The Radius-Luminosity Relationship Depends on Optical Spectra in Active Galactic Nuclei
DD RAFT VERSION O CTOBER
1, 2019Typeset using L A TEX twocolumn style in AASTeX62
The Radius-Luminosity Relationship Depends on Optical Spectra in Active Galactic Nuclei P U D U AND J IAN -M IN W ANG
1, 2, 31
Key Laboratory for Particle Astrophysics, Institute of High Energy Physics, Chinese Academy of Sciences, 19B Yuquan Road, Beijing 100049, China National Astronomical Observatories of China, Chinese Academy of Sciences, 20A Datun Road, Beijing 100020, China School of Astronomy and Space Science, University of Chinese Academy of Sciences, 19A Yuquan Road, Beijing 100049, China
ABSTRACTThe radius-luminosity ( R H β – L ) relationship of active galactic nuclei (AGNs) established by the rever-beration mapping (RM) observations has been widely used as a single-epoch black hole mass estimator in theresearch of large AGN samples. However, the recent RM campaigns discovered that the AGNs with high ac-cretion rates show shorter time lags by factors of a few comparing with the predictions from the R H β – L relationship. The explanation of the shortened time lags has not been finalized yet. We collect 8 different single-epoch spectral properties to investigate how the shortening of the time lags correlate with those properties andto understand what is the origin of the shortened lags. We find that the flux ratio between Fe II and H β emissionlines shows the most prominent correlation, thus confirm that accretion rate is the main driver for the shortenedlags. In addition, we establish a new scaling relation including the relative strength of Fe II emission. This newscaling relation can provide less biased estimates of the black hole mass and accretion rate from the single-epochspectra of AGNs. Keywords: galaxies: active; galaxies: nuclei - quasars: supermassive black holes INTRODUCTIONIn the past 40 years, reverberation mapping (RM; e.g.,Bahcall et al. 1972; Blandford & McKee 1982; Peterson etal. 1993) has become a powerful tool to investigate the ge-ometry and kinematics of the broad-line regions (BLRs) inactive galactic nuclei (AGNs) and to measure the masses ofsupermassive black holes (BHs). Through long-term spec-troscopic monitoring of an AGN, the size scale ( R BLR ) ofits BLR can be directly obtained by measuring the delayedresponse ( τ BLR ) of the emission line with respect to the vari-ation of the continuum, where R BLR = cτ BLR and c is thespeed of light. Fortunately, the RM observations of ∼ objects (e.g., Peterson et al. 1993, 1998, 2002, 2004; Kaspi etal. 2000, 2007; Bentz et al. 2008, 2009a; Denney et al. 2009;Barth et al. 2011, 2013, 2015; Rafter et al. 2011, 2013; Duet al. 2014, 2015, 2016a, 2018a,b; Wang et al. 2014a; Shenet al. 2016b; Fausnaugh et al. 2017; Grier et al. 2012, 2017b;De Rosa et al. 2018; Woo et al. 2019) lead to a correlation be-tween the time lag of H β emission line (or the radius R H β ofthe H β -emitting region) and the monochromatic luminosity( λL λ ) at 5100 ˚A (hereafter L ) with the form of R H β = α(cid:96) β , (1)where (cid:96) = L / erg s − (e.g., Kaspi et al. 2000;Bentz et al. 2009b, 2013). This correlation makes it possi-ble to estimate BLR radius from a single-epoch spectrum. Itis called the R H β – L relationship, and has been widelyadopted as a single-epoch BH mass estimator in the researchof large AGN samples (e.g., McLure & Dunlop 2004; Vester-gaard & Peterson 2006; Kollmeier et al. 2006; Greene & Ho 2007; Shen et al. 2011). However, there is growing evidencefor increasing scatters of the R H β – L relationship fromongoing RM campaigns.Reverberation of broad emission lines to the continuumconfirms photoionization as the major radiation mechanismin the BLR. As a canonic model of BLR, photoionizationdefined by the ionization parameter Ξ ∝ L ion /R n e T di-rectly indicates R BLR ∝ L / , where L ion is the ionizingluminosity, n e and T are electron density and temperature ofthe ionized gas, respectively. If we take L as a proxyof L ion , we have R BLR ∝ L / , agreeing with the obser-vations (see also Bentz et al. 2013). We would like to em-phasize here the necessary conditions for this canonic rela-tion: (1) ionizing source should be isotropic or at least quasi-isotropic so that the BLR clouds receive the same luminos-ity with observers; (2) L ion ∝ L should always work;(3) ionizing luminosity comes from a point source, which ismuch smaller than the distances of the BLR clouds to thecentral BH. Condition (1) is broken in the AGNs powered byslim accretion disks (Wang et al. 2014c), where the puffed-up inner region may lead to non-isotropic ionizing radiation(Wang et al. 2014c). Condition (2) relies on spectral energydistributions (only holds for pro-grade accretion AGNs pow-ered by the Shakura-Sunyaev disks), and does not work inones with retro-grade accretion (Wang et al. 2014b; Czernyet al. 2019). Low spin (also low accretion rate and large BHmass) may lead to the deficit of the UV photons and non-linear relation between L ion and L (Wang et al. 2014b;Czerny et al. 2019). And the L variation shows a littlelag with respect to the L ion variation in the RM of accretion a r X i v : . [ a s t r o - ph . GA ] S e p ˚ A)0 . . . . . . . . F l u x ( − e r g s − c m − ˚ A − ) [ O III ] H β H e II F e II F e II SDSS J081456 4500 5000 5500 6000 6500Wavelength ( ˚ A)0 . . . . . . . . F l u x ( − e r g s − c m − ˚ A − ) [ O III ][ O III ] H β H e II F e II F e II Mrk 1310
Figure 1.
Two fitting examples. In each panel, the black line is the spectrum in the rest frame after Galactic extinction correction. The red lineis the best fit. The orange line is the power law of the continuum. The purple line is the broad H β component. The blue lines are the narrowemission lines ([O III ] λλ , , H β , and He II ). The green, brown, and yellow lines are the Fe II template, broad He II , and the templateof host galaxy, respectively. The zoom-in panels show the detailed fitting around H β . Some strong emission lines are labeled. disks (Edelson et al. 2015; McHardy et al. 2018; Cackett etal. 2018, e.g.,). About Condition (3), the size of accretiondisk, although small, has been successfully measured and isnot infinitesimal (Edelson et al. 2015; McHardy et al. 2018;Cackett et al. 2018). Therefore, the R H β – L relationshipis expected to depend on accretion situation or some otherproperties.Recently, Super-Eddington Accreting Massive Black Hole(SEAMBH) campaign discovered that many objects withstrong Fe II and narrow H β emission lines, which are thoughtto be the AGNs with high accretion rates, lie below the R H β – L relationship (Du et al. 2015, 2016a, 2018a). Theyfound that the time lags of the AGNs with high accretionrates become shortened by factors of ∼ relative to thenormal-accretion-rate AGNs with the same luminosities, andthe shortening itself shows correlation with the accretion rate(Du et al. 2015, 2016a, 2018a). Wang et al. (2014c) proposedthat the anisotropic radiation of the slim accretion disk mayprobably result in the shortened time lags. The Sloan Digi-tal Sky Survey Reverberation Mapping (SDSS-RM) Projectalso reported many AGNs have time lags shorter than ex-pected from the R H β – L relationship (Grier et al. 2017b),but cautioned that selection effects may arise at least in somecases (see Grier et al. 2019).Although the detailed physical explanation causing theshortened time lags is not yet finalized (Wang et al. 2014c;Grier et al. 2017b, 2019), more and more objects deviatingfrom the traditional R H β – L relationship are being dis-covered (Grier et al. 2017b; Du et al. 2018a). It is urgentto investigate the origin of the shortened lags in more detail.In this paper, we investigate how the deviation of an AGNfrom the R H β – L relationship correlates with the proper-ties in the single-epoch spectrum, and try to establish a newscaling relationship including the influence of single-epochspectral properties. We describe the sample, data, and mea- surements in Section 2. A new scaling relation is establishedand presented in Section 3. Some discussions are providedin Section 4, and a brief summary is given in Section 5. Weadopt the standard Λ CDM cosmology and the parameters of H = 67 km s − Mpc − , Ω Λ = 0 . , and Ω m = 0 . (Planck Collaboration et al. 2014, 2018) in this paper. DATA AND MEASUREMENT2.1.
Sample
The analysis in the present paper is mainly based on thesamples: (1) the RM measurements compiled in Bentz etal. (2013) from the previous literatures, (2) the AGNs withhigh accretion rates of the SEAMBH campaign published inDu et al. (2014, 2015, 2016a, 2018a); Wang et al. (2014a);Hu et al. (2015), (3) some other AGNs published after 2013:Mrk 1511 from Barth et al. (2013), NGC 5273 from Bentz etal. (2014), KA 1858+4850 from (Pei et al. 2014), MCG +06-30-015 from (Bentz et al. 2016b; Hu et al. 2016), UGC 06728from (Bentz et al. 2016a), and MCG +08-11-011, NGC 2617,3C 382 and Mrk 374 from Fausnaugh et al. (2017) . The col-lection in Bentz et al. (2013) includes 41 AGNs monitoredsuccessively since the late 1980’s, most of which have rela-tively weaker Fe II emission and broader H β lines comparedto the SEAMBH objects (Du et al. 2018a). The SEAMBHcampaign, as a dedicated RM project for AGNs with highaccretion rates, published the time lags of 25 objects (totally30 measurements, some objects have more than one measure-ment). Including the objects published after 2013, we havetotally 75 objects with 117 measurements. The time lags, We also include the new RM observations of the previous mapped ob-jects after 2013 in the following analysis: NGC 4593 (Barth et al. 2013),NGC 7469 (Peterson et al. 2014), NGC 5548 (Lu et al. 2016; Pei et al.2017), NGC 4051 (Fausnaugh et al. 2017), PG 1226+023 (3C 273, Zhang etal. 2019), and PG 2130+099 (Hu et al. 2019) β , and [O III ] luminosities, equivalent width (EW)of H β and [O III ], and H β FWHM of the corresponding cam-paigns are listed in Table 1. Some objects have been mappedmore than once, in order to understand the population proper-ties better, we have to equalize the weights of the individualobjects in the following analysis. We average the multiplemeasurements by taking into account their measurement un-certainties (see more details in Du et al. 2015). The averagemeasurements, etc. lags, FWHM, luminosities, are also listedin Table 1.2.2.
Data of Single-epoch Spectral Properties
To investigate how the single-epoch spectral propertiescontrol the deviation of AGNs from the R H β – L rela-tionship, we compile 8 different parameters from the spectraaround the H β region in the optical band: (1) the flux ratiobetween Fe II and H β , which is denoted as R Fe = F Fe /F H β , (2)where F Fe is the flux of Fe II from 4434 ˚A to 4684 ˚A and F H β is the flux of broad H β , (2) FWHM of H β emissionline ( FWHM H β ), (3) equivalent width (EW) of [O III ] λ emission line ( EW [OIII] ), (4) the ratio between FWHM and σ line (second moment of the line profile) of the H β line( D H β = FWHM H β /σ H β ), (5) the ratio between the FWHMof Fe II and the FWHM of H β ( FWHM
Fe II / FWHM H β ),(6) the asymmetry of H β line defined by A = [ λ c (3 / − λ c (1 / / FWHM H β , where λ c (3 / and λ c (1 / are thecentral wavelengths at the 3/4 and 1/4 of the H β peak height(De Robertis 1985; Boroson & Green 1992; Brotherton 1996;Du et al. 2018b), (7) the EW of He II ( EW HeII ), and (8) theEW of H β ( EW H β ). These parameters are referred to as“single-epoch spectral properties”, because they can be mea-sured simply from the single-epoch spectra rather than fromthe time-domain observations like RM . In order to estab-lish some new scaling relationships which can be applied tothe large AGN samples obtained in the spectroscopic surveysof SDSS or Dark Energy Spectroscopic Instrument (DESI)in the near future, we need to find the correlations betweenthe single-epoch properties and the deviation of AGNs fromthe R H β – L relationship. The FWHM and EW of the H β line, and the EW of the [O III ] line of each RM campaign arecollected and listed in Table 1. We use them (and the aver-age values for the objects with multiple RM measurements)directly in the following analysis. We search the other param-eters in the literatures and list the values in Table 2. For theparameters that we can not find in the literatures, we fit thespectra of the objects found in the public archive and measurethose parameters by ourselves (see also Table 2).2.3.
Fitting the spectra Of course, if we have the RM data, we can definitely measure them froman individual spectrum in the RM campaign or the mean spectrum (can betreated as the average of the values from the individual spectra). But if wedon’t have the RM data, we can still measure them from the single-epochspectra found in some other literatures or databases.
We use the following components in the spectral fitting: (1)a power law to model the AGN continuum, (2) two Gaussiansto model the broad H β emission line, (3) a template con-structed from the Fe II spectrum of the narrow-line Seyfert 1(NLS1) galaxy I Zw 1 by Boroson & Green (1992) for the Fe II emission, (4) one or two Gaussians for each of the narrowemission lines, e.g., [O III ] λλ β , He II (if neces-sary), (5) one or two Gaussians to model the broad He II and(6) a simple stellar population model from Bruzual & Char-lot (2003) as a template for the contribution of the host galaxyif necessary. The fitting is mainly performed in the windowsof 4170–4260 ˚A and 4430–5550 ˚A in the rest frame. If thehost contribution is significant, we supplement the window of6050–6200 ˚A to give a better constraint to the fitting of stellartemplate. All of the narrow-line components in each objectare fixed to have the same velocity width and shift, exceptfor those showing very different width/shift. [O III ] λ III ] λ β ). In the fit-ting for the spectra of NLS1s, we also fix the flux of narrowH β to be one-tenth of the [O III ] λ ANALYSIS3.1.
Pairwise Correlations between Different Properties
Before discussing the correlations between the single-epoch spectral properties and the deviation from the R H β – L relationship, we first present the pairwise correlationsbetween different properties in Figure 2. Although thereare many similar discussions using different samples in thehistorical literatures (e.g., Boroson & Green 1992), it is stillvaluable to do this demonstration for the RM objects. Spear-man’s rank correlation coefficients ( ρ ) and the correspondingtwo-sided p -value for a null hypothesis test (two sets of dataare uncorrelated) are marked in the panels of Figure 2. The ρ and p values of the significant correlations (with p < . )are marked with red color.Among all of the correlations, FWHM H β versus D H β isthe most significant one, which has Spearman’s correlationcoefficient ρ = 0 . . It means that if the width of H β lineis smaller, its profile tends to be more Lorentzian-like. Thiscorrelation has been demonstrated by, e.g., Kollatschny &Zetzl (2011, 2013). The sample in the present paper is largerthan that used in Kollatschny & Zetzl (2011, 2013), but theresult is almost the same. The secondarily-significant corre-lations are R Fe versus FWHM
FeII / FWHM H β ( ρ = 0 . ), We select the simple stellar population model which can get the smallest χ in the fitting. R Fe E W H β ( ˚ A ) ρ = − . p =1 . × − FWHM H β (km s − ) ρ =0 . p =0 . × − EW [OIII] ( ˚ A) ρ =0 . p =2 . × − D H β ρ =0 . p =1 . × − . . FWHM
FeII / FWHM H β ρ = − . p =0 . × − − .
25 0 .
00 0 . Asymmetry ρ = − . p =1 . × − EW HeII ( ˚ A) ρ =0 . p =0 . × − E W H e II ( ˚ A ) ρ = − . p =2 . × − ρ = − . p =2 . × − ρ =0 . p =0 . × − ρ =0 . p =0 . × ρ = − . p =0 . × ρ = − . p =0 . × − − . . . . A s y mm e tr y ρ =0 . p =2 . × − ρ = − . p =1 . × − ρ = − . p =0 . × − ρ = − . p =2 . × − ρ =0 . p =2 . × − . . F W H M F e II / F W H M H β ρ =0 . p =0 . × − ρ = − . p =2 . × − ρ = − . p =2 . × − ρ = − . p =0 . × − D H β ρ = − . p =2 . × − ρ =0 . p =0 . × − ρ =0 . p =0 . × − E W [ O III ] ( ˚ A ) ρ = − . p =3 . × − ρ =0 . p =0 . × − F W H M H β ( k m s − ) ρ = − . p =1 . × − Figure 2.
Pairwise correlations of the single-epoch spectral properties for the sample in the present paper. The Spearman’s rank correlationcoefficients ( ρ ) and the corresponding null probabilities (see Section 3.1) are marked in the corner of each panel (with red color if p < . ). R Fe versus H β Asymmetry ( ρ = 0 . ), R Fe versus EW H β ( ρ = − . ), R Fe versus FWHM H β ( ρ = − . ), and FWHM H β versus FWHM
Fe II / FWHM H β ( ρ = − . ).The previous three correlations mean that, if the relativestrength of Fe II is higher, the widths of Fe II and H β aremore similar, H β line tends to have stronger blue wing, andthe EW of H β line is weaker. Similar to the correlation be-tween R Fe and FWHM
Fe II / FWHM H β , FWHM H β versus FWHM
Fe II / FWHM H β means that the objects with nar-rower H β lines also have more similar Fe II and H β widths. In addition, the correlations of R Fe versus EW [OIII] , R Fe versus D H β , FWHM H β versus H β Asymmetry, EW [OIII] versus EW HeII , and
FWHM
Fe II / FWHM H β versus H β Asymmetry are also significant. R Fe versus ( FWHM H β and EW [OIII] ) are the prominent correlations in the famous AGNeigenvector 1 sequence (e.g., Boroson & Green 1992; Su-lentic et al. 2000; Marziani et al. 2001, 2003, 2018; Shen &Ho 2014; Sun, & Shen 2015). The detailed physical processof this sequence is still under some debate (e.g., Panda etal. 2019). The correlation between R Fe versus H β Asym- R H β ( l t − d a y s ) . . . . R F e R H β ( l t − d a y s ) F W H M H β ( k m s − ) R H β ( l t − d a y s ) . . . . . . l og E W [ O III ] ( ˚ A ) R H β ( l t − d a y s ) . . . . . D H β R H β ( l t − d a y s ) . . . . . F W H M F e II / F W H M H β R H β ( l t − d a y s ) − . . . A s y mm e tr y L (erg s − ) R H β ( l t − d a y s ) E W H e II ( ˚ A ) L (erg s − ) R H β ( l t − d a y s ) . . . . . l og E W H β ( ˚ A ) Figure 3.
The R H β – L relationship color-coded by the spectral properties. The dotted lines are the R H β – L relationship for the low-accretion-rate AGNs in Du et al. (2018a). It is obvious that some objects deviate from the dotted lines. The colors show clear trend with R Fe and EW H β , which means the deviation correlates with R Fe and EW H β . metry has been demonstrated using the PG quasar samplein Boroson & Green (1992), and is also associated with theeigenvector 1 sequence. The RM sample reproduces thiscorrelation. Besides, there are some other weak correlations,please see Figure 2. More discussions about the pairwisecorrelations are provided in Section 4.3.2. Deviation from the R H β – L Relationship
In Figure 3, we first show the R H β – L relationshipcolor-coded by the properties we collected. The deviationfrom the R H β – L relationship shows clear correlationwith R Fe and EW H β , both of which show obvious varia-tion trend across the R H β axis in Figure 3. In order to furtherinvestigate the significance of the correlations, we define thedeviation from the R H β – L relationship as ∆ R H β = log ( R H β /R H β, R − L ) , (3) R Fe − . − . − . − . . . . ∆ R H β = l og ( R H β / R H β , R − L ) ρ = − . p =2 . × − FWHM H β (km s − ) ρ =0 . p =2 . × − EW [OIII] ( ˚ A) ρ =0 . p =0 . × − D H β ρ =0 . p =0 . × − . . . FWHM
FeII / FWHM H β − . − . − . − . . . . ∆ R H β = l og ( R H β / R H β , R − L ) ρ = − . p =0 . × − − . . . . Asymmetry ρ = − . p =0 . × − EW HeII ( ˚ A) ρ =0 . p =0 . × − EW H β ( ˚ A) ρ =0 . p =1 . × − Figure 4.
The correlations between ∆ R H β and the properties for the sample in the present paper. The Spearman’s rank correlation coefficients( ρ ) and the corresponding null probabilities (see Section 3.1) are marked in the corner of each panel (with red color if p < . ). R Fe showssignificant correlation with ∆ R H β , and EW H β shows a moderate correlation. The grey dashed lines show ∆ R H β = 0 . The black dotted lineand the orange color show the linear regression and its confidence band (2 σ ). The two grey points are MCG +06-26-012 and MCG +06-30-015(see more details in Section 3.3). We correct the intrinsic reddening MCG +06-30-015 and use its corrected luminosity in the analysis. We donot include MCG +06-26-012 in the analysis. where R H β, R − L is the prediction from the R H β – L rela-tionship. Here, we adopt log R H β, R − L = 1 .
53 + 0 .
51 log (cid:96) obtained by Du et al. (2018a) for the AGN with dimension-less accretion rate ˙ M < ( ˙ M is defined by the followingEquation (8), please see Section 4.3) as the fiducial R H β – L relationship. It should be noted that using the R H β – L relationship in Bentz et al. (2013) doesn’t change thediscussion and conclusion in this paper.The correlations between ∆ R H β and the single-epochproperties are shown in Figure 4. Again, the Spearman’scorrelation coefficients and the corresponding null proba-bilities are marked in the corners of the panels in Figure 4.The correlation between ∆ R H β and R Fe is the most signif-icant one. ∆ R H β shows a strong anti-correlation with R Fe with Spearman’s correlation coefficient ρ = − . and thenull probability p = 2 × − . ∆ R H β and EW H β showsa weaker correlation with ρ = 0 . and p = 1 × − .In addition, the low- D H β or small- FWHM H β objects showmore extended distribution of ∆ R H β , while the high- D H β orlarge- FWHM H β objects have relatively narrower ∆ R H β distribution and the average ∆ R H β more close to zero.But the Spearman’s coefficients of ∆ R H β versus D H β and FWHM H β are not high enough. Actually, this complex dis-tribution in the ∆ R H β versus FWHM H β (or D H β ) planemay be caused by the eigenvector 1 sequence (the corre-lation between R Fe and FWHM H β , see, e.g., Boroson & Green 1992; Sulentic et al. 2000; Marziani et al. 2001, 2003;Shen & Ho 2014; Sun, & Shen 2015, or the recent review inMarziani et al. 2018). The large- FWHM H β (or high- D H β )objects have low R Fe values, but the small- FWHM H β (orlow- D H β ) objects have large range of R Fe (span from low tohigh R Fe ). Therefore, the ∆ R H β distribution in the small- FWHM H β (or low- D H β ) objects is more extended. Somemore discussions about this are provided in Section 4.4. Thecorrelations between ∆ R H β and the other parameters arenot significant given the present data ( ∆ R H β shows weaka correlation with FWHM
Fe II / FWHM H β ). Because thecorrelation coefficient of R Fe is the highest among all ofthe spectral properties, R Fe can be regarded as the primaryparameter that controls the deviation of an AGN from the R H β – L relationship.The eigenvector 1 sequence (or main sequence) of AGNshas been extensively investigated in the past decades, andcontains the information of the evolution or systematic varia-tion of AGNs (see the recent review in Marziani et al. 2018).Through the analysis of the eigenvector 1 sequence, R Fe hasbeen demonstrated as a probe of accretion rate/Eddingtonratio (e.g., Boroson & Green 1992; Sulentic et al. 2000;Marziani et al. 2001, 2003; Shen & Ho 2014; Sun, & Shen2015), thus a primary physical driver of the shortened timelags is the accretion rate. . . . . . . .
65 + 0 .
45 log ‘ − . R Fe . . . . . . l og ( R H β / l t − d a y s ) σ = 0 . Figure 5.
New scaling relation. The scatter of the new scaling re-lation is σ = 0 . , and marked in the lower right corner. Thetwo grey points are MCG +06-26-012 and MCG +06-30-015 (seemore details in Section 3.3). We correct the intrinsic reddeningMCG +06-30-015 and use its corrected luminosity in the analysis.We do not include MCG +06-26-012 in the analysis. As a simple test, we provide here the linear regressionof the correlation between ∆ R H β and R Fe . We adopt theBCES method (Akritas, & Bershady 1996, the orthogonalleast squares) to perform the linear regression, which takesinto account both of the error bars in x and y axis. MCG +06-26-012 has a relatively low sampling cadence in the first 80days in its light curve of Wang et al. (2014a) and Hu et al.(2015), which makes its time lag may bias towards longervalue. We do not use it in the regression. And the intrinsicreddening of MCG +06-30-015 is strong in light of its highBalmer decrement (Hu et al. 2016). We correct its intrinsicreddening and use the corrected luminosity and the corre-sponding R H β, R − L . The MCG +06-26-012 and MCG +06-30-015 are marked as grey points in Figure 4 (also in thefollowing Figure 5). The linear regression is yielded as ∆ R H β = − (0 . ± . R Fe + (0 . ± . . (4)The regression and the corresponding confidence band ( σ )are shown in Figure 4. We have also tested that the residual ∆ R H β − ∆ R H β ( R Fe ) does not show any correlations withall of the spectral properties (with Spearman’s coefficients | ρ | < . ), where ∆ R H β ( R Fe ) is the ∆ R H β value deducedfrom R Fe by Equation (4).3.3. A New Scaling Relation
Because the strongest correlation is the relation between ∆ R H β and R Fe , we can add R Fe as a new parameter into the R H β – L relationship to establish a new scaling relation with smaller scatter. We fit the RM sample with the followingnew scaling relation: log ( R H β / lt − days) = α + β log (cid:96) + γ R Fe . (5)In order to obtain the uncertainties of the parameters, we em-ploy the bootstrap technique. A subset is generated by re-sampling N points from the RM sample with replacement( N is the number of the objects in the RM sample). Then,we calculate the best parameters for this subset using theLevenberg-Marquardt method (Press et al. 1992), and repeatthis procedures for 5000 times to generate the distributions of α , β , and γ . The final best parameters and the correspondinguncertainties are obtained from the α , β , and γ distributions.The fit is shown in Figure 5, and the best parameters are: α = 1 . ± . , β = 0 . ± . , γ = − . ± . . (6)The scatter of the new scaling relation is σ = 0 . , whichis much smaller than the original scatter of the R H β – L relationship ( σ ∼ . , see Du et al. 2018a). DISCUSSIONS4.1.
Some Discussions on Pairwise Correlations
In Section 3.1, we showed the pairwise correlations be-tween the parameters we compiled. Some of the correlationshave been presented in the literatures using different sam-ples of AGNs, and some have been discussed directly or in-directly. A σ H β - D H β correlation, which is a width-profilecorrelation of the H β line similar to the FWHM H β - D H β inthis paper, was presented in Collin et al. (2006) using the RMsample at that time. It was also discussed by Kollatschny& Zetzl (2011, 2013) and explained as the different contri-butions from the rotation/Keplerian motions and the turbu-lent velocities in the objects with different line widths (Kol-latschny & Zetzl 2011, 2013). The correlation in Figure 2 isgenerally the same as in Kollatschny & Zetzl (2011), but hasmore objects at the narrow-width (small FWHM H β ) end be-cause the current sample has more NLS1s or high-accretion-rate objects. However, it should be noted that the parameter D H β involves FWHM H β as the numerator, thus may intro-duce a certain degree of self correlation to the FWHM H β - D H β relation.The comparison between FWHM
Fe II and
FWHM H β hasbeen presented for the quasar sample in the Sloan DigitalSky Survey (SDSS) in, e.g., Hu et al. (2008a,b); Craccoet al. (2016). The FWHM
Fe II is systematically smallerthan
FWHM H β , which was demonstrated and explained bythe contribution from a intermediate-line region in Hu et al.(2008a,b). The Fe II emission and the intermediate-widthcomponent of H β line are both from this intermediate-lineregion, and H β has an extra very broad component (Hu etal. 2008a,b). In addition, Hu et al. (2015) shows a com-parison between the time lags of Fe II and H β using theSEAMBH sample, which is also a direct evidence for therelatively larger size of Fe II -emitting region and the smallerH β -emitting region. And Hu et al. (2015) also shows thatthe lag ratio between Fe II and H β correlates with R Fe . The R Fe − FWHM
Fe II / FWHM H β correlation in Figure 2 mayhas the same physical origin as the correlation between thelag ratio and R Fe (The gas in the Fe II region has a larger sizeand a smaller FWHM). Similarly, FWHM H β correlates with FWHM
Fe II / FWHM H β (see in Figure 2) because of the R Fe - FWHM H β correlation (Eigenvector 1 sequence, e.g.,Boroson & Green 1992; Sulentic et al. 2000; Marziani et al.2001, 2003, 2018; Shen & Ho 2014).The principal component analysis in Boroson & Green(1992) has shown a weak correlation between R Fe and EW H β using the PG quasar sample, however with a rela-tively small correlation coefficient of − . (see more de-tails in Boroson & Green 1992). The R Fe and EW H β of theRM sample presented in this paper show a slightly strongercorrelation with Spearman’s coefficient of ρ = − . . Thiscorrelation may be related to the Baldwin effect of H β line(e.g., Baldwin 1977; Korista, & Goad 2004), and especiallywith the intrinsic Baldwin effect (e.g., Gilbert, & Peterson2003; Raki´c et al. 2017). The intrinsic Baldwin effect showsan anti-correlation between EW of the emission line and theluminosity, and is also equivalent to an anti-correlation be-tween EW and the accretion rate because the BH mass keepsa constant during the observation campaign (e.g., Gilbert, &Peterson 2003; Raki´c et al. 2017). The R Fe parameter is cor-related with Eddington ratio/accretion rate, thus is naturallycorrelated with the EW H β .The asymmetry- R Fe correlation has been shown in Boro-son & Green (1992), and discussed in the context of eigen-vector 1 sequence (Sulentic et al. 2002). The high- R Fe objects tend to have stronger blue H β wings, and vise versa.The FWHM
Fe II / FWHM H β -asymmetry and FWHM H β -asymmetry correlations can also attribute to the asymmetry- R Fe correlation. The origin of the H β asymmetry must besubject to the geometry and kinematics of the BLRs, but isstill under some debate because of the degeneracy of H β pro-files with different BLR geometry and kinematics. A recentdedicated RM campaign project for the BLR kinematics ofthe AGNs with H β asymmetry has started (Du et al. 2018b),and may provide more observations for the velocity-resolvedRM measurements in the future.The EW [OIII] - EW HeII correlation is a natural result of thephotoionization physics. Both of [O
III ] and He II have highionization energy (54.9 eV for [O III ] and 54.4 eV for He II ,respectively), thus are sensitive to the variation of the AGNcircumstances (e.g., spectral energy distribution, SED) in asimilar way.4.2. The New Scaling Relation and BH Mass Measurement
Through the analysis in this paper, we found that the R Fe parameter is the dominant observational property in the scat-ter of the R H β – L relationship. This confirms the state- They may have some special BLR kinematics or inhomogeneous gasdistribution, or even binary BHs in their centers (see more details in Du etal. 2018b). M • , RM /M (cid:12) M • , S E / M (cid:12) new scaling relation ( σ = 0 . R H β − L relationship ( σ = 1 . M • , SE /M • , RM N u m b e r Figure 6.
Single-epoch BH mass versus RM BH mass. The bluepoints are estimated by the new scaling relation, while the orangepoints are obtained by the traditional R H β – L relationship. Theembedded panel shows the distributions of M • , SE /M • , RM of thenew scaling relation and the R H β – L relationship. The BHmasses estimated by the R H β – L relationship are biased towardsto higher values with respect to those from the new scaling relation.The standard deviations (simply denoted by σ ) of the distributionsare provided in the upper-left corner. We do not plot the error barsin order to show the differences more clearly. ment that the AGNs with high accretion rates tend to haveshortened lags in Du et al. (2015, 2016a, 2018a). The short-ened time lags in high- R Fe /high-accretion-rate AGNs implysmaller BLR scale sizes and smaller BH mass estimates withrespect to the traditional R H β – L relationship.The R H β – L relationship was heavily utilized as asingle-epoch BH mass estimator in large AGN samples, andhelped establish our paradigm for AGN evolution (McLure& Dunlop 2004; Vestergaard & Peterson 2006; Kollmeier etal. 2006; Greene & Ho 2007; Shen et al. 2011). The short-ened time lags in high- R Fe /high-accretion-rate AGNs makeit vital to take this into account in the BH mass estimation.Therefore, we suggest calculate the BH masses from single-epoch spectra following the steps below: (1) obtain the R H β from the luminosities and the strength of Fe II ( R Fe ) usingEquation (5), (2) get the BH mass by the following Equation(7) from the line width and the estimated R H β . Then, thedimensionless accretion rate can be easily estimated by thefollowing Equation (8) in Section 4.3.In combination with the velocity width ( ∆ V ) of the emis-sion line, RM measurement yields an estimate of BH mass as M • = f BLR ∆ V R BLR G (7) − − − log ˙ M RM − − − l og ˙ M S E new scaling relation ( σ = 0 . R H β − L relationship ( σ = 0 . . . log ( ˙ M SE / ˙ M RM ) N u m b e r − − − log ( L Bol /L Edd ) RM − − − l og ( L B o l / L E dd ) S E new scaling relation ( σ = 0 . R H β − L relationship ( σ = 0 . − log [( L Bol /L Edd ) SE / ( L Bol /L Edd ) RM ] N u m b e r Figure 7.
Accretion rates and Eddington ratios from the single-epoch spectra and the RM observations. The blue points are deduced by thenew scaling relation, while the orange points are obtained by the R H β – L relationship. The embedded panel shows the correspondingdistributions of log ( ˙ M SE / ˙ M RM ) and log [( L Bol /L Edd ) SE / ( L Bol /L Edd ) RM ] . The accretion rates and Eddington ratios estimated from the R H β – L relationship are biased downward in the high- ˙ M end. The standard deviations (simply denoted by σ ) of the distributions areprovided in the upper-left corner. We do not plot the error bars in order to show the differences more clearly. (e.g., Wandel et al. 1999; Peterson & Wandel 1999; Petersonet al. 2004), where G is the gravitational constant, and f BLR is the virial factor related to the geometry and kinematics ofthe BLR (e.g., Onken et al. 2004; Park et al. 2012; Grier etal. 2013a; Ho & Kim 2014; Woo et al. 2015). Although mea-suring BH mass through RM technique is feasible for a smallnumber of objects, it is not easy to apply RM to large AGNsamples because RM is fairly time-consuming (always con-tinuing for months to years). Some multi-object RM cam-paigns based on fiber-fed telescopes, e.g., the RM campaignsof SDSS (Shen et al. 2015) and OzDES (King et al. 2015), arecommitted to enlarge the sample of RM objects but still on-going. Fortunately, the R H β – L relationship can be usedto obtain R BLR from the single-epoch spectra very simply.The geometry and kinematics of the BLRs determine thevirial factor f BLR in BH mass estimate (in Equation 7). Com-paring the RM AGNs with stellar velocity dispersion mea-surements with M • − σ ∗ relation of inactive galaxies gives f BLR ∼ if the velocity width of H β is measured from FWHM H β (e.g., Onken et al. 2004; Ho & Kim 2014; Wooet al. 2015). The virial factor for high- R Fe /high-accretion-rate AGNs is still a matter of some debate. Through fittingAGN spectral energy distribution by accretion disk model,Mej´ıa-Restrepo et al. (2018) found a correlation between thevirial factor and the width of emission line ( f BLR ∼ if FWHM H β ∼ − and f BLR ∼ . if FWHM H β ∼ − ). Yu et al. (2019) also shows potential a corre- lation between f BLR and line width. From the BLR modelingresults of the small samples in Pancoast et al. (2014); Grier etal. (2017b); Williams et al. (2018), the virial factor is roughlyconsistent with the value derived from the M • − σ ∗ relation,and the virial factors of individual objects do no show signif-icant correlation with the Eddington ratios (or show potentialand weak anti-correlation, namely smaller virial factor forhigher Eddington ratio). NLS1s (thought to have smaller BHmasses and higher accretion rates) tend to host pseudobulges(e.g., Mathur et al. 2012). Ho & Kim (2014) classified theAGN sample to classical bulges/ellipticals and pseudobulges,and derived a virial factor of the AGNs with pseudobulgessmaller than 1. Woo et al. (2015) found that NLS1s haveno significant differences from the other AGNs, and derived f BLR = 1 . if using FWHM H β as the line-width measure-ment. Thus, adopting f BLR ∼ and acknowledging its largeuncertainty is acceptable at present.As a simple test, we compare the BH masses measuredby RM ( M • , RM ) with the masses estimated from the single-epoch spectra ( M • , SE ) using both of the new scaling relationin Section 3.3 and the traditional R H β – L relationship(see Section 3.2) in Figure 6. Here, we adopt f BLR = 1 forsimplicity and list the RM BH masses in Table 1. The scat-ter is quantified by the standard deviation of M • , SE /M • , RM .We do not plot the error bars in order to show clearly thedifferences between the estimates from the new scaling rela-tion and the R H β – L relationship. The scatter of M • , SE − − M SE − − − − l og ˙ M R M new scaling relation ( σ = 0 . . (2016) ( σ = 0 . . . log ( ˙ M SE / ˙ M RM ) N u m b e r − − − L Bol /L Edd ) SE − − − l og ( L B o l / L E dd ) R M new scaling relation ( σ = 0 . . (2016) ( σ = 0 . . . log [( L Bol /L Edd ) SE / ( L Bol /L Edd ) RM ] N u m b e r Figure 8.
Comparisons with the accretion rates and Eddington ratios estimated from the fundamental plane of BLR. The blue points are deducedby the new scaling relation, while the orange points are obtained by the FP in Du et al. (2016b). The embedded panel shows the correspondingdistributions of log ( ˙ M SE / ˙ M RM ) and log [( L Bol /L Edd ) SE / ( L Bol /L Edd ) RM ] . The accretion rates and Eddington ratios estimated from theFP show much larger scatter. The standard deviations (simply denoted by σ ) of the distributions are provided in the upper-left corner. We donot plot the error bars of the blue points for clarity. from the new scaling relation is 0.52, while the scatter ob-tained by the traditional R H β – L relationship is 1.18.The embedded plot in Figure 6 shows the distributions of M • , SE /M • , RM . The M • , SE /M • , RM distribution is biased to-ward the value larger than 1 for the traditional R H β – L relationship, which means the BH masses are overestimatedif using this simple relationship.Recently, Mart´ınez-Aldama et al. (2019) found a correc-tion for the time delay based on the dimensionless accretionrate ( ˙ M in the following Section 4.3) considering the anti-correlation between f BLR and line width, established a corre-lation between the corrected time lag ( R corrBLR ) and L , anddiscussed the measurements of the cosmological distancesusing this correlation. Their correction ( R corrBLR ) relies on themeasured L , FWHM H β , and R BLR itself. The new scal-ing relation in the present paper is established in a differentperspective (based on spectral properties), and can deduce R BLR directly from L and R Fe .4.3. Accretion Rate and Eddington Ratio
Because of the shortened lags, the accretion rates or Ed-dington ratios would be underestimated by factors of a few ifusing the traditional R H β – L relationship. From the stan-dard disk model (Shakura, & Sunyaev 1973), the accretionrate can be estimated by the formula of ˙ M = 20 . (cid:18) (cid:96) cos i (cid:19) / m − , (8) where m = M • / M (cid:12) , and i is inclination angle ofthe accretion disk (here we adopt cos i = 0 . as an av-erage value for all of the AGNs, see the discussions in Duet al. 2014; Wang et al. 2014a; Du et al. 2016a). Here,we compare the ˙ M and Eddington ratio ( L Bol /L Edd ) es-timates obtained from the single-epoch spectra (using thenew scaling relation and the R H β – L relationship) withthose values calculated from RM (listed in Table 1) inFigure 7, where L Bol is the bolometric luminosity and L Edd = 1 . × ( M • /M (cid:12) ) is the Eddington luminos-ity for the gas with solar composition. Here we simplyadopt L Bol = 10 L (Kaspi et al. 2000), but it should benoted that the bolometric correction factor depends on ac-cretion rate or BH mass (Jin et al. 2012). We do not drawthe error bars in order to show the differences at high- ˙ M end more clearly. It is obvious that the points of the newscaling relation at high-accretion-rate end are much closerto the diagonal, while those estimated by the traditional R H β – L relationship are biased downward. The distribu-tions and the standard deviations of log ( ˙ M SE / ˙ M RM ) and log [( L Bol /L Edd ) SE / ( L Bol /L Edd ) RM ] are also provided inFigure 7. The new scaling relation should be preferentiallyused in the estimation of accretion rates or Eddington ratiosfrom the single-epoch spectra in the statistical study of largeAGN samples.Du et al. (2016b) established a bivariate correlation be-tween ˙ M and ( R Fe , D H β ), which can be used to estimate1the accretion rate directly from the BLR properties. It iscalled the fundamental plane (FP) of BLR. The FP can de-duce ˙ M or Eddington ratio estimates without any luminositymeasurements (see also Negrete et al. 2018), however hasfairly large uncertainties (the scatter of the FP is 0.7 for ˙ M and 0.48 for Eddington ratio, respectively, see more detailsin Du et al. 2016b). As a comparison, we plot the ˙ M esti-mates from the FP and the new scaling relation in Figure 8.The single-epoch ˙ M and Eddington ratios corresponding tothe FP method are estimated from the R Fe and D H β listed inTable 2 and the FP in Du et al. (2016b). It should be notedthat the we switch x and y axes of Figure 8 (with respect toFigure 7) for easier comparison with the figures in Du et al.(2016b). Again, we do not draw the error bars of the ˙ M from the new scaling relation for clarity. The scatters of the ˙ M and L Bol /L Edd estimated from the FP are much larger.The FP connect the BLR physics with the accretion status ofAGNs, however the different temperature, number density,metallicity of the BLR in different AGN introduces large un-certainties. The FP (Du et al. 2016b) is a good beginning thatsearches for direct indicator of accretion rate from the BLRobservational properties, but its scatter and accuracy shouldbe improved by including more observational properties inthe future.The strong correlation between accretion rates ˙ M and R Fe has been explored by Panda et al. (2018, 2019), but still re-mains open. Accretion flows to the central BHs supplied byeither star formation from torus (Wang et al. 2010), or asym-metric dynamics (Begelman, & Shlosman 2009) will havedifferent dependence on metallicity, however, R Fe is not aunique function of metallicity (Baldwin et al. 2004; Verneret al. 2004). More details of photoionization are necessaryto investigate R Fe dependence on BLR clouds (density, tem-perature and metallicity) and the SEDs of accretion disks.4.4. Accretion Rate or Orientation?
The eigenvector 1 sequence can do some help to break thedegeneracy of accretion rate and orientation. It was demon-strated that the accretion rate or Eddington ratio drives thevariation of R Fe , while the orientation effect dominantly con-trols the dispersion in FWHM H β at fixed R Fe (e.g., Marzianiet al. 2001; Shen & Ho 2014). We plot the eigenvector 1sequence of the RM sample color-coded by ∆ R H β in Fig-ure 9. It is obvious that the objects with the most significantlag deviations are located in the lower right corner (with thestrongest R Fe ). Furthermore, there is no significant trend inthe FWHM H β -axis at fixed R Fe . It means that the accretionrate definitely plays the primary role in the shortening of thetime lags but the orientation does not contribute much. InFigure 4, the dispersion of the ∆ R H β at lower FWHM H β is larger, this is caused by the higher R Fe in those objects,which is clearly shown in Figure 9. SUMMARYIn this paper, we systematically investigate the dependenceof the R H β – L relationship on optical spectra for a wide . . . . . . R Fe F W H M H β ( k m s − ) O r i e n t a t i o n Accretion Rate − . − . − . . . ∆ R H β = log ( R H β /R H β, R − L ) Figure 9.
Main Sequence of AGNs. The points are color-coded by ∆ R H β . Accretion rate/Eddington ratio drives the variation of R Fe ,while the orientation effect dominantly controls the dispersion in FWHM H β at fixed R Fe (Marziani et al. 2001; Shen & Ho 2014).It is obvious that accretion rate definitely plays the primary role inthe shortening of the time lags (see more details in Section 4.4). range of AGN parameters. The reverberation mapping cam-paign of the AGNs with high accretion rates show many ob-jects deviate significantly from the traditional R H β – L re-lationship (Du et al. 2015, 2016a, 2018a). We collect 8 dif-ferent single-epoch spectral properties to investigate how thedeviation of an AGN from the R H β – L relationship cor-relates with those properties and to understand what is theorigin of the shortened lags. • The flux ratio between Fe II and H β lines ( R Fe ) isconfirmed to be the most prominent property that cor-relates with the deviation of an AGN from the R H β – L relationship. R Fe is thought to be the indicatorof accretion rate, therefore, accretion rate is the driverfor the shortened lags. FWHM H β , which is inducedby the orientation of an AGN to line of sight, does notshow clear trend with the lag shortening. Thus, theorientation is not a dominant factor. • We established a new scaling relation with the form of log R H β = α + β log (cid:96) + γ R Fe , which can be usedas a single-epoch estimator of BH mass and accretionrate, where α = 1 . ± . , β = 0 . ± . , and γ = − . ± . . 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Objects τ H β log L FWHM log ( M • /M (cid:12) ) log ˙ M log L H β EW(H β ) log L [O III] EW([O
III ]) Ref.(days) (erg s − ) (km s − ) (erg s − ) ( ˚A) (erg s − ) ( ˚A)Mrk 335 . +1 . − . . ± .
06 2096 ±
170 6 . +0 . − . . +0 . − . . ± .
06 110 . ± . . ± .
06 31 . ± .
1, 2, 3 . +4 . − . . ± .
06 1792 ± . +0 . − . . +0 . − . . ± .
06 119 . ± . . ± .
06 36 . ± .
4, 5, 6, 7 . +6 . − . . ± .
06 1679 ± . +0 . − . . +0 . − . . ± .
06 111 . ± . . ± .
06 31 . ± .
4, 5, 6, 7 . +0 . − . . ± .
06 1724 ±
236 6 . +0 . − . . +0 . − . . ± .
07 89 . ± . . ± .
06 38 . ± .
4, 8, 9 . + . − . . ± .
07 1707 ±
79 6 . + . − . . + . − . . ± .
09 108 . ± . . ± .
10 34 . ± . ...PG 0026+129 . +24 . − . . ± .
02 2544 ±
56 8 . +0 . − . . +0 . − . . ± .
04 46 . ± . . ± .
02 21 . ± .
4, 5, 6, 10PG 0052+251 . +24 . − . . ± .
03 5008 ±
73 8 . +0 . − . − . +0 . − . . ± .
05 107 . ± . . ± .
02 48 . ± .
4, 5, 6, 10Fairall9 . +3 . − . . ± .
04 5999 ±
66 8 . +0 . − . − . +0 . − . . ± .
04 249 . ± . . ± .
03 75 . ± .
4, 5, 6, 11Mrk 590 . +3 . − . . ± .
06 2788 ±
29 7 . +0 . − . − . +0 . − . . ± .
06 107 . ± . . ± .
06 25 . ± .
4, 5, 6, 7 . +8 . − . . ± .
09 3729 ±
426 7 . +0 . − . − . +0 . − . . ± .
16 142 . ± . . ± .
06 74 . ± .
4, 5, 6, 7 . +4 . − . . ± .
07 2744 ±
79 7 . +0 . − . − . +0 . − . . ± .
07 119 . ± . . ± .
06 42 . ± .
4, 5, 6, 7 . +3 . − . . ± .
06 2500 ±
43 7 . +0 . − . − . +0 . − . . ± .
07 94 . ± . . ± .
06 22 . ± .
4, 5, 6, 7 . + . − . . ± .
21 2716 ±
202 7 . + . − . − . + . − . . ± .
12 108 . ± . . ± .
06 29 . ± . ...Mrk 1044 . +3 . − . . ± .
10 1178 ±
22 6 . +0 . − . . +0 . − . . ± .
09 101 . ± . . ± .
09 11 . ± .
1, 2, 33C 120 . +21 . − . . ± .
05 2327 ±
48 7 . +0 . − . . +0 . − . . ± .
06 100 . ± . . ± .
05 81 . ± .
4, 5, 6, 7 . +2 . − . . ± .
05 3529 ±
176 7 . +0 . − . − . +0 . − . . ± .
05 135 . ± . . ± .
05 134 . ± .
4, 8, 9 . + . − . . ± .
10 2472 ±
729 7 . + . − . − . + . − . . ± .
04 118 . ± . . ± .
07 116 . ± . ...IRAS 04416+1215 . +13 . − . . ± .
03 1522 ±
44 6 . +0 . − . . +0 . − . . ± .
02 55 . ± . . ± .
05 23 . ± .
1, 2, 3Ark 120 . +8 . − . . ± .
06 6042 ±
35 8 . +0 . − . − . +0 . − . . ± .
05 211 . ± . . ± .
05 18 . ± .
4, 5, 6, 7 . +4 . − . . ± .
08 6246 ±
78 8 . +0 . − . − . +0 . − . . ± .
07 321 . ± . . ± .
05 41 . ± .
4, 5, 6, 7 . + . − . . ± .
25 6077 ±
147 8 . + . − . − . + . − . . ± .
13 244 . ± . . ± .
05 22 . ± . ...MCG +08-11-011 . +0 . − . . ± .
11 4139 ±
207 7 . +0 . − . − . +0 . − . . ± .
09 108 . ± . . ± .
07 276 . ± . . +5 . − . . ± .
04 4980 ±
249 7 . +0 . − . − . +0 . − . . ± .
04 58 . ± . . ± .
04 33 . ± . . +8 . − . . ± .
07 5056 ±
85 7 . +0 . − . − . +0 . − . . ± .
07 92 . ± . . ± .
07 56 . ± .
4, 5, 6, 7 . +6 . − . . ± .
07 4760 ±
31 7 . +0 . − . − . +0 . − . . ± .
08 78 . ± . . ± .
07 44 . ± .
4, 5, 6, 7 . +6 . − . . ± .
07 4766 ±
71 7 . +0 . − . − . +0 . − . . ± .
07 86 . ± . . ± .
07 51 . ± .
4, 5, 6, 7 . + . − . . ± .
07 4793 ±
145 7 . + . − . − . + . − . . ± .
05 85 . ± . . ± .
07 50 . ± . ...SDSS J074352 . +5 . − . . ± .
02 3156 ±
36 7 . +0 . − . . +0 . − . . ± .
01 65 . ± . . ± .
05 5 . ± . . +18 . − . . ± .
01 1904 ± . +0 . − . . +0 . − . . ± .
03 51 . ± . . ± .
02 5 . ± . . +15 . − . . ± .
05 1495 ±
67 7 . +0 . − . . +0 . − . . ± .
03 68 . ± . . ± .
03 25 . ± . . +5 . − . . ± .
04 1679 ±
35 7 . +0 . − . . +0 . − . . ± .
04 70 . ± . . ± .
03 16 . ± . . + . − . . ± .
09 1645 ±
139 7 . + . − . . + . − . . ± .
10 69 . ± . . ± .
05 19 . ± . ...Mrk 382 . +2 . − . . ± .
08 1462 ±
296 6 . +0 . − . . +0 . − . . ± .
05 39 . ± . . ± .
05 32 . ± .
1, 2, 3SDSS J075949 . +17 . − . . ± .
03 1807 ±
11 7 . +0 . − . . +0 . − . . ± .
02 97 . ± . . ± .
04 8 . ± . . +11 . − . . ± .
06 1783 ±
17 7 . +0 . − . . +0 . − . . ± .
04 98 . ± . . ± .
07 6 . ± . . + . − . . ± .
03 1800 ±
19 7 . + . − . . + . − . . ± .
02 97 . ± . . ± .
10 8 . ± . ...SDSS J080101 . +9 . − . . ± .
03 1930 ±
18 6 . +0 . − . . +0 . − . . ± .
02 105 . ± . . ± .
05 3 . ± . . +8 . − . . ± .
04 1188 ± . +0 . − . . +0 . − . . ± .
03 64 . ± . . ± .
07 13 . ± . . +14 . − . . ± .
04 1290 ±
13 6 . +0 . − . . +0 . − . . ± .
05 52 . ± . . ± .
04 20 . ± . . + . − . . ± .
04 1194 ±
70 6 . + . − . . + . − . . ± .
09 59 . ± . . ± .
12 17 . ± . ...PG 0804+761 . +18 . − . . ± .
02 3053 ±
38 8 . +0 . − . . +0 . − . . ± .
03 122 . ± . . ± .
03 16 . ± .
4, 5, 6, 10SDSS J081441 . +12 . − . . ± .
07 1615 ±
22 6 . +0 . − . . +0 . − . . ± .
03 132 . ± . . ± .
07 11 . ± . . +7 . − . . ± .
04 1782 ±
16 7 . +0 . − . . +0 . − . . ± .
02 140 . ± . . ± .
04 11 . ± . . + . − . . ± .
06 1730 ±
121 7 . + . − . . + . − . . ± .
03 137 . ± . . ± .
06 11 . ± . ...SDSS J081456 . +7 . − . . ± .
04 2409 ±
61 7 . +0 . − . . +1 . − . . ± .
03 74 . ± . . ± .
04 24 . ± . . +1 . − . . ± .
16 8026 ±
401 7 . +0 . − . − . +0 . − . . ± .
12 165 . ± . . ± .
10 41 . ± . . +5 . − . . ± .
02 1758 ±
16 6 . +0 . − . . +0 . − . . ± .
02 56 . ± . . ± .
03 12 . ± . . +3 . − . . ± .
04 1243 ±
13 6 . +0 . − . . +0 . − . . ± .
05 55 . ± . . ± .
04 10 . ± . Table 1 continued Table 1 (continued)
Objects τ H β log L FWHM log ( M • /M (cid:12) ) log ˙ M log L H β EW(H β ) log L [O III] EW([O
III ]) Ref.(days) (erg s − ) (km s − ) (erg s − ) ( ˚A) (erg s − ) ( ˚A) . +7 . − . . ± .
02 1297 ±
12 6 . +0 . − . . +0 . − . . ± .
03 61 . ± . . ± .
03 8 . ± . . + . − . . ± .
02 1273 ±
39 6 . + . − . . + . − . . ± .
03 60 . ± . . ± .
07 9 . ± . ...PG 0844+349 . +13 . − . . ± .
07 2694 ±
58 7 . +0 . − . . +0 . − . . ± .
05 111 . ± . . ± .
03 13 . ± .
5, 6, 10, 15SDSS J085946 . +19 . − . . ± .
03 1718 ±
16 7 . +0 . − . . +1 . − . . ± .
02 63 . ± . . ± .
04 16 . ± . . +5 . − . . ± .
04 1543 ± . +0 . − . . +0 . − . . ± .
05 139 . ± . . ± .
05 78 . ± .
4, 5, 6, 7 . +10 . − . . ± .
04 1658 ± . +0 . − . . +0 . − . . ± .
05 94 . ± . . ± .
05 66 . ± .
4, 5, 6, 7 . +14 . − . . ± .
05 1600 ±
39 7 . +0 . − . . +0 . − . . ± .
04 139 . ± . . ± .
05 110 . ± .
4, 5, 6, 7 . + . − . . ± .
12 1634 ±
83 7 . + . − . . + . − . . ± .
08 123 . ± . . ± .
05 81 . ± . ...SDSS J093302 . +3 . − . . ± .
13 1800 ±
25 7 . +0 . − . . +0 . − . . ± .
05 31 . ± . . ± .
06 4 . ± . . +2 . − . . ± .
04 1209 ±
16 6 . +0 . − . . +0 . − . . ± .
04 53 . ± . . ± .
13 8 . ± . . +21 . − . . ± .
01 3071 ±
27 8 . +0 . − . . +0 . − . . ± .
04 64 . ± . . ± .
02 18 . ± .
4, 5, 6, 10SDSS J100402 . +43 . − . . ± .
01 2088 ± . +0 . − . . +0 . − . . ± .
01 53 . ± . . ± .
04 4 . ± . . +23 . − . . ± .
02 2311 ± . +0 . − . . +0 . − . . ± .
02 52 . ± . . ± .
11 2 . ± . . +0 . − . . ± .
11 4112 ±
206 7 . +0 . − . − . +0 . − . . ± .
10 71 . ± . . ± .
10 142 . ± .
4, 9, 16SDSS J102339 . +19 . − . . ± .
03 1733 ±
29 7 . +0 . − . . +0 . − . . ± .
03 57 . ± . . ± .
04 8 . ± . . +1 . − . . ± .
06 1588 ±
58 6 . +0 . − . . +0 . − . . ± .
04 55 . ± . . ± .
04 10 . ± .
1, 2, 3 . +0 . − . . ± .
04 1462 ± . +0 . − . . +0 . − . . ± .
05 57 . ± . . ± .
04 21 . ± .
4, 17 . + . − . . ± .
04 1462 ±
86 6 . + . − . . + . − . . ± .
06 56 . ± . . ± .
27 18 . ± . ...NGC 3516 . +1 . − . . ± .
20 5384 ±
269 7 . +0 . − . − . +0 . − . . ± .
18 94 . ± . . ± .
17 59 . ± .
4, 9, 16SBS 1116+583A . +0 . − . . ± .
23 3668 ±
186 6 . +0 . − . − . +0 . − . . ± .
07 186 . ± . . ± .
06 117 . ± .
4, 17Arp 151 . +0 . − . . ± .
10 3098 ±
69 6 . +0 . − . − . +0 . − . . ± .
11 130 . ± . . ± .
07 80 . ± .
4, 17NGC 3783 . +3 . − . . ± .
18 3770 ±
68 7 . +0 . − . − . +0 . − . . ± .
18 144 . ± . . ± .
17 126 . ± .
4, 5, 6, 18MCG +06-26-012 . +8 . − . . ± .
11 1334 ±
80 6 . +0 . − . − . +0 . − . . ± .
06 114 . ± . . ± .
05 38 . ± .
1, 2, 3UGC 06728 . +0 . − . . ± .
08 1642 ±
161 5 . +0 . − . . +0 . − . . ± .
05 49 . ± . . ± .
02 33 . ± . . +0 . − . . ± .
14 2409 ±
24 6 . +0 . − . − . +0 . − . . ± .
10 94 . ± . . ± .
08 293 . ± .
4, 17NGC 4051 . +0 . − . . ± .
19 851 ±
277 5 . +0 . − . . +1 . − . . ± .
18 86 . ± . . ± .
17 79 . ± .
4, 9, 16 . +0 . − . . ± .
18 1145 ±
192 5 . +0 . − . . +1 . − . . ± .
18 71 . ± . . ± .
17 108 . ± . . + . − . . ± .
15 1076 ±
277 5 . + . − . . + . − . . ± .
19 78 . ± . . ± .
12 104 . ± . ...NGC 4151 . +1 . − . . ± .
21 6371 ±
150 7 . +0 . − . − . +0 . − . . ± .
20 150 . ± . . ± .
17 1637 . ± .
4, 5, 6, 20PG 1211+143 . +25 . − . . ± .
08 2012 ±
37 7 . +0 . − . . +0 . − . . ± .
06 100 . ± . . ± .
03 16 . ± .
5, 6, 10, 15Mrk 202 . +1 . − . . ± .
14 1471 ±
18 6 . +0 . − . . +0 . − . . ± .
09 70 . ± . . ± .
07 87 . ± .
4, 17NGC 4253 . +1 . − . . ± .
12 1609 ±
39 6 . +0 . − . . +0 . − . . ± .
12 81 . ± . . ± .
12 326 . ± .
4, 17PG 1226+023 . +8 . − . . ± .
05 3314 ±
59 8 . +0 . − . . +0 . − . . ± .
03 80 . ± . . ± .
02 10 . ± . . +27 . − . . ± .
05 3828 ±
54 8 . +0 . − . − . +0 . − . . ± .
06 209 . ± . . ± .
03 58 . ± .
4, 5, 6, 10NGC 4593 . +0 . − . . ± .
18 5143 ±
16 7 . +0 . − . − . +0 . − . . ± .
18 101 . ± . . ± .
17 30 . ± .
4, 6, 22 . +1 . − . . ± .
18 4395 ±
362 7 . +0 . − . − . +0 . − . . ± .
18 115 . ± . . ± .
18 51 . ± .
9, 23 . + . − . . ± .
37 5142 ±
572 7 . + . − . − . + . − . . ± .
33 108 . ± . . ± .
22 39 . ± . ...IRAS F12397+3333 . +5 . − . . ± .
05 1802 ±
560 6 . +0 . − . . +0 . − . . ± .
04 54 . ± . . ± .
04 66 . ± .
1, 2, 3NGC 4748 . +1 . − . . ± .
12 1947 ±
66 6 . +0 . − . . +0 . − . . ± .
10 136 . ± . . ± .
10 300 . ± .
4, 17PG 1307+085 . +36 . − . . ± .
02 5059 ±
133 8 . +0 . − . − . +0 . − . . ± .
06 98 . ± . . ± .
02 39 . ± .
4, 5, 6, 10MCG +06-30-015 . +1 . − . . ± .
23 1958 ±
75 6 . +0 . − . − . +0 . − . . ± .
13 60 . ± . . ± .
12 104 . ± . . +3 . − . . ± .
12 1933 ±
81 6 . +0 . − . − . +0 . − . . ± .
12 108 . ± . . ± .
12 110 . ± . . + . − . . ± .
11 1947 ±
58 6 . + . − . − . + . − . . ± .
19 91 . ± . . ± .
09 110 . ± . ...NGC 5273 . +1 . − . . ± .
16 5688 ±
163 7 . +0 . − . − . +1 . − . . ± .
11 82 . ± . . ± .
08 46 . ± . . +3 . − . . ± .
07 5354 ±
32 7 . +0 . − . − . +0 . − . . ± .
06 132 . ± . . ± .
06 36 . ± .
4, 5, 6, 27PG 1411+442 . +61 . − . . ± .
02 2801 ±
43 8 . +0 . − . − . +0 . − . . ± .
03 99 . ± . . ± .
03 21 . ± .
4, 5, 6, 10NGC 5548 . +1 . − . . ± .
10 4674 ±
63 7 . +0 . − . − . +0 . − . . ± .
10 128 . ± . . ± .
09 89 . ± .
4, 5, 6, 28 . +2 . − . . ± .
11 5418 ±
107 8 . +0 . − . − . +0 . − . . ± .
13 151 . ± . . ± .
09 156 . ± .
4, 5, 6, 28 . +2 . − . . ± .
09 5236 ±
87 7 . +0 . − . − . +0 . − . . ± .
10 119 . ± . . ± .
09 97 . ± .
4, 5, 6, 28 . +1 . − . . ± .
11 5986 ±
95 7 . +0 . − . − . +0 . − . . ± .
17 144 . ± . . ± .
09 185 . ± .
4, 5, 6, 28 . +1 . − . . ± .
10 5931 ±
42 7 . +0 . − . − . +0 . − . . ± .
09 135 . ± . . ± .
09 103 . ± .
4, 5, 6, 28
Table 1 continued Table 1 (continued)
Objects τ H β log L FWHM log ( M • /M (cid:12) ) log ˙ M log L H β EW(H β ) log L [O III] EW([O
III ]) Ref.(days) (erg s − ) (km s − ) (erg s − ) ( ˚A) (erg s − ) ( ˚A) . +3 . − . . ± .
09 7378 ±
39 8 . +0 . − . − . +0 . − . . ± .
10 114 . ± . . ± .
09 91 . ± .
4, 5, 6, 28 . +2 . − . . ± .
09 6946 ±
79 8 . +0 . − . − . +0 . − . . ± .
09 102 . ± . . ± .
09 65 . ± .
4, 5, 6, 28 . +1 . − . . ± .
09 6623 ±
93 8 . +0 . − . − . +0 . − . . ± .
10 106 . ± . . ± .
09 80 . ± .
4, 5, 6, 28 . +2 . − . . ± .
10 6298 ±
65 8 . +0 . − . − . +0 . − . . ± .
10 153 . ± . . ± .
09 125 . ± .
4, 5, 6, 28 . +4 . − . . ± .
09 6177 ±
36 8 . +0 . − . − . +0 . − . . ± .
10 98 . ± . . ± .
09 56 . ± .
4, 5, 6, 28 . +3 . − . . ± .
09 6247 ±
57 8 . +0 . − . − . +0 . − . . ± .
09 106 . ± . . ± .
09 68 . ± .
4, 5, 6, 28 . +5 . − . . ± .
11 6240 ±
77 7 . +0 . − . − . +0 . − . . ± .
13 172 . ± . . ± .
09 170 . ± .
4, 5, 6, 28 . +5 . − . . ± .
11 6478 ±
108 8 . +0 . − . − . +0 . − . . ± .
14 139 . ± . . ± .
09 167 . ± .
4, 5, 6, 28 . +2 . − . . ± .
13 6396 ±
167 7 . +0 . − . − . +0 . − . . ± .
10 74 . ± . . ± .
09 240 . ± .
4, 29 . +0 . − . . ± .
11 12771 ±
71 8 . +0 . − . − . +0 . − . . ± .
10 105 . ± . . ± .
09 211 . ± .
4, 17 . +2 . − . . ± .
11 11481 ±
574 8 . +0 . − . − . +0 . − . . ± .
10 96 . ± . . ± .
09 220 . ± .
4, 16 . +1 . − . . ± .
09 9912 ±
362 8 . +0 . − . − . +0 . − . . ± .
09 160 . ± . . ± .
09 107 . ± . . +0 . − . . ± .
09 9496 ±
418 7 . +0 . − . − . +0 . − . . ± .
09 91 . ± . . ± .
09 64 . ± . . + . − . . ± .
19 7256 ± . + . − . − . + . − . . ± .
21 117 . ± . . ± .
07 91 . ± . ...PG 1426+015 . +29 . − . . ± .
02 7113 ±
160 8 . +0 . − . − . +0 . − . . ± .
04 80 . ± . . ± .
03 19 . ± .
4, 5, 6, 10Mrk 817 . +3 . − . . ± .
05 4711 ±
49 7 . +0 . − . − . +0 . − . . ± .
05 98 . ± . . ± .
05 28 . ± .
4, 5, 6, 7 . +3 . − . . ± .
05 5237 ±
67 7 . +0 . − . − . +0 . − . . ± .
06 108 . ± . . ± .
05 37 . ± .
4, 5, 6, 7 . +6 . − . . ± .
05 4767 ±
72 8 . +0 . − . − . +0 . − . . ± .
05 91 . ± . . ± .
05 36 . ± .
4, 5, 6, 7 . +3 . − . . ± .
05 5627 ±
30 7 . +0 . − . − . +0 . − . . ± .
05 43 . ± . . ± .
05 24 . ± .
4, 9, 16 . + . − . . ± .
09 5348 ±
536 7 . + . − . − . + . − . . ± .
14 78 . ± . . ± .
04 31 . ± . ...Mrk 1511 . +0 . − . . ± .
06 4171 ±
137 7 . +0 . − . − . +0 . − . . ± .
06 115 . ± . . ± .
05 36 . ± .
9, 23Mrk 290 . +1 . − . . ± .
06 4543 ±
227 7 . +0 . − . − . +0 . − . . ± .
06 153 . ± . . ± .
06 150 . ± .
4, 9, 16Mrk 486 . +7 . − . . ± .
05 1942 ±
67 7 . +0 . − . . +0 . − . . ± .
04 135 . ± . . ± .
04 22 . ± .
1, 2, 3Mrk 493 . +1 . − . . ± .
08 778 ±
12 6 . +0 . − . . +0 . − . . ± .
05 87 . ± . . ± .
05 12 . ± .
1, 2, 3PG 1613+658 . +15 . − . . ± .
02 9074 ±
103 8 . +0 . − . − . +0 . − . . ± .
03 86 . ± . . ± .
02 33 . ± .
4, 5, 6, 10PG 1617+175 . +29 . − . . ± .
02 6641 ±
190 8 . +0 . − . − . +0 . − . . ± .
05 114 . ± . . ± .
03 14 . ± .
4, 5, 6, 10PG 1700+518 . +45 . − . . ± .
01 2252 ±
85 8 . +0 . − . . +0 . − . . ± .
02 78 . ± . . ± .
02 3 . ± .
4, 5, 6, 103C 382 . +8 . − . . ± .
10 7652 ±
383 8 . +0 . − . − . +0 . − . . ± .
03 259 . ± . . ± .
03 61 . ± . . +6 . − . . ± .
10 12694 ±
13 8 . +0 . − . − . +0 . − . . ± .
05 206 . ± . . ± .
03 169 . ± .
4, 5, 6, 30 . +3 . − . . ± .
03 13211 ±
28 9 . +0 . − . − . +0 . − . . ± .
04 97 . ± . . ± .
03 30 . ± .
4, 31 . + . − . . ± .
58 12796 ±
361 9 . + . − . − . + . − . . ± .
35 108 . ± . . ± .
05 31 . ± . ...KA 1858+4850 . +2 . − . . ± .
05 1820 ±
79 6 . +0 . − . . +0 . − . . ± .
04 146 . ± . . ± .
03 47 . ± . . +0 . − . . ± .
28 3323 ± . +0 . − . − . +0 . − . . ± .
28 121 . ± . . ± .
28 69 . ± .
4, 17Mrk 509 . +6 . − . . ± .
05 3015 ± . +0 . − . − . +0 . − . . ± .
04 132 . ± . . ± .
05 77 . ± .
4, 5, 6, 7PG 2130+099 . +2 . − . . ± .
04 2101 ±
100 7 . +0 . − . . +0 . − . . ± .
04 142 . ± . . ± .
04 38 . ± .
4, 8, 9NGC 7469 . +3 . − . . ± .
11 4369 ± . +0 . − . − . +0 . − . . ± .
10 63 . ± . . ± .
09 77 . ± .
4, 33N
OTE —The objects are sorted in the order of right ascension. References: (1) Du et al. (2014), (2) Wang et al. (2014a), (3) Hu et al. (2015), (4) Bentz et al. (2013), (5) Collin et al.(2006), (6) Kaspi et al. (2005), (7) Peterson et al. (1998), (8) Grier et al. (2012), (9) Du et al. (2015), (10) Kaspi et al. (2000), (11) Santos-Lle´o et al. (1997), (12) Fausnaugh et al.(2017), (13) Du et al. (2018a), (14) Du et al. (2016a), (15) Bentz et al. (2009b), (16) Denney et al. (2010), (17) Bentz et al. (2009a), (18) Stirpe et al. (1994), (19) Bentz et al. (2016a),(20) Bentz et al. (2006), (21) Zhang et al. (2019), (22) Denney et al. (2006), (23) Barth et al. (2013), (24) Bentz et al. (2016b), (25) Hu et al. (2016), (26) Bentz et al. (2014), (27)Santos-Lle´o et al. (2001), (28) Peterson et al. (2002), (29) Bentz et al. (2007), (30) Lu et al. (2016), (31) Pei et al. (2017), (32) Dietrich et al. (1998), (33) Dietrich et al. (2012), (34)Pei et al. (2014), (35) Peterson et al. (2014). For the objects with multiple observations, the ˙ M of the individual campaigns are calculated based on the averaged M • . Table 2 . Single-epoch spectral properties
Objects R Fe D H β FWHM Fe / FWHM H β Asymmetry EW(He II ) Ref.( ˚A)Mrk 335 .
62 1 . ± .
05 0 . ± . − . ± .
007 23 . ± .
1, 2, 3 a PG 0026+129 .
33 1 . ± .
09 0 . ± .
02 0 . ± .
006 13 . ± .
1, 4, 5 a PG 0052+251 .
12 2 . ± .
05 0 . ± . − . ± .
004 28 . ± .
1, 4, 6 a Fairall9 .
49 2 . ± .
03 0 . ± . − . ± .
022 12 . ± .
1, 8, 9 a Mrk 590 .
45 1 . ± .
07 0 . ± . − . ± .
010 47 . ± .
1, 10 a Mrk 1044 .
99 1 . ± .
03 0 . ± .
03 0 . ± .
003 23 . ± .
1, 2, 11 a
3C 120 .
39 1 . ± .
05 0 . ± . − . ± .
020 17 . ± .
1, 12 a IRAS 04416+1215 .
96 1 . ± .
06 0 . ± .
04 0 . ± . too weak 1, 2, 11 a , 13 a Ark 120 .
83 1 . ± .
01 0 . ± . − . ± . too weak 1, 12 a , 14 a MCG +08-11-011 .
29 1 . ± .
11 0 . ± . − . ± .
020 14 . ± . a Mrk 374 .
88 1 . ± .
10 0 . ± . − . ± .
016 5 . ± . a Mrk 79 .
33 2 . ± .
06 0 . ± . − . ± .
018 22 . ± .
1, 16 a SDSS J074352 .
11 1 . ± .
02 0 . ± . − . ± . too weak 11 a , 17 a SDSS J075051 .
22 1 . ± .
01 1 . ± .
06 0 . ± .
018 5 . ± . a , 17 a SDSS J075101 .
97 1 . ± .
08 1 . ± .
10 0 . ± .
020 5 . ± .
1, 11 a Mrk 382 .
75 1 . ± .
36 0 . ± .
24 0 . ± .
019 15 . ± .
1, 2, 11 a , 13SDSS J075949 .
02 1 . ± .
04 0 . ± .
03 0 . ± .
011 15 . ± . a , 17 a SDSS J080101 .
01 1 . ± .
08 0 . ± .
02 0 . ± .
005 6 . ± .
1, 11 a , 18SDSS J080131 .
49 1 . ± .
03 1 . ± .
07 0 . ± .
012 5 . ± .
1, 11 a , 19PG 0804+761 .
61 2 . ± .
04 0 . ± . − . ± .
001 1 . ± .
4, 5 a , 6SDSS J081441 .
46 1 . ± .
08 0 . ± . − . ± .
003 20 . ± .
1, 11 a , 17SDSS J081456 .
31 1 . ± .
09 0 . ± .
01 0 . ± .
007 8 . ± .
1, 11 a , 18NGC 2617 .
31 2 . ± .
18 0 . ± . − . ± . too weak 12, 16 a , 20SDSS J083553 .
57 1 . ± .
02 0 . ± .
02 0 . ± .
006 12 . ± .
11, 17 a SDSS J084533 .
11 1 . ± .
03 0 . ± .
03 0 . ± .
017 5 . ± .
1, 11 a , 19PG 0844+349 .
78 1 . ± .
04 0 . ± .
02 0 . ± .
009 9 . ± .
1, 4, 6, 21 a SDSS J085946 .
39 1 . ± .
04 0 . ± .
03 0 . ± .
020 5 . ± .
1, 11 a , 19Mrk 110 .
14 1 . ± .
09 0 . ± . − . ± .
005 21 . ± .
1, 21 a SDSS J093302 .
44 1 . ± .
02 1 . ± .
04 0 . ± .
005 16 . ± . a , 17SDSS J093922 .
48 1 . ± .
06 1 . ± . − . ± .
021 16 . ± .
1, 11 a , 18PG 0953+414 .
27 1 . ± .
04 0 . ± . − . ± .
004 7 . ± .
1, 4, 6, 21 a SDSS J100402 .
17 1 . ± .
01 0 . ± .
02 0 . ± .
007 0 . ± . a , 17SDSS J101000 .
17 1 . ± .
00 0 . ± .
03 0 . ± .
015 4 . ± . a , 17NGC 3227 .
46 2 . ± .
17 0 . ± . − . ± .
002 33 . ± .
1, 7, 9 a , 22 a SDSS J102339 .
03 1 . ± .
04 0 . ± . − . ± .
012 20 . ± .
1, 11 a , 19Mrk 142 .
14 1 . ± .
26 0 . ± .
06 0 . ± .
005 14 . ± .
1, 11 a NGC 3516 .
66 2 . ± .
17 0 . ± . − . ± .
004 47 . ± .
1, 7, 9 a SBS 1116+583A .
59 2 . ± .
13 0 . ± . − . ± .
008 31 . ± .
1, 11 a , 23Arp 151 .
32 1 . ± .
04 0 . ± . − . ± .
008 41 . ± .
1, 23 a NGC 3783 .
04 2 . ± .
05 0 . ± . − . ± .
011 57 . ± .
1, 4, 24, 25 a MCG +06-26-012 .
04 1 . ± .
11 0 . ± .
07 0 . ± .
001 16 . ± .
1, 2, 13 a UGC 06728 .
11 0 . ± .
12 0 . ± . − . ± .
020 14 . ± . a Mrk 1310 .
46 1 . ± .
07 0 . ± . − . ± .
002 32 . ± .
1, 23 a NGC 4051 .
18 1 . ± .
15 1 . ± .
12 0 . ± .
009 8 . ± .
1, 7, 27, 28 a NGC 4151 .
22 2 . ± .
07 0 . ± .
02 0 . ± .
017 21 . ± .
1, 4, 9 a , 29PG 1211+143 .
42 1 . ± .
04 0 . ± .
03 0 . ± .
007 15 . ± .
1, 4, 6, 21 a Mrk 202 .
57 1 . ± .
08 0 . ± . − . ± .
006 21 . ± .
1, 11 a , 23NGC 4253 .
99 1 . ± .
06 0 . ± . − . ± .
001 29 . ± .
1, 23 a PG 1226+023 .
64 1 . ± .
03 0 . ± . − . ± .
004 0 . ± . a PG 1229+204 .
53 2 . ± .
05 0 . ± . − . ± .
007 13 . ± .
1, 4, 6, 21 a NGC 4593 .
89 2 . ± .
66 0 . ± .
07 0 . ± .
020 6 . ± .
1, 31 a IRAS F12397+3333 .
48 1 . ± .
52 0 . ± .
30 0 . ± .
010 19 . ± .
1, 2, 3, 11 a Table 2 continued Table 2 (continued)
Objects R Fe D H β FWHM Fe / FWHM H β Asymmetry EW(He II ) Ref.( ˚A)NGC 4748 .
99 1 . ± .
08 0 . ± .
05 0 . ± .
002 34 . ± .
1, 23 a PG 1307+085 .
21 2 . ± .
09 0 . ± . − . ± .
006 15 . ± .
1, 4, 6 a MCG +06-30-015 .
93 2 . ± .
08 1 . ± .
49 0 . ± .
020 16 . ± . a , 32NGC 5273 .
58 3 . ± .
13 0 . ± .
04 0 . ± . too weak 1, 22 a , 33Mrk 279 .
55 2 . ± .
03 0 . ± .
10 0 . ± .
006 7 . ± .
1, 4, 34, 35 a PG 1411+442 .
63 1 . ± .
04 0 . ± . − . ± .
004 11 . ± .
1, 4, 5 a , 6NGC 5548 .
10 2 . ± .
33 0 . ± . − . ± .
012 8 . ± .
1, 4, 11 a , 36PG 1426+015 .
46 2 . ± .
09 0 . ± . − . ± .
003 10 . ± .
1, 4, 6 a Mrk 817 .
69 2 . ± .
29 0 . ± . − . ± .
057 22 . ± .
1, 4, 37, 38 a Mrk 1511 .
80 2 . ± .
13 0 . ± . − . ± .
020 6 . ± .
1, 31 a Mrk 290 .
29 2 . ± .
18 0 . ± . − . ± .
004 21 . ± .
1, 7, 21, 23 a Mrk 486 .
54 1 . ± .
06 0 . ± .
06 0 . ± .
003 15 . ± .
1, 2, 11 a , 13 a Mrk 493 .
13 1 . ± .
03 1 . ± .
02 0 . ± .
018 5 . ± .
1, 2, 11 a , 13 a PG 1613+658 .
38 2 . ± .
05 0 . ± . − . ± .
003 7 . ± .
1, 4, 6 a , 21PG 1617+175 .
74 2 . ± .
12 0 . ± . − . ± .
020 1 . ± .
1, 4, 6 a PG 1700+518 .
32 1 . ± .
08 0 . ± .
06 0 . ± . too weak 1, 4, 5, 6, 27, 393C 382 .
31 1 . ± .
13 0 . ± . − . ± . too weak 15, 35 a
3C 390.3 .
12 2 . ± .
66 0 . ± . − . ± . too weak 1, 18, 40, 41KA 1858+4850 .
11 2 . ± .
13 0 . ± .
05 0 . ± .
020 35 . ± .
1, 42 a NGC 6814 .
45 1 . ± .
03 0 . ± . − . ± .
020 23 . ± .
1, 18, 23 a Mrk 509 .
13 1 . ± .
01 0 . ± . − . ± .
022 44 . ± .
1, 4, 9 a , 37PG 2130+099 .
96 1 . ± .
11 0 . ± .
06 0 . ± .
001 14 . ± .
1, 5 a , 43NGC 7469 .
43 1 . ± .
18 0 . ± . − . ± .
032 17 . ± .