation between source luminosity and FWHM Hβ BC. We identified published spectra for a total of 53 sources (z 2.0) that showed an obvious inflection between the broad and narrow components of CIVλ1549 or where the line showed a rounded top consistent with the absence of any narrow component. Analog processing of this sample (obviously biased toward broader CIVλ1549 profiles) yielded FWHM estimates likely to be accurate within ±500 km s 1. We obtained a mean (RL+RQ) FWHM(CIVλ1549 BC ) 7300 ± 2100 km s 1, which is similar to the low-redshift values presented earlier. The available evidence suggests no significant difference in profile width between high- and low-redshift samples of AGN. CIVλ1549 also maintains the same profile width as W(CIVλ1549) systematically decreases with redshift/source luminosity (Section 3.1) FWZI Broad line profiles sometimes show inflections that complicate our interpretation of simple measures like FWHM. Inflections raise the possibility that broad lines are a composite of several kinematically and/or geometrically distinct emitting regions. The frequent mismatch in profile wings (Romano et al 1996) also point

8 528 SULENTIC ET AL in this direction. The best evidence for a second distinct component in broadline profiles involves a very broad component (VBC) that sometimes underlies Hβ BC and CIVλ1549 BC. It is unclear whether the VBC component sometimes seen in HeIIλ4686 (Ferland et al 1990, Marziani & Sulentic 1993) shows the same properties. We are aware of only two studies that have attempted to measure the underlying VBC (Corbin 1995: Hβ; Laor et al 1994: CIVλ1549). The former used the mostly RQ BG92 sample and found mean FWHM(Hβ VBC) 9500 ± 7500 km s 1 for 67 sources, almost 3 times the mean FWHM of the classical Hβ BC. The latter studied FOS spectra for five sources and obtained FWHM(CIVλ1549 VBC) in the range 12 24,000 km s 1 (no FeII UV was subtracted). An alternative method for studying the VBC involves available full width at zero intensity (FWZI) measures of Hα where FeII emission is not a serious problem (Osterbrock & Shuder 1982, Shuder 1984). We analog extracted FWZI values from other samples of high s/n Hα spectra (Stirpe 1989, Sulentic et al 1998), which yielded an additional sample of n = 47 mixed RL/RQ sources. A large complementary (n = 94) sample of FWZI (Hα) measures (Eracleous & Halpern 1994: EH94) includes all of the best known double-peaked and complex RL profiles. The mean values ( ±2000 km s 1 ) for the combined RL and RQ samples are FWZI RL (Hα) km s 1 and FWZI RQ (Hα) 17,950 km s 1. These are surprisingly similar, considering (1) a bias in the EH94 RL sample for unusually broad profiles, and (2) the significantly larger mean FWHM found for RL samples. It is unclear how we should interpret these results because some profiles with smooth, very broad wings show FWZI values just as large as obviously inflected profiles. 2.2 Line Profile Velocity Shifts Velocity displacement of broad lines is one of the most striking and potentially model-constraining profile diagnostics. Broad-line shifts are often measured relative to one another especially for the HIL, because strong UV forbidden lines are rare (Gaskell 1982, Wilkes 1984, Junkkarinen 1989). Because all broad lines show velocity shifts (MgIIλ2800 appears to be the most stable; Junkkarinen 1989), relative measures are dangerous and difficult to interpret. Ideally, one would like to know the mean line shift and dispersion for each line relative to the rest frame. Several comparisons exist between [OIII]λλ4959,5007 and radio HI and/or absorption line measures of the host galaxies (Boroson & Oke 1984, Vrtilek & Carleton 1985, Wilson & Heckman 1985, Appenzeller & Östreicher 1988). These studies suggest that the narrow lines show velocity displacements of ±100 km s 1 or less relative to the host rest frame. This supports the cautious use of the narrow line redshift as a measure of the local rest frame in an AGN (for exceptions see Erkens et al 1997 and Section 3.2.1). Figure 3 shows the distribution of line shifts for Hβ BC and CIVλ1549 BC in low redshift RQ and RL sources. The Hβ BC sample is a combination of data from Marziani et al 1996 (M96: n = 30 RL and n = 21 RQ), Eracleous & Halpern

9 PHENOMENOLOGY OF AGN EMISSION LINES 529 Figure 3 The distribution of rest frame peak line shifts for (upper)hβ BC (z 0.5) and (lower) CIVλ1549 BC. Radio-quiet (RQ) sources are hatched in each histogram (EH94: Hα measures for n = 75 RL sources), and Boroson & Green 1992 (BG92: n = 63 mostly RQ) samples. The CIVλ1549 BC data come from M96 and Sulentic et al 2000 (plus two AGN from Laor et al 1994), who measured CIVλ1549 BC shifts relative to [OIII]λλ4959,5007 and subtracted CIVλ1549 NC. These sources are supplemented by 21 lower-resolution IUE CIVλ1549 spectra (Espey et al 1989, Rodriguez-Pascual et al 1997a). The former measures the shifts relative to MgIIλ2800 and the latter measures those relative to the CIVλ1549 NC. CIVλ1549 BC and Hβ BC show line shifts in many sources, whereas Hβ NC and CIVλ1549 NC do not. The most impressive shift behavior involves the apparent systematic blueshift of CIVλ1549 BC in RQ sources. This trend was first revealed in the early 1980s (Gaskell 1982, Wilkes 1984). The mean CIVλ1549 BC shift

10 530 SULENTIC ET AL relative to [OIII]λλ4959,5007 (n = 21 RQ sources from M96) is v r 790 ± 890 km s 1. The distribution of shifts indicates that CIVλ1549 BC is not blueshifted by a constant value relative to rest frame, but instead shows a distribution of values from zero to at least 4000 km s 1. The CIVλ1549 BC blueshift in RQ sources is usually present at all levels in the profile from peak to 1/4 maximum. In other studies, the RQ blueshift is masked/muted by the variable contribution of CIVλ1549 NC and/or measures relative to a broad line (e.g. Junkkarinen 1989; Corbin 1990, 1991; Steidel & Sargent 1991; Corbin & Boroson 1996). Most of these samples still show evidence for a mean blueshift in the range 60 to 2000 km s 1. The mean Hβ BC shift values for RQ sources are 20 ± 349 and 48 ± 228 km s 1 (consistent with a mean of zero) for the M96 and BG92 samples, respectively. The sigma values given above are sample standard deviations that provide a measure of the range in amplitude of observed shifts. Typical (1σ ) uncertainties for shift measures are 50 and 100 km s 1 for Hβ BC and CIVλ1549 BC. These lines, respectively, show mean (and sample standard deviation) RL shifts +500 ± 1200 and 140 ± 620 km s 1 (n = 40 objects from M96 and Sulentic et al 2000), the latter mean consistent with zero. Both RL values are redshifted relative to the RQ values. It may be significant that CIVλ1549 BC consequently maintains a blueshift relative to Hβ BC in almost every measured source (see M96 for a few possible exceptions). It is premature to generalize, but present data suggest that relative to the source rest frame, CIVλ1549 BC changes from a net RQ blueshift to a mean RL shift consistent with zero. At the same time, Hβ BC changes from a mean RQ shift near zero to a mean RL redshift. Although generally less than ±600 km s 1, the RQ Hβ BC shifts are real, as profiles with a strong NC-BC inflection shown in Figure 4 make clear. Examples of a large RQ CIVλ1549 BC blueshift and RL Hβ BC redshift are shown in upper left panel Figure 5 and lower left panel Figure 6, respectively. Figure 4 Examples of small Hβ BC line shifts in relatively symmetric RQ sources that show a very clear inflection between the broad- and narrow-line components. From left to right, the peak line shifts are 565, +270 (marginal), and +420 km s 1, respectively. Vertical dotted line marks the NLR redshift assumed to be the local rest frame.

11 PHENOMENOLOGY OF AGN EMISSION LINES 531 As always happens in AGN phenomenology, there are exceptions to the rule. For example, RQ source H (Laor et al 1994) shows CIVλ1549 BC v r km s 1 relative to [OIII]λλ4959,5007 and the RL source 3C227 (Sulentic 1989; EH94) shows a v r (Hβ BC ) 1800 km s 1 (RL Hβ BC shows a known line shift range from 2000 to km/s). The more serious exception (H ) is not certain because strong FeII UV in this source was not subtracted. M96 concluded that the RQ RL mean shift differences were statistically significant. The M96 RL mean v r (Hβ BC ) +490 km s 1 can be compared with a larger RL sample (n = 75; EH94) where v r +186 km s 1, and a larger RQ sample (BG92; n = 63) where v r 48 km s 1. There are a small number of sources in common between M96 and these two samples. Other HIL are very difficult to disentangle from satellite lines (Lyα,CIII]λ1909, MgIIλ2800) and FeII UV emission, as a recent detailed study of IZw1 (Laor et al 1997a) makes clear. The general result is that all of the HIL are blueshifted relative to MgIIλ2800 in RQ samples. Most sources show a line shift within the envelope of the observed profile width (i.e. v r /FWHM 0.5; Sulentic 1989) Dependence on Redshift Currently available measures suggest that both line width and shift show the same phenomenologies over a wide range in redshift. Espey et al (1989) was the first to compare HIL and LIL data for the same high z sources, using IR measures of Hα. Their results are fully consistent (despite a variable CIVλ1549 NC contribution) with the low-redshift results of M96 (see Figure 3). Examples of unambiguous large blueshifts (e.g. Q ; Sargent et al 1988) and redshifts (e.g. Q ; Barthel et al 1990) can be found in the high-redshift sources listed in Table 1. The role of MgIIλ2800 may be pivotal in some sources if the results for Q are typical (Carswell et al 1991). In that (z 2) source, the HIL show blueshifts ( v r 1800 km s 1 for CIVλ1549 BC ) and the LIL redshifts (+700 for Hα) all relative to [OIII]λλ4959,5007. Evidence has been advanced (from studies of high-redshift samples) that line shift correlates directly with line ionization level (Tytler & Fan 1992). This requires further study because the higher ionization line shifts were measured relative to Hβ BC, which is a kinematically volatile line Shifts at Zero Intensity Limited studies of the underlying VBC discussed in the previous section suggest very uncertain mean shifts of v r +800 km s 1 for Hβ (not significant with an rms scatter of 3100 km s 1 ; Corbin 1995) and ± 3100 km s 1 for CIVλ1549 (Laor et al 1994). Apparently contradictory results come from a large sample of RL sources (EH94) where the mean shift at ZI for Hβ is found to be only v r +270 ± 1070 km s 1. It is interesting that the EH94 profiles, an extreme (doublepeaked or boxy ) sample, show little evidence for a redshifted VBC often seen in other samples of RL sources. M96 found Hβ BC v r (ZI ) +80 ± 2000 km s 1 for RQ and v r (ZI ) ± 1900 km s 1 for RL, which is more consistent

12 532 SULENTIC ET AL with the frequent observation of a red shelf underlying [OIII]λλ4959,5007. The different results may indicate that a larger Balmer decrement exists in the RL VBC or that a VBC component is not present in all sources. 2.3 Line Asymmetry and Shape Higher-order moments of the broad lines suffer from larger uncertainties while offering a better possibility to constrain models. Some of the first Hβ BC measures (de Robertis 1985, Sulentic 1989) showed that large numbers of both red and blue asymmetric profiles exist. Examination of the 15 RL sources in the latter survey showed a preference for red asymmetries (the asymmetry index A.I. = [C( 3 4 ) C( 1 )]/FWHM, where C(i) are profile centroid measures at different levels). We made a comparison of the measured asymmetry and shape parame- 4 ters for sources in common in the samples previously discussed. This comparison suggests that measures of asymmetry are good to only one significant figure (0.1). The preference for red asymmetries in the RL AGN is the most robust result among the higher order moments. BG92 and M96 found A.I ± 0.10 and 0.06 ± 0.10, respectively, for RL sources. Two large studies of RL sources confirm the RL red asymmetry excess with A.I (Brotherton 1996) and (Corbin 1997). The lobe-dominated (steep-spectrum) RL population shows evidence for slightly more extreme red asymmetries. Both BG92 and M96 find A.I ± 0.10 for the RQ samples,with no preference for red or blue asymmetries. RQ sources show a scatter of CIVλ1549 BC asymmetries between ±0.2 (one source with +0.32) with A.I (M96). This lack of a strong preference for red or blue asymmetries is particularly remarkable and much in need of confirmation, considering the systematic blueshift that CIVλ1549 BC shows in RQ sources. Numerous claims for systematic CIVλ1549 BC asymmetries, especially in RQ samples, exist in the literature (e.g. Corbin 1992; Wills et al 1993, 1995; Corbin & Boroson 1996). Measurements of the composite CIVλ1549 line with systematically blueshifted broad component and unshifted NC component would be expected to produce strong blue asymmetry measures. The RL population shows the same preference for red CIVλ1549 asymmetries (as was found for Hβ BC ) with A.I (M96; Sulentic et al 2000). A qualitative description of profile shapes ranges from extremely peaked to very boxy, which may be more common among RL AGN (e.g. the Hβ BC profiles of I Zw 1, and NGC 5548, respectively; see Figure 2). This description is based on examination of reasonably symmetric profiles. The presence of a strong asymmetry makes it difficult to assign a meaningful profile shape. These differences can be quantified by taking the ratio of the line width at different fractional intensities in a profile (for example, 1 4 and 3 peak intensity). The M96 4 curtosis parameter (defined as FW 3 4 M/FW 1 M) measures between 0.2 and 0.6 in 4 all populations. Considering the current reproducibility of such measures ( ±0.1), there is little sensitivity to possible trends. Also, curtosis may be intrinsically

13 PHENOMENOLOGY OF AGN EMISSION LINES 533 Figure 5 Emission-line profiles for (lower) NGC 5548, a RQ Population B Seyfert 1 source, and (upper) I Zw 1, a RQ Population A NLS1 source. Left panels show the respective CIVλ1549 BC profiles, and right panels show Hβ BC. Ordinate is flux in units of ergs s 1 cm 2 AA 1, and abscissa is radial velocity in km s 1 relative to the source rest frame (dot-dashed line). Dotted lines are the residual spectra after FeII opt was subtracted (upper right) or after the Hβ BC (thick solid line, all other panels) was subtracted. Note (lower left) that the apparent shift to the red of CIVλ1549 NC in NGC 5548 is a result of self-absorption. ambiguous if an inflection suggestive of multiple components is present in the line profile. Studies of profile wings are another way of characterizing profile shape (van Groningen & van Weeren 1989, Stirpe 1989, Penston et al 1990, Romano et al 1996). A sample of about 100 Hα line profiles (4 to 1 RL bias; Romano et al 1996, Robinson 1995) shows not only a mix of log and power-law functions needed to fit the profile wings but also profiles that frequently show significantly different red and blue wing shapes. The latter probably reflects the presence of multiple

14 534 SULENTIC ET AL Figure 6 Examples of sources with unusual Balmer line profiles after underlying continuum subtraction. Vertical lines at bottom of the lower panels indicate positions of contaminating narrow emission lines. Upper left: Arakelian 120, an apparently small-separation, double-peaked RQ source. Upper right: Arp 102B, a prototype broad separation double-peak RL source. Lower left: OQ 208, a single-peaked RL source with large-peak redshift. Lower right: Pictor A, an RL source Balmer lines that recently changed from single- to double-peak source. components, such as the redshifted VBC discussed earlier. Higher s/n profile samples, easily obtainable with the new generation of large telescopes (see Arav et al 1997, 1998), promise to greatly clarify the profile shapes and composite nature of broad-line profiles. Line profiles with the most unusual shape tend to be RL sources and include the rare objects that show double-peaked structure in the Balmer lines with wide peak separation v r 1000 km s 1, (e.g. Arp102B, see Figure 6; Chen et al 1989, Halpern et al 1996), 3C390.3 (Zheng et al 1991, O Brien et al 1998), and

15 PHENOMENOLOGY OF AGN EMISSION LINES 535 Pictor A (see Figure 6; Sulentic et al 1995). These objects show FWHM values approaching the mean FWZI values discussed earlier (EH94), indicating extreme boxiness. Double-peaked sources provide a special opportunity to search for a corresponding HIL component in a source where the LIL have a unique shape. Above referenced IUE and HST data for Arp 102B and 3C390.3 suggest that any HIL with the same shape as the double-peaked LIL profiles is quite weak [(I dp (CIVλ1549) 0.2 I(Hβ BC ); Halpern et al 1996)]. Instead, the CIVλ1549 BC is dominated by a single-peaked component that is, again, broader than the secondary single-peaked component seen in the LIL. Very broad (FWHM 10 4 km s 1 ) and double-peaked line profiles are primarily RL sources (EH94); however, at least two recent RQ examples of transient double-peaked Balmer lines are now known in sources that showed no previous BC (NGC1097: Storchi-Bergmann et al 1995, M81: Bower et al 1996). In RL Pictor A (Halpern & Eracleous 1994, Sulentic et al 1995), the double-peaked LIL structure appeared in a source previously showing only single-peaked BC profiles. 3. BASIC LINE-RELATED CORRELATIONS One of the golden rules of statistical analysis involves the necessity that two variables be experimentally independent of one another so that redundant, false, or mixed correlations can be avoided. For example, a correlation between apparent magnitude and redshift for a sample of quasars does not contain useful information about BLR physics. Another golden rule is that neither of the variables should contain some obvious source of variance that could create an apparent correlation when one is correlated against the other. Anticorrelations involving FWHM versus equivalent width for HIL lines (e.g. Francis et al 1992, Wills et al 1993, Corbin & Francis 1994) are driven by the relative strength of the narrow component in that line (i.e. the same correlation would be found between FWHM and (Hβ) if the NLR component were not subtracted). Correlations between FWHM Hβ BC and FWHM CIVλ1549 (e.g. Wilkes et al 1999) are driven by the relative strengths of the broad and narrow components in CIVλ1549. It is known that the narrow component of CIVλ1549, like [OIII]λλ4959,5007, is stronger in RL sources (e.g. Brotherton & Francis 1999). This means that CIVλ1549 will be measured systematically narrower in RL sources where Hβ BC is broadest. If CIVλ1549 NC is not subtracted, an anticorrelation will be found between FWHM (Hβ BC ) and FWHM (CIVλ1549). If there were a physically meaningful correlation between Hβ BC and CIVλ1549 BC (M96 concluded there is not for RQ AGN), it could be quantified only by comparing the measured values of the broad components of the two lines without additional sources of variance. Several correlations displayed in different contexts by different authors are driven by biased samples, i.e. samples for which the ranges covered by the independent variables are not sampled uniformly (e.g. Baldwin 1977a, Wang et al 1996, Corbin 1991, M96, Laor et al 1997a). Biased samples are extremely prone to spurious correlations. Consider, for example, the addition of two symmetrically

16 536 SULENTIC ET AL displaced points, each displaced by 5σ relative to the mean, in a sample of 30 uncorrelated observations that show a Gaussian distribution of two parameters. This situation will produce a formally significant correlation 99.9% of the time. If the points are located at ±4σ a significant correlation still results 91% of the time. This is not to say that correlations based on biased samples are necessarily erroneous, but assignment of a statistical significance to the correlation coefficient is invalid. Any statistical dependence based on a biased sample is of heuristic value. 3.1 Luminosity Correlations: The Baldwin Effect Baldwin (1977a) reported an anticorrelation between rest frame W(CIVλ1549) and the 1450 Å continuum luminosity W(CIVλ1549 α L 2 3 ν ) (the Baldwin effect, hereafter BE). The original Baldwin (1977a) sample of 20 quasars was dominated by flat spectrum RL with a continuum luminosity range of ergs s 1 Hz 1. The relation was soon confirmed with a rather small rms scatter of 0.6 magnitude (Baldwin et al 1978) and later with two larger flat spectrum samples (Wampler et al 1984). Baldwin et al (1989) found W L 1 3 ν for a sample of BQS and PKS quasars. Combining five RQ and RL samples, Kinney et al (1990) found evidence for an even flatter correlation in the range log L ν (six decades in luminosity) with a factor of 10 change in W(CIVλ1549) ( 250 to 25 Å in the rest frame). Multiple observations for 5 RQ and 2 RL sources in that sample showed a source-specific anticorrelation for each source, with an average slope of 0.64, almost identical to the value originally found by Baldwin (1977a). This work led to the concept of two Baldwin effects: (1) a global BE with small slope, W L 0.18 ν, and significant only over a wide range of luminosity (L ν, max /L ν, min 10 2 ); and (2) an intrinsic BE reflecting the change in source equivalent width with change in continuum luminosity. The intrinsic BE contributes significantly to the scatter in the global BE (e.g. Kenney et al 1990, Pogge & Peterson 1992). More recent studies have usually confirmed a weak effect in large samples of optically selected AGN (Cristiani & Vio 1990, Francis et al 1992, Tytler & Fan 1992, Osmer et al 1994, Francis & Koratkar 1995, Laor et al 1995, Wang et al 1998, Chavushyan 1995) with a slope of Detection of the BE in small samples has remained controversial, with several negative results (Steidel & Sargent 1991, Wilkes et al 1999, Wills et al 1999). Attention has also been focused on other UV lines that may follow a trend similar to W(CIVλ1549). A BE has been reported in Lyα (Baldwin 1977a, Cristiani & Vio 1990, Kinney et al 1990, Green 1996), CIII]λ1909 (Tytler & Fan 1992, Green 1996, Steidel & Sargent 1991), HeIIλ1640 (Tytler & Fan 1992), HeIIλ4686 (Heckman 1980, Boroson & Green 1992-their Eigenvector 2, Zheng & Malkan 1993, Green 1996) and MgIIλ2800 (Baldwin 1977a, Tytler & Fan 1992, Thompson et al 1999, Steidel & Sargent 1991). A controversial detection for Nvλ1240 (Tytler & Fan 1992) was not subsequently confirmed (Osmer et al 1994, Steidel & Sargent 1991, Laor et al 1995). Another controversial BE involves the λ1400 blend with reports of no detection (Osmer et al 1994), and a detection that is apparently the strongest luminosity

17 PHENOMENOLOGY OF AGN EMISSION LINES 537 anticorrelation among the UV lines of Laor et al Zheng et al (1995) reported a BE for the HIL OVIλ1034 that was not confirmed by Green (1996), while Laor et al (1995) found a shallow trend (0.15 ± 0.08). More data, however, suggested that the strength of the BE was related to the ionization potential of the relevant ion (Espey et al 1993, Lanzetta et al 1993, Zheng & Malkan 1993): the OVIλ1034 line shows the strongest BE (steepest slope) of any previously examined line, with MgIIλ2800 and the Balmer lines showing the weakest (or no) effect. These results have led to a standard scenario in which the BE occurs in all measurable HIL except NVλ1240, and slope increases with ionization potential (e.g. Osmer & Shields 1999; see Hamann & Ferland 1999 for the implications of the apparent lack of any BE for NVλ1240). No convincing BE has been reported for the Balmer lines or FeII opt (Espey et al 1989, Osterbrock 1978, Shuder 1981, Yee & Oke 1981, BG92, Elston et al 1994). Rather, high-redshift RQ quasars that have been observed show evidence of strong FeII opt and FeII UV (Hill et al 1993, Elston et al 1994). Observations of 12 high-redshift quasars with z suggest that the FeII( Å)/MgIIλ2800 ratio is approximately the same and possibly slightly increasing relative to the LBQS (Francis et al 1992, Thompson et al 1999; see Kawara et al 1996 for a similar result on B , a RL AGN at z 3.6). Kinney et al (1990) and Zheng & Malkan (1993) found no difference in the BE for RQ and RL subsamples. On the contrary, Osmer et al (1994) suggested that RL sources have a more negative power-law index than optically selected samples, although they concluded that the difference in slope was not significant. The BQS and PKS samples show an important difference: the correlation coefficient between W(CIVλ1549) and continuum luminosity is significantly nonzero for the PKS but not for the BQS sources (Baldwin et al 1989). These last results suggest that RQ and RL AGN should be considered separately in BE studies BE Selection Effects A BE for core-dominated quasars was criticized as a selection effect because these sources show strong intrinsic continuum variability (Murdoch 1983). Indeed, it is possible to ascribe the original BE entirely to selection effects. Core-dominated objects are a small, biased subset of a minority quasar population (RL AGN). The correlation may be stronger in small, radio-selected samples because of the increased probability of finding beamed sources. This leads to a BE-like correlation for sources with varying degrees of continuum beaming. Even the Kinney et al (1991) BE could be the result of a statistical bias. The weak, shallow anticorrelation produces a difference that is comparable in size to the observed scatter in W(CIVλ1549) at any given luminosity. The intrinsic BE that is related to source continuum variability may enhance the scatter in UV luminosity. A spurious correlation could also arise if the dispersion in continuum luminosity is much larger than the dispersion in line luminosity (Yuan et al 1998a). We made Monte-Carlo simulations (assuming a uniform distribution over the same equivalent width-luminosity range as for Kinney et al 1991) to test the

18 538 SULENTIC ET AL likelihood of finding a BE in small samples. Our results suggest that randomly extracted samples of 20 objects covering two decades in luminosity will show a significant BE-like correlation 25% of the time, always with a slope larger than that of the dataset from which they were extracted. Thus, the alternating positive and negative results reported above are not surprising. If we add the data obtained by Wu et al (1983) and Wilkes et al (1999) (negative results) into the Kinney et al (1990) diagram, we find that they fit the BE derived there. However, there are two important caveats; the reality of the BE depends on the real absence of (1) low EW and luminosity AGN at low z and (2) high EW and high luminosity AGN at high z. These objects could be underrepresented because of biases in selection techniques. Magnitude-limited high-redshift samples are biased toward intrinsically over-luminous quasars, which will be biased toward below-average line equivalent widths (see Section 4.7 for a possible physical interpretation). At the other extreme, low-luminosity NLS1 with low W(CIVλ1549) are found in significant numbers (e.g. Grupe et al 1999; Brotherton & Francis 1999). They have a tendency to blur the Baldwin effect at the low-luminosity end. In addition, the Kinney et al (1990) results were based on IUE spectra of AGN that often detected the UV continuum slightly above noise, a fact that tends to increase the observed W(CIVλ1549). It is important that HST archival data be used in confirmatory studies. Another issue involves the intensity of the narrow core in the CIVλ1549 line, which contributes non-negligibly to W(CIVλ1549), and which may be redshift (luminosity) dependent. It is not clear whether this is the case (Steidel & Sargent 1991, Osmer et al 1994, Brotherton & Francis 1999) but this component should not be included in W(CIVλ1549) measures evaluating the BE. This narrow component is also stronger in RL and X-ray selected AGN (see e.g. Chun et al 1999, Laor et al 1997b, Green 1998), which further heightens the danger of bias when mixing RQ RL samples. 3.2 Eigenvector 1 Correlations Boroson & Green (1992) identified a series of related correlations from principal component analysis (PCA) of their sample correlation matrix. The measures included FeII opt,[oiii]λλ4959,5007, Hβ BC, logr (defined as ratio of 6 cm to 4400 Å flux densities; Kellermann et al 1989), and the optical to X-ray spectral index, α ox. Their sample contained 87 PG quasars with z 0.5 (17 of them radio-loud). The first PCA eigenvector (hereafter E1) strongly anticorrelates with W(FeIIλ4570)/W(Hβ BC ) and luminosity of [OIII]λλ4959,5007. In an effort to clarify the meaning of E1, we have searched for the correlation diagram that shows maximal discrimination between the various AGN subclasses. The best E1 correlation space that we can identify involves (1) FeII λ4570 strength, defined as the ratio R Fell = W(FeIIλ4570)/W(Hβ BC ), and (2) FWHM(Hβ BC ), supplemented by (3) the soft X-ray photon index, Ɣ soft (measured between 0.8 and 2.4 KeV on ROSAT spectra). Appreciation of the significance of this 3D correlation space has been growing over the past few years (see e.g. Wang et al 1996, Laor et al 1997b, Lawrence et al 1997, Brotherton & Francis We focus on the BG92

19 PHENOMENOLOGY OF AGN EMISSION LINES 539 sample because it is nearest to a complete sample of quasars with state-of-the-art resolution and s/n measures of Hβ BC and FeII opt. We find that variables such as W(FeII λ4570) and W(Hβ BC ) show weak or illdefined correlations with FWHM(Hβ BC ). The lower right panel of Figure 7 shows a plot of W(FeII λ4570) versus FWHM(Hβ BC ) for the BG92 quasars. Correlation coefficients found for W(FeII λ4570) versus FWHM(Hβ BC ) include r 0.44 (BG92), r 0.54 (Wang et al 1996) and r 0.69 (Corbin & Boroson 1996) with considerable overlap between these samples. The diagram is obviously not a simple correlation and is best described as two clusters of sources that we will call Populations A and B. Population A [FWHM(Hβ BC ) 4000 km s 1 ]is an almost pure RQ phenomenology centered near FWHM(Hβ BC ) 2200 km s 1 and W(FeII λ4570) 60 Å with approximate ranges km s 1 and Å. This population includes 63% of the RQ sample. Sources within Population A show no correlation between W(FeII opt ) vs. FWHM(Hβ BC ). Population B [FWHM(Hβ BC ) 4000 km s 1 ] is centered near FWHM(Hβ BC ) 5500 km s 1 and W(Hβ BC ) = 30 Å with approximate ranges km s 1 and 0 60 Å. This population includes 26% of the BG92 RQ sample. There is no significant correlation within this population. The remaining BG92 sources scatter over a wide range of parameter space: W(FeIIλ4570) Å and FWHM(Hβ BC ) ,000 km s 1. There may be a continuous distribution of RQ sources with FeII opt decreasing as FWHM(Hβ BC ) increases; however, the Population A and B distinction is useful because it identifies those RQ quasars that show optical properties most similar to the bulk of RL sources (see Table 2; see also Section 4.5). TABLE 2 FWHM Hβ b WHβ d W FeII e Ɣ soft g W [OIII] h AGN a (km 1 ) N c (Å) (Å) RFB f Photon Index (Å) A-RQ ± ± ± ± 0.4 (17) 15 ± 12 A RQ ± ± ± ± 0.2 (17) 28 ± 38 A RQ ± ± ± ± 0.1 (8) 20 ± 20 B RQ ± ± ± ± 0.2 (14) 33 ± 26 RL ± ± ± ± 0.2 (3) 11 ± 6 RL ± ± ± ± 0.3 (7) 27 ± 15 RL ± ± ± ± 0.5 (17) a ABN sample type. A and B refer to populations A and B. RQ and RL refer to radio-quiet and radio-loud sources. b The arbitrarily chosen range of FWHM. c The number of sources. d Mean equivalent width and standard deviation for the broad component of Hβ. e Mean equivalent width and standard deviation for the optical FeIIλλ Å blend. f The ratio column 5/column 4. g The soft X-ray spectral index from sources given in the text. Value derived with N H as a free parameter when available. h Mean equivalent width and standard deviation for [OIII]λ5007 Å.

20 540 SULENTIC ET AL The R FeII parameter shows a significant correlation with FWHM(Hβ BC ) (Gaskell 1985, Zheng & O Brien 1990, r 0.55 in BG92, r 0.74 in Wang et al 1996). The lack of any significant correlation between W(Hβ BC ) and FWHM (Hβ BC )(r +0.18) in the BG92 sample lessens concerns about correlate crosstalk in this part of E1. W(FeII λ4570) will be measured systematically low relative to W(Hβ BC ) because the average continuum increases toward the blue. (M96 included FeII blends on both side of Hβ, but this choice also has drawbacks). Figure 7 shows 2D projections of the FWHM (Hβ) versus R Fell, Ɣ soft versus R Fell, and FWHM(Hβ BC ) versus Ɣ soft. We have supplemented the BG92 RL sample with an additional 18 sources (comparable s/n spectra) taken from M96 with W(FeII opt ) measures converted to W(FeII λ4570) as in by BG92. It has been known for some time that one can distinguish between RL and RQ quasars on the basis of both hard (Piccinotti et al 1982, Brandt et al 1997) and soft X-ray measures (Wang et al 1996, Brinkmann et al 1997a). Soft X-ray photon indices are available for a large fraction of the BG92 sample (Boller et al 1996, Brinkmann et al 1997b, Siebert et al 1998, Wang et al 1996, Yuan et al 1998b). For Figure 7, we have preferred values derived with N H taken as a free parameter. Early claims of a correlation between Ɣ soft and W(FeII λ4570) emission (Wilkes et al 1987) were challenged (Boroson 1989, Zheng & O Brien 1990, Walter & Fink 1993), but a significant correlation is clearly present in the larger BG92 sample (see also Shastri et al 1993). A better-defined correlation was found between Ɣ soft and R FeII (Puchnarewicz et al 1992; Laor et al 1994, 1997b; Wang et al 1996; Boller et al 1996; Grupe et al 1999). Previous estimates of the correlation strength yielded (1) Ɣ soft versus FWHM(Hβ BC ): r 0.73; and (2) Ɣ soft versus R FeII :r 0.65 (Wang et al who did not distinguish between RL and RQ sources in their correlation analysis). A significant difference between the R FeII versus FWHM(Hβ BC ) and W(FeII λ4570) versus FWHM(Hβ BC ) correlations (Figure 7) is found because Population A quasars show a much larger scatter along the R FeII axis. Population A shows a much larger range in R FeII ( ) than Population B (R FeII ). The objects in the high R FeII tail are often objects with unusually low W(Hβ BC ) rather than unusually strong W(FeII λ4570). Low W(Hβ) in narrow line Seyfert 1 nuclei (see next section) was recognized a long time ago (Osterbrock & Pogge 1985, Gaskell 1985). Objects with R FeII 0.9 show mean W(Hβ) 67 Å while those with R FeII 0.9 show W(Hβ) 114 Å, almost twice as large. The W (Hβ BC ) range in RQ Population B is also large: W(Hβ) = Å and W(FeII λ4570) = 0 69 Å, respectively. However, the ratio R FeII is surprisingly constant, with R FeII > 0.3 ± 0.25 which is similar to the mean for RL sources. This is a justification for retaining our Population A-B distinction. We find no significant difference in mean RQ optical luminosity between Populations A and B, which is consistent with the ortogonality between E1 and luminosity implied by the BG92 PCA analysis. Table 2 summarizes some relevant mean parameters for RQ and RL sources shown