A DIRECT LINKAGE BETWEEN AGN OUTFLOWS IN THE NARROW-LINE REGIONS AND THE X-RAY EMISSION FROM THE ACCRETION DISKS

It is now generally believed that the feedback from central active galactic nuclei (AGNs) plays an important role in galaxy evolution (see reviews in Veilleux et al. 2005, Fabian 2012, and Heckman & Best 2014). Self-regulated super-massive black hole (SMBH) growth and star formation in the host galaxy is potentially realized by sweeping out circumnuclear gas through the feedback process in both galaxy merger and secular evolution scenarios (e.g., Alexander & Hickox 2012; Kormendy & Ho 2013). Both semianalytic models and numerical simulations indicate that this feedback process is useful not only for reproducing the observed M_BH–σ_* relation, luminosity functions of quasars, and normal galaxies (e.g., Haehnelt et al. 1998; Silk & Rees 1998; Fabian 1999; Kauffmann & Haehnelt 2000; Granato et al. 2004; Di Matteo et al. 2005, 2007; Springel et al. 2005; Croton et al. 2006; Hopkins et al. 2007, 2008; Khalatyan et al. 2008; Menci et al. 2008; Somerville et al. 2008), but also for solving the "over cooling" problem in the Λ cold dark matter (ΛCDM) galaxy formation model in which the cooling predicted in galaxy groups and clusters is stronger than the observed one (e.g., Ciotti & Ostriker 2007; Somerville et al. 2008, Hirschmann et al. 2014). By using the SDSS spectroscopic survey, Wang et al. (2011) and Wang (2015) recently suggested a coevolution between the feedback and host galaxy based on the fact that an AGN that has an [O iii]λ5007 emission line with stronger blue asymmetry tends to be associated with a younger stellar population.

Various methods of the feedback process have been proposed in past decades by a mixture of observational and theoretical studies. These methods include AGN wind (e.g., Crenshaw et al. 2003; Pounds et al. 2003; Ganguly et al. 2007; Reeves et al. 2009; Dunn et al. 2010; Tombesi et al. 2012), radiation pressure (e.g., Granato et al. 2004; Alexander et al. 2010), and mechanical energy outflow caused by a collimated radio jet (e.g., Best et al. 2006; Holt et al. 2008; Nesvadba et al. 2008; Rosario et al. 2010; Guillard et al. 2012).

From an observational standpoint, there is accumulating evidence that feedback from the central SMBH can drive outflows on various scales (see reviews in Veilleux et al. 2005 and Fabian 2012). In addition to the blueshifted absorption lines in optical, UV, and soft-X-ray spectra (e.g., Crenshaw et al. 2003; Hamann & Sabra 2004; Wang & Xu 2015 and references therein), the outflows can be conveniently traced by the blue asymmetry of the [O iii]λλ4959, 5007 doublet and its bulk blueshift with respect to the local system (Heckman et al. 1981; Veron-Cetty et al. 2001; Zamanov et al. 2002; Marziani et al. 2003; Aoki et al. 2005; Bian et al. 2005; Boroson 2005; Komossa et al. 2008; Xu & Komossa 2009; Mullaney et al. 2013; Zhang et al. 2013).

Using the [O iii] line profile as a diagnostic of the outflow occurring in the AGN's narrow-line region (NLR) enables studies of the origin of the outflow (feedback) based on the large optical spectroscopic sample. Different results are, however, obtained by various authors. On the one hand, there is ample evidence that the [O iii] line blue asymmetry and bulk velocity blueshift are strongly correlated with the Eddington ratio ($L/{L}_{\mathrm{Edd}}$, where ${L}_{\mathrm{Edd}}=1.26\times {10}^{38}({M}_{{}_{\mathrm{BH}}}/{M}_{\odot })\ \mathrm{erg}\ {{\rm{s}}}^{-1}$ is the Eddington luminosity) of the central AGNs: the higher the $L/{L}_{\mathrm{Edd}}$, the stronger the blue asymmetry and the larger the bulk velocity blueshift will be (e.g., Bian et al. 2005; Boroson 2005; Zhang et al. 2011; Wang et al. 2011; Wang 2015). In addition, the extreme "blue outliers," usually defined as the objects with [O iii] bulk blueshift larger than $250\;\mathrm{km}\;{{\rm{s}}}^{-1}$ (e.g., Zamanov et al. 2002; Zhou et al. 2006; Komossa et al. 2008), are found to exclusively occur in the AGNs associated with high $L/{L}_{\mathrm{Edd}}$. On the other hand, the relation between the [O iii] line profile and radio emission has been reported in the previous studies. Early studies show that the [O iii] line width is found to be correlated with radio luminosity at 1.4 GHz (${L}_{1.4\mathrm{GHz}}$) for a sample of flat-spectrum radio galaxies (e.g., Heckman et al. 1984; Whittle 1985). Mullaney et al. (2013) and Zakamska & Greene (2014) recently analyzed the [O iii] line profile in a large sample of both type I and type II AGNs detected by the SDSS spectroscopic survey. Their results suggest that the [O iii] line width is more strongly related to ${L}_{1.4\mathrm{GHz}}$ than the other AGN parameters (e.g., line luminosity and $L/{L}_{\mathrm{Edd}}$).

The origin of feedback in local type I AGNs is studied in this paper by focusing on the relationship between the outflow traced by the [O iii] line profile and the X-ray emission from the central AGNs. AGNs are well known to be luminous X-ray emitters up to 100 keV. X-ray emission is a powerful tool for identifying nuclear SMBH accretion activity and for studying the accretion process that fuels AGNs, because the luminous X-ray emission from a rapidly accreting AGN is produced in the region very close to the SMBH (e.g., Haardt & Maraschi 1991; Zdziarski et al. 2000; Kawaguchi et al. 2001; Cao 2009).

The paper is organized as follows. The sample selection and spectral analysis are presented in Sections 2 and 3, respectively. The statistical results are shown in Section 4, and the implications are discussed in Section 5. A ΛCDM cosmology with parameters ${H}_{0}=70\ \mathrm{km}\ {{\rm{s}}}^{-1}\ {\mathrm{Mpc}}^{-1}$, ${{\rm{\Omega }}}_{{\rm{m}}}=0.3$, and ${{\rm{\Omega }}}_{{\rm{\Lambda }}}=0.7$ is adopted throughout the paper.

The sample of X-ray-selected type-I AGNs used in the current study is extracted from the 2XMMi/SDSS-DR7 catalog, which was originally cross-matched by Pineau et al. (2011) between the incremental Second XMM-Newton Serendipitous Source Catalog (2XMMi, Watson et al. 2009) and the SDSS-DR7 catalog (Abazajian et al. 2009). In the cross-match, the optical counterpart of an X-ray source is identified by the probability of spatial coincidence estimated from the traditionally adopted likelihood ratio estimator. There are in total 221, 012 unique, serendipitous X-ray sources in the 2XMMi catalog. A ∼90% completeness can be achieved for the catalog at a sensitivity of $1\times {10}^{-14}\ \mathrm{erg}\ {{\rm{s}}}^{-1}\ {\mathrm{cm}}^{-2}$ and $9\times {10}^{-14}\ \mathrm{erg}\ {{\rm{s}}}^{-1}\ {\mathrm{cm}}^{-2}$ in the 0.5–2.0 keV and 2.0–12.0 keV bandpasses, respectively. The localization accuracy of the X-ray sources is typical of $2^{\prime\prime}$. More than 30,000 X-ray point-like sources with a localization accuracy $\geqslant 5^{\prime\prime}$ have an SDSS-DR7 optical counterpart with an identification probability larger than 90%.

A sub-sample of X-ray type-I AGNs is selected from the 2XMMi/SDSS-DR7 catalog by requiring that (1) the probability of an identification is no less than 95%, (2) the angular distance between an individual XMM-Newton X-ray source and its corresponding optical counterpart is less than 3'', taking into account the SDSS fiber aperture, (3) the redshift is smaller than 0.2, (4) the g-band brightness is brighter than 19 mag, which is necessary for a proper modeling of the optical continuum, and (5) the X-ray flux in the 0.2–12 keV bandpass is larger than $1\times {10}^{-14}\ \mathrm{erg}\ {{\rm{s}}}^{-1}\ {\mathrm{cm}}^{-2}$, which excludes the extracted X-ray spectra with low photon count rates. By focusing on the objects that are classified as quasars according to the spectral type classification given by the SDSS pipelines (Bromley et al. 1998; Glazebrook et al. 1998), there are a total of 82 type-I AGNs fulfilling the above selection criteria after removing SDSS J093249.57+472522.8, which shows a spectrum typical of a CV star and is incorrectly classified as a quasar by the SDSS pipelines (Yip et al. 2004). By examining the SDSS spectra one by one by eye, 19 out of the 82 sources are then excluded from the subsequent optical and X-ray spectral analysis, because of their Seyfert-1.8-like spectra (i.e., partially obscured AGNs) in which the continuum is dominated by the starlight from their host galaxies.

3.1. XMM-Newton EPIC Spectra

Our X-ray spectral analysis focuses on the XMM-Newton EPIC PN (Struder et al. 2001) data. The data are reduced by the SAS v11.0 software³ and by the corresponding calibration files. The events with patterns of 0–4 are extracted from the PN data for all the 63 X-ray type-I AGNs selected in Section 2, except for one object. Single pixel events (i.e., pattern = 0) are extracted for one bright object (SDSS J103438.59+393828.2 = KUG 1031+398) since these events are less sensitive to pile-up than other patterns. In the extraction, the bad and hot pixels are removed from the original image, and the CCD chip gaps are avoided. The source spectrum is extracted from a circular aperture at the detected source position. The aperture has a radius of 25''–40'' depending on the brightness of the object. The background is determined from a circular source-free region that is offset from, but close to, the source. The pile-up in the data is checked by the SAS task epatplot. The tasks rmfgen and arfgen are used to generate the needed response files.

There are in total 47 X-ray selected type-I AGNs for the subsequent X-ray spectral modelings. The other 16 objects are excluded either because of the coincidence of the CCD gap or because of the bad X-ray spectral quality beyond 2–3 keV due to their faintness. We fit the extracted spectra over the 0.3–10 keV band by the XSPEC package (Arnaud 1996). The absorption caused by our own galaxy is included in the spectral fitting for each object by taking the column density value from the Leiden/Argentine/Bonn (LAB) Survey (Kalberla et al. 2005). A basic model expressed as wabs*zwabs* (N×blackbody+powerlaw) is adopted in the fitting for all of the objects. The power-law photon spectrum is defined as $N(E)\propto {E}^{-{\rm{\Gamma }}}$, where E is the photon energy and Γ is the photon index. We also attempt to reproduce each of the spectra by the neutral reflection model (pexrav; Magdziarz & Zdziarski 1995) instead of the simple power law, though no significant improvement (i.e., the difference of the reduced ${\chi }^{2}$ is less than 10%) can be obtained from this complicated model. In addition to the basic model, additional components are required for some objects to reproduce the observed spectra. A Gaussian profile is required in three objects (SDSS J091848.61+211717.0, SDSS J140700.40+282714.6 = Mark 668, and SDSS J143452.45+483942.7 = NGC 5683) to model their broad iron Kα emission lines at 6.4 keV (rest frame). At the low-energy end, the X-ray spectra of SDSS J091848.61+211717.0, SDSS J111706.39+441333.3 (PG 1114+445), and SDSS J124210.61+331702.6 (WAS 61) are best fitted by an edge-like absorption (zedge) due to the K-edge of ions. The best-fit edge energy ${E}_{\mathrm{edge}}$ in the rest frame is 0.67 ± 0.04, $0.72\pm 0.01\ \mathrm{keV}$, and $0.73\pm 0.01\ \mathrm{keV}$ for SDSS J091848.61+211717.0, SDSS J111706.39+441333.3, and SDSS J124210.61+331702.6, respectively. These modeled edge energies suggest that the edge-like absorptions in the three objects are likely caused by the K-shell absorption of O vii ions at 0.74 keV. Our best-fit model of PG 1114+445 is highly consistent with that obtained from the ROSAT PSPC observation (Laor et al. 1994). As an illustration, Figure 1 shows the X-ray spectral modelings for four objects.⁴

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3.2. SDSS Optical Spectroscopy

We analyze the one-dimensional optical spectra of the 47 X-ray-selected type-I AGNs using the IRAF⁵ package as follows. At the beginning, the correction of galactic extinction is applied to each spectrum according to the color excess, the parameter $E(B-V)$ taken from the Schlegel, Finkbeiner, and Davies Galactic reddening map (Schlegel et al. 1998). An extinction law of the Milky Way with ${R}_{{\rm{V}}}=3.1$ (Cardelli et al. 1989) is adopted in the correction. Each spectra is then transformed to the rest frame, along with the flux correction due to the relativity effect, according to the redshift provided by the SDSS pipelines.

3.2.1. Continuum Modeling and Removal

The continuum of each rest-frame spectrum is modeled by the linear sum of several components: (1) a broken power law from a central AGN, in which the wavelength of the break point and the two spectral indices are not fixed in the continuum modeling, (2) a template of both high-order Balmer emission lines and a Balmer continuum from the broad-line region (BLR), (3) an empirical template of the optical Fe ii complex, and (4) the eigenspectra of the host galaxy that are built from the standard single stellar population spectral library through the principal component analysis method (e.g., Francis et al. 1992; Hao et al. 2005; Wang & Wei 2008). The used spectral library was developed by Bruzual & Charlot (2003). A galactic extinction curve with R_V = 3.1 is involved in the modeling to account for the intrinsic extinction due to the host galaxy. A ${\chi }^{2}$ minimization is iteratively performed over the rest-frame wavelength range from 3700 to 7500 Å, except for the regions with strong emission lines. Figure 2 illustrates the modeling and removal of the continuum for two typical objects.

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The template of the high-order Balmer lines (i.e., ${{\rm{H}}}_{7}-{{\rm{H}}}_{50}$) is taken from the case B recombination model with an electron temperature of ${T}_{e}=1.5\times {10}^{4}\;{\rm{K}}$ and an electron density of ${n}_{e}={10}^{8-10}\ {\mathrm{cm}}^{-3}$. The model is calculated by Storey & Hummer (1995). The line width of the high-order Balmer lines is fixed to be that of the broad component of Hβ, which is determined by our line profile modeling (see below). The Balmer continuum is approximately modeled by the emission from a partially optically thick cloud by following Dietrich et al. (2002), see also in Grandi (1982) and Malkan & Sargent (1982):

$$\begin{eqnarray}&&{f}_{\lambda }^{\mathrm{BC}}={f}_{\lambda }^{\mathrm{BE}}{B}_{\lambda }({T}_{e})(1-{e}^{-{\tau }_{\lambda }});\lambda \leqslant {\lambda }_{\mathrm{BE}}\end{eqnarray} \tag{ 1 }$$

where ${f}_{\lambda }^{\mathrm{BE}}$ is the continuum flux at the Balmer edge ${\lambda }_{\mathrm{BE}}=3646$ Å, ${B}_{\lambda }({T}_{e})$ is the Planck function at an electron temperature⁶ of ${T}_{e}=1.0\times {10}^{4}$ K, and ${\tau }_{\lambda }$ is the optical depth at wavelength λ. The optical depth at λ is related to the one at the Balmer edge ${\tau }_{\mathrm{BE}}$ as ${\tau }_{\lambda }={\tau }_{\mathrm{BE}}{(\lambda /{\lambda }_{\mathrm{BE}})}^{3}$. ${\tau }_{\mathrm{BE}}$ ranges from 0.1 to 2.0, and a typical value of ${\tau }_{\mathrm{BE}}=0.5$ is adopted in the current Balmer continuum fitting.

We model the optical Fe ii complex by using the empirical templates provided in Veron-Cetty et al. (2004). Both the broad and narrow components of the Fe ii templates are included in the modeling. Again, the line widths of the broad and narrow Fe ii emission are determined from the line profile modeling of the Hβ emission line (see below). With the modeled Fe ii complex, its flux (i.e., Fe iiλ4570) is measured from the modeled spectrum in the rest-frame wavelength range from 4434 to 4684 Å.

3.2.2. Emission Line Profile Modeling

After the removal of the continuum, the emission-line profiles are modeled on each emission-line isolated spectrum for both Hα and Hβ regions (i.e., in the wavelength ranges of λλ6350–6750 and λλ4800–5050) by the SPECFIT task (Kriss 1994) in the IRAF package. For each object, each emission line is profiled by a linear combination of a set of several Gaussian functions. The line flux ratios of the [O iii] and [N ii] doublets are fixed to their theoretical values. The line width of the narrow Hα (Hβ) component is fixed to equal that of the [N ii] ([O iii] core) line, if the resulting two widths are significantly different. In order to properly isolate the [O iii]$\lambda 5007$ line profile, the broad He i $\lambda 5016$ emission line (Veron et al. 2002) is additionally required in the line profile modelings in three objects (i.e., SDSS J111830.28+402554.0, SDSS J134022.86+274058.5, and SDSS J155909.62+350147.4).

The line modelings are schematically presented in the left and right panels of Figure 3 for the Hα and Hβ regions, respectively. As shown in the figure, a linear combination of two or three broad Gaussian functions is usually required to adequately reproduce the observed broad Balmer line profile in most cases. A residual line profile, which is obtained by subtracting the modeled narrow-line component (including the modeled forbidden lines) from the observed profile, is then used to measure the line width and integrated line flux of either broad Hα or Hβ emission lines. Figure 4 schematically shows the [O iii]λ5007 line profiles modeled by a sum of n Gaussian functions for four typical cases. As illustrated by the figure, generally speaking, a sum of two or three Gaussian functions are adequate to model all of the observed [O iii] line profiles well.

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Table 1 lists the results obtained from the above X-ray and optical spectral modelings for the 47 X-ray selected type-I AGNs. The identification of each object and the corresponding redshift given by the SDSS pipelines are listed in columns (1) and (2), respectively. The full width at half maximum (FWHM) of the broad Hβ emission is tabulated in column (3). Columns (4) and (11) list the modeled broad Hα emission-line luminosity (${L}_{{\rm{H}}\alpha }$) and the spectral-fitting-inferred intrinsic X-ray luminosity in a bandpass of 2–10 keV (${L}_{2-10\mathrm{keV}}$), respectively. The calculated line luminosities are corrected for the local extinction that is inferred from the narrow-line ratio Hα/Hβ by assuming a Balmer decrement for standard case B recombination and a Galactic extinction curve with R_V = 3.1. Figure 5 shows a tight relationship between ${L}_{{\rm{H}}\alpha }$ and ${L}_{2-10\mathrm{keV}}$, which is not related to AGN NLR. A distribution of the inferred ${L}_{2-10\mathrm{keV}}$ is plotted in the inert panel of the figure, which shows a range from 10⁴² to ${10}^{44}\;\mathrm{erg}\ {{\rm{s}}}^{-1}$ for ${L}_{2-10\mathrm{keV}}$. The parameter of RFe defined as the flux ratio between the Fe iiλ4570 and Hβ broad component is tabulated in column (5). The fitted spectral photon index ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ and the corresponding errors at a confidence level of 90% are listed in column (6). These errors reported by the XSPEC package are obtained from our spectral modelings. The measured ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ has an average value of 1.86 and a standard deviation of 0.31, which are highly consistent with the typical value of ${\rm{\Gamma }}\sim 1.9$ for radio-quiet AGNs (e.g., Zdziarski et al. 1995; Reeves & Turner 2000; Piconcelli et al. 2005; Dadina 2008; Panessa et al. 2008; Mateos et al. 2010; Zhou & Zhang 2010; Corral et al. 2011).

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Table 1.  Properties of the XMM-Newton/SDSS-DR7 Type-I AGNs

SDSS	z	${\mathrm{FWHM}}_{{\rm{H}}\beta }$	$\mathrm{log}\displaystyle \frac{{L}_{{\rm{H}}\alpha }}{\mathrm{erg}\;{{\rm{s}}}^{-1}}$	RFe	${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$	${\rm{\Delta }}\upsilon$ a	$\mathrm{log}\displaystyle \frac{{M}_{\mathrm{BH}}}{{M}_{\odot }}$	$L/{L}_{\mathrm{Edd}}$	$\delta \upsilon$	$\mathrm{log}\displaystyle \frac{{L}_{{\rm{X}}}}{\mathrm{erg}\ {{\rm{s}}}^{-1}}$
		$\mathrm{km}\ {{\rm{s}}}^{-1}$				$\mathrm{km}\ {{\rm{s}}}^{-1}$			$\mathrm{km}\ {{\rm{s}}}^{-1}$
(1)	(2)	(3)	(4)	(5)	(6)	(7)	(8)	(9)	(10)	(11)
J010712.04+140844.9	0.0769	1420 ± 220	41.3	0.55 ± 0.09	${2.15}_{-0.10}^{+0.10}$	−25.1 ± 39.7	6.19	0.33	−39.1 ± 13.6	42.5
J015950.24+002340.8	0.1627	3240 ± 150	42.7b	0.57 ± 0.03	${1.96}_{-0.19}^{+0.17}$	56.3 ± 25.0	8.09	0.19	−332.1 ± 185.5	43.7
J030639.58+000343.2	0.1074	1970 ± 50	43.7	0.003 ± 0.002	${1.80}_{-0.04}^{+0.04}$	108.3 ± 34.9	8.13	0.38	−85.0 ± 40.9	43.4
J091848.61+211717.0	0.1493	1740 ± 190	43.8	0.05 ± 0.01	${2.04}_{-0.19}^{+0.19}$	103.7 ± 45.9	7.80	1.03	−212.8 ± 37.3	42.9
J092247.02+512038.0	0.1598	2610 ± 150	42.3	2.46 ± 0.23	${2.23}_{-0.10}^{+0.10}$	−361.2 ± 45.1	6.92	0.44	−349.1 ± 169.6	43.4
J092343.00+225432.6	0.0332	2750 ± 100	41.5	0.25 ± 0.02	${1.90}_{-0.02}^{+0.02}$	−42.6 ± 17.6	6.86	0.09	−42.4 ± 29.6	43.6
J093922.89+370943.9	0.1859	1900 ± 150	42.6	0.61 ± 0.05	${2.10}_{-0.39}^{+0.35}$	−29.0 ± 42.7	7.13	0.50	−424.7 ± 362.5	43.2
J094439.88+034940.1	0.1554	4410 ± 30	43.0	0.49 ± 0.06	${1.90}_{-0.18}^{+0.19}$	−31.4 ± 68.6	8.24	0.08	−165.3 ± 108.6	43.2
J100035.47+052428.5*	0.0786	4020 ± 130	41.8	0.06 ± 0.05	${1.41}_{-0.22}^{+0.22}$	22.7 ± 8.9	6.78	0.20	−7.0 ± 3.9	43.2
J102822.84+235125.7*	0.1734	5340 ± 200	42.5	0.01 ± 0.01	${1.89}_{-0.20}^{+0.19}$	98.2 ± 18.8	8.05	0.04	+5.5 ± 23.8	43.4
J103059.09+310255.7*	0.1781	5560 ± 330	43.5	0.00 ± 0.00	${1.59}_{-0.07}^{+0.07}$	−0.6 ± 27.5	8.71	0.07	+22.2 ± 12.9	44.3
J103349.93+631830.4	0.1555	1950 ± 410	42.8	0.03 ± 0.01	${2.11}_{-0.13}^{+0.14}$	⋯	7.48	0.28	−42.9 ± 40.8	43.1
J103438.59+393828.2	0.0431	870 ± 90	42.1	0.16 ± 0.02	${2.56}_{-0.17}^{+0.17}$	−27.3 ± 15.7	6.61	0.53	−191.0 ± 32.9	42.3
J105143.89+335926.7	0.1671	3880 ± 160	43.3	0.18 ± 0.01	${1.92}_{-0.04}^{+0.04}$	75.2 ± 41.0	8.15	0.17	−1.1 ± 5.7	44.0
J110101.77+110248.9*	0.0356	7210 ± 250	41.9	0.01 ± 0.01	${1.66}_{-0.03}^{+0.05}$	0.1 ± 16.1	7.79	0.02	+5.5 ± 8.0	43.0
J111706.39+441333.3	0.1438	5180 ± 170	44.2	0.02 ± 0.01	${1.52}_{-0.02}^{+0.02}$	−12.6 ± 49.0	8.88	0.21	−124.7 ± 83.5	44.0
J111830.28+402554.0	0.1545	2100 ± 110	43.6	0.43 ± 0.02	${2.14}_{-0.11}^{+0.10}$	−109.8 ± 40.4	7.70	0.54	−116.1 ± 48.2	43.8
J112328.11+052823.2	0.1013	1660 ± 350	42.2	0.18 ± 0.04	${2.00}_{-0.11}^{+0.11}$	11.5 ± 35.4	7.04	0.25	−71.6 ± 7.7	42.8
J114008.71+030711.4	0.0811	1350 ± 80	41.5	0.78 ± 0.07	${1.84}_{-0.28}^{+0.27}$	−3.7 ± 22.9	6.52	0.22	−62.0 ± 34.4	42.6
J120442.10+275411.7*	0.1651	5080 ± 420	43.3	0.06 ± 0.02	${1.58}_{-0.08}^{+0.07}$	38.4 ± 9.4	8.28	0.12	−232.0 ± 27.7	44.4
J121356.19+140431.3	0.1539	4970 ± 460	42.6	0.15 ± 0.01	${1.48}_{-0.20}^{+0.18}$	⋯	8.20	0.04	−53.7 ± 113.3	43.4
J121930.87+064334.4	0.0804	1780 ± 70	43.0	0.08 ± 0.01	${2.11}_{-0.08}^{+0.08}$	11.5 ± 43.9	7.42	0.54	−61.0 ± 22.5	43.0
J122137.93+043026.1*	0.0947	8530 ± 410	42.0	0.25 ± 0.01	${1.43}_{-0.19}^{+0.17}$	21.6 ± 15.9	8.18	0.01	+2.1 ± 4.3	42.9
J123113.66+151127.9	0.1919	3500 ± 190	42.3	0.47 ± 0.03	${2.20}_{-0.15}^{+0.16}$	17.9 ± 44.1	7.48	0.12	−83.7 ± 48.5	43.3
J124013.80+473354.8	0.1174	1840 ± 240	41.7	0.75 ± 0.06	${1.48}_{-0.48}^{+0.45}$	4.7 ± 51.2	6.72	0.20	−33.1 ± 10.7	43.0
J124210.60+331702.6	0.0437	1470 ± 120	43.0	0.06 ± 0.01	${2.15}_{-0.01}^{+0.01}$	−2.9 ± 35.1	7.32	0.60	−173.1 ± 44.5	43.3
J124635.24+022208.7	0.0482	1270 ± 90	41.9	0.45 ± 0.05	${2.36}_{-0.10}^{+0.10}$	−118.4 ± 32.8	6.72	0.31	−62.3 ± 43.8	43.0
J130022.15+282402.6	0.0911	3550 ± 140	42.3	0.39 ± 0.01	${1.70}_{-0.12}^{+0.11}$	7.5 ± 48.7	7.46	0.12	−7.9 ± 8.2	43.4
J130947.00+081948.2	0.1543	4320 ± 130	43.7	0.10 ± 0.01	${1.46}_{-0.04}^{+0.04}$	41.4 ± 24.6	8.43	0.20	−137.1 ± 25.9	44.0
J133141.02–015212.4	0.1454	1570 ± 280	42.1	0.24 ± 0.05	${1.72}_{-0.41}^{+0.39}$	−38.7 ± 31.0	7.01	0.22	−27.4 ± 26.8	43.0
J134351.06+000434.7*	0.0737	3300 ± 150	41.7	0.69 ± 0.01	${1.43}_{-0.23}^{+0.23}$	18.3 ± 30.8	6.84	0.16	+0.9 ± 4.5	41.6
J134834.94+263109.8*	0.0589	1580 ± 90	42.3	0.11 ± 0.01	${1.79}_{-0.23}^{+0.23}$	11.7 ± 37.6	7.11	0.28	−119.5 ± 56.9	42.7
J135435.68+180517.4	0.1509	4190 ± 110	43.9	0.04 ± 0.01	${1.83}_{-0.11}^{+0.11}$	⋯	8.52	0.25	−142.5 ± 48.9	44.0
J135553.52+383428.7*	0.0502	6500 ± 220	42.4	0.04 ± 0.01	${1.56}_{-0.07}^{+0.07}$	22.6 ± 7.3	7.75	0.07	−25.3 ± 7.4	43.2
J135724.51+652505.9	0.1063	1590 ± 190	41.6	0.30 ± 0.05	${1.96}_{-0.23}^{+0.18}$	−15.5 ± 27.9	6.40	0.36	−13.2 ± 9.8	42.9
J140251.19+263117.5	0.1875	5900 ± 80	44.1	0.07 ± 0.01	${1.49}_{-0.10}^{+0.10}$	104.5 ± 33.2	9.01	0.11	−84.3 ± 26.4	44.2
J140621.89+222346.5	0.0979	4280 ± 110	42.5	1.43 ± 0.04	${1.95}_{-0.42}^{+0.44}$	−47.2 ± 44.3	7.83	0.07	−283.1 ± 105.0	42.5
J140700.40+282714.6	0.0766	8240 ± 450	43.5	0.09 ± 0.05	${1.22}_{-0.10}^{+0.10}$	61.5 ± 20.6	9.18	0.03	−29.6 ± 15.0	42.5
J141519.50–003021.5	0.1347	2260 ± 190	41.8	1.21 ± 0.11	${2.15}_{-0.09}^{+0.09}$	−304.5 ± 58.3	6.94	0.14	−99.6 ± 67.6	42.9
J141700.82+445606.3	0.1136	2590 ± 150	42.8	0.86 ± 0.07	${1.93}_{-0.09}^{+0.09}$	⋯	7.73	0.18	−276.9 ± 49.7	43.5
J143452.45+483942.7	0.0365	4520 ± 190	43.0	0.004 ± 0.003	${1.58}_{-0.23}^{+0.21}$	−39.5 ± 45.4	8.23	0.08	+14.8 ± 6.3	43.1
J145108.76+270926.9	0.0645	3040 ± 100	42.6	0.79 ± 0.03	${2.27}_{-0.06}^{+0.06}$	9.9 ± 12.9	7.50	0.20	−103.7 ± 10.6	43.3
J150626.44+030659.9	0.1734	2270 ± 400	42.8	0.15 ± 0.01	${1.79}_{-0.25}^{+0.02}$	⋯	7.49	0.28	−147.7 ± 20.8	43.6
J151600.96+000949.7	0.1712	4440 ± 190	42.3	0.09 ± 0.11	${1.68}_{-0.07}^{+0.07}$	⋯	7.81	0.05	−45.2 ± 38.0	43.4
J155909.63+350147.5	0.0311	1700 ± 110	41.7	1.53 ± 0.14	${2.15}_{-0.07}^{+0.06}$	−70.9 ± 21.7	6.42	0.41	−18.7 ± 23.2	42.9
J160452.45+240241.6*	0.0876	2710 ± 780	42.7	0.17 ± 0.04	${1.78}_{-0.07}^{+0.07}$	42.7 ± 58.4	7.56	0.21	−17.9 ± 6.0	43.1
J221918.53+120753.1	0.0815	1110 ± 80	42.4	0.19 ± 0.01	${2.47}_{-0.13}^{+0.12}$	−26.3 ± 32.6	6.80	0.70	−63.7 ± 27.2	43.1

Notes. The object that shows a Seyfert-1.5-like spectrum is marked by a star. The redshifts given in column (2) are provided by the SDSS pipelines.

^aThe values of ${\rm{\Delta }}\upsilon$ are not available for a few objects due to the bad constraint on their Hβ narrow peaks. ^bThe value is given for Hβ rather than Hα because of the poorly observed Hα line profile. See Section 4.3.1 in the text for the details.

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[O iii]$\lambda 5007$ bulk velocity shift. Column (7) tabulates the calculated [O iii]$\lambda 5007$ line bulk relative velocity shift defined as ${\rm{\Delta }}\upsilon =c{\rm{\Delta }}\lambda /{\lambda }_{0,[{\rm{O}}{\rm{III}}]}$, where ${\lambda }_{0,[{\rm{O}}{\rm{III}}]}$ and ${\rm{\Delta }}\lambda$ denote the rest-frame wavelength in vacuum of the [O iii]λ5007 emission line and the wavelength shift with respect to the narrow Hβ line, respectively. ${\rm{\Delta }}\lambda$ is calculated from the modeled line centers as ${\rm{\Delta }}\lambda =({\lambda }_{[{\rm{O}}\;{\rm{III}}]}^{\mathrm{ob}}-{\lambda }_{{\rm{H}}\beta }^{\mathrm{ob}})-({\lambda }_{0,[{\rm{O}}{\rm{III}}]}-{\lambda }_{0,{\rm{H}}\beta })$, where ${\lambda }_{[{\rm{O}}\;{\rm{III}}]}^{\mathrm{ob}}$ (${\lambda }_{{\rm{H}}\beta }^{\mathrm{ob}}$) and ${\lambda }_{0,[{\rm{O}}{\rm{III}}]}$ (${\lambda }_{0,{\rm{H}}\beta }$) are the observed line center and the line wavelength in vacuum of the [O iii] (Hβ) line, respectively. The narrow Hβ line shows a very small velocity shift relative to the galaxy rest frame (e.g., Komossa et al. 2008), though this point has been argued against in recent studies⁷ (e.g., Hu et al. 2008; Bae & Woo 2014; Wang & Xu 2015). A negative value of ${\rm{\Delta }}\upsilon$ corresponds to a blueshift, and a positive value corresponds to a redshift.

[O iii]$\lambda 5007$ line asymmetry. Various "asymmetry" indices are commonly used in previous studies to quantify the asymmetry of the [O iii] emission line (e.g., Heckman et al. 1981; Whittle 1985; Veilleux 1991; Wang et al. 2011; Liu et al. 2013; Harrison et al. 2014). Briefly speaking, most of these indices (e.g., AI₂₀ and AI) give a quantified asymmetry by comparing the measured line widths/centers (in units of either wavelength or velocity) at different line flux levels. By modeling the [O iii]λ5007 emission-line profile by a sum of several Gaussian profiles, we parametrize the asymmetry of the [O iii] lines by a velocity of $\delta \upsilon$, defined as

$$\begin{eqnarray}&&\delta \upsilon =\displaystyle \frac{{\displaystyle \sum }_{k=1}^{n}{a}_{k}({\upsilon }_{k}-{\upsilon }_{p})}{{\displaystyle \sum }_{k=1}^{n}{a}_{k}}\end{eqnarray} \tag{ 2 }$$

where a_k and ${\upsilon }_{k}$ is the modeled flux and velocity of the kth Gaussian function, respectively. ${\upsilon }_{p}$ denotes the velocity of the Gaussian profile that reproduces the peak of the observed line profile. A negative value of $\delta \upsilon$ denotes a blue asymmetry, and a positive one denotes a red asymmetry. The fitted parameters $\delta \upsilon$ are listed in column (10) in Table 1. Figure 6 compares the value of $\delta \upsilon$ used in this study with the parameter of velocity offset $\delta {\upsilon }_{{\rm{H}}14}$ defined in Harrison et al. (2014). We calculate the values of $\delta \upsilon$_H14 for all of the 47 X-ray selected type I AGNs by following the definition in Harrison et al. (2014): $\delta {\upsilon }_{{\rm{H}}14}=({\upsilon }_{05}+{\upsilon }_{95})/2$, where ${\upsilon }_{05}$ and ${\upsilon }_{95}$ are the velocities at the 5% and 95% percentiles of the overall emission line, respectively. Although there is a small systematical difference, one can see from the figure a good correlation between the two parameters.

Uncertainty estimation. Except for ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$, all of the quoted errors in Table 1 correspond to a 1σ significance level. A proper error propagation is considered in the derivation of the uncertainties for some parameters. Specifically speaking, the uncertainty of RFe is estimated as

$$\begin{eqnarray}&&{\rm{\Delta }}\mathrm{RFe}=\mathrm{RFe}\sqrt{{\left(\displaystyle \frac{{\rm{\Delta }}{f}_{\mathrm{Fe}{\rm{II}}}}{{f}_{\mathrm{Fe}{\rm{II}}}}\right)}^{2}+{\left(\displaystyle \frac{{\rm{\Delta }}{f}_{{\rm{H}}\beta }}{{f}_{{\rm{H}}\beta }}\right)}^{2}},\end{eqnarray} \tag{ 3 }$$

where ${f}_{\mathrm{Fe}{\rm{II}}({\rm{H}}\beta )}$ and ${\rm{\Delta }}{f}_{\mathrm{Fe}{\rm{II}}({\rm{H}}\beta )}$ are the measured flux of the optical Fe ii complex (broad Hβ emission) and the corresponding uncertainty, respectively. We estimate the final uncertainty of ${\rm{\Delta }}\upsilon$ as

$$\begin{eqnarray}&&\delta {\rm{\Delta }}\upsilon =c\sqrt{{\left(\displaystyle \frac{{\rm{\Delta }}{\lambda }_{[{\rm{O}}{\rm{III}}]}}{{\lambda }_{[{\rm{O}}{\rm{III}}]}^{0}}\right)}^{2}+{\left(\displaystyle \frac{{\rm{\Delta }}{\lambda }_{{\rm{H}}\beta }}{{\lambda }_{{\rm{H}}\beta }^{0}}\right)}^{2}},\end{eqnarray} \tag{ 4 }$$

where ${\rm{\Delta }}{\lambda }_{[{\rm{O}}{\rm{III}}]({\rm{H}}\beta )}$ and ${\lambda }_{[{\rm{O}}\;{\rm{III}}]({\rm{H}}\beta )}^{0}$ are the uncertainty of the measured line center and the rest frame wavelength in vacuum of the [O iii] (narrow Hβ) emission line, respectively, and c is the light speed. The uncertainty of $\delta \upsilon$ is determined through the formula

$$\begin{eqnarray}&&\mathrm{std}.\ \delta \upsilon =\sqrt{\displaystyle \sum _{k=1}^{n}({{\rm{\Delta }}}^{2}{a}_{k}^{\prime }+{{\rm{\Delta }}}^{2}{\upsilon }_{k}^{\prime })}(k\;\ne \;p),\end{eqnarray} \tag{ 5 }$$

where ${\rm{\Delta }}{a}_{k}^{\prime }$ is the uncertainty of the weight ($={a}_{k}/\sum {a}_{k}$) of the kth Gaussian function and ${\rm{\Delta }}{\upsilon }_{k}^{\prime }=\sqrt{{{\rm{\Delta }}}^{2}{\upsilon }_{k}+{{\rm{\Delta }}}^{2}{\upsilon }_{p}}$ ($k\;\ne \;p$). ${\rm{\Delta }}{\upsilon }_{k}$ and ${\rm{\Delta }}{\upsilon }_{p}$ are the uncertainty of the velocity of the kth Gaussian function and that of the Gaussian function that reproduces the line peak.

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4.1.  ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ versus the [O iii]λ5007 Line Profile

The aim of this paper is to study the relationship between the properties of AGN X-ray emission and the strength of the feedback traced by the outflow in AGN NLR. The main results are presented in Figure 7 by using the [O iii]$\lambda 5007$ emission-line profile as a diagnose for the outflow occurring in the NLR. The related statistical results are listed in Table 2.

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Table 2.  Correlations Coefficient Matrix Related with the Hard-X-Ray Photon Index vs. [O iii]λ5007 Emission Line Profile Correlations

Property	${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$	$\delta \upsilon$	${\rm{\Delta }}\upsilon$
(1)	(2)	(3)	(4)
(1) ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$	⋯	−0.411(0.0054)	−0.483(0.0022)
(2) RFe	0.408(0.0057)	−0.268(0.0694)	−0.508(0.0011)
(3) ${\mathrm{FWHM}}_{{\rm{H}}\beta }$	−0.673($\lt {10}^{-4}$)	0.269(0.0685)	0.416(0.0077)
(4) ${M}_{\mathrm{BH}}$	−0.497(0.0008)	−0.073(0.6218)	0.509(0.0011)
(5) $L/{L}_{\mathrm{Edd}}$	0.620(< 10−4)	−0.418(0.0046)	−0.266(0.0886)
(6) ${L}_{1.4\mathrm{GHz}}$ a	−0.240(0.1510)	−0.170(0.3103)	0.281(0.0929)
(7) ${L}_{1.4\mathrm{GHz}}$ b	0.010(0.8882)	−0.107(0.1362)	0.085(0.2371)

Notes.

^aThe statistics are based on the 20 sources with detected radio flux taken from the FIRST catalog. ^bThe statistics are based on all 47 sources through survival analysis.

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The measured hard-X-ray photon index ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ is plotted against the resulting [O iii] line profile asymmetry parameter $\delta \upsilon$ in the left panel of Figure 7. One can see that there is a moderate anti-correlation between the two variables. The relationship means that the stronger the blue asymmetry of the [O iii]$\lambda 5007$ emission line, the steeper the X-ray spectrum will be. A Spearman rank-order test returns a correlation coefficient of ${r}_{s}=-0.411$. The corresponding probability of null correlation from two-tailed is calculated to be p_s = 0.0054, which corresponds to a significant level at $2.78\sigma$.

The right panel of Figure 7 shows a similar relationship between the X-ray photon index ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ and bulk velocity shift ${\rm{\Delta }}\upsilon$ of the [O iii]$\lambda 5007$ line: the larger the bulk blueshift of the [O iii] line, the steeper the X-ray spectrum will be. The two points marked with horizontal arrows at the left side of the plot are two objects (SDSS J092247.02+512038.0 and SDSS J140621.89+222346.5) with remarkable [O iii] line bulk blueshifts, i.e., ${\rm{\Delta }}\upsilon \leqslant -300\ \mathrm{km}\ {{\rm{s}}}^{-1}$. These values allow us to classify the two objects as "blue outliers" in which the [O iii] bulk blueshift is defined to be larger than $250\;\mathrm{km}\ {{\rm{s}}}^{-1}$ (e.g., Zamanov et al. 2002; Komossa et al. 2008). The measured ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ of the two objects are as large as 2.15 ± 0.09 and 2.23 ± 0.10. The Spearman rank-order test yields a correlation coefficient of ${r}_{s}=-0.483$ with a probability of null correlation of p_s = 0.0022 (i.e., a significance level at $3.06\sigma$). The statistics is slightly degraded to ${r}_{s}=-0.426$ with p_s = 0.0087 (i.e., a significance level at $2.62\sigma$) when the two objects with extremely large bulk blue velocity shifts are excluded.

4.2. Relation with Eigenvector-I Space

We examine the relation between the identified ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ versus the [O iii] line profile correlations and the well documented AGN Eigenvector-I (EI) space in this section. The EI space is one of the key properties of AGN phenomena. It was first introduced by Boroson & Green (1992; hereafter BG92) for a sample of 87 bright Palomar-Green quasars. In addition to the original anti-correlation between the intensities of the optical Fe ii blends and [O iii] emission, the space has been subsequently extended to infrared, UV and soft-X-ray bands (e.g., Wang et al. 1996; Laor et al. 1997; Lawrence et al. 1997; Grupe et al. 1999; Sulentic et al. 2000, 2002, 2004; Vaughan et al. 2001; Zamanov et al. 2002; Marziani et al. 2003; Xu et al. 2003, 2012; Grupe 2004; Wang et al. 2006; Marziani & Sulentic 2012). Currently, the best EI space that is widely accepted involves three parameters: the RFe, the FWHM of the Hβ broad component, and the photon index in soft-X-ray, which means that the relation with EI space can be studied equivalently by focusing on the two parameters: the RFe and FWHM(Hβ) in the current sample. The measured RFe in this paper is in fact found to be correlated with the FWHM(Hβ). With the Spearman rank-order test, the correlation coefficient and corresponding probability of null correlation are calculated to be ${r}_{s}=-0.397$ and ${\rho }_{s}=0.0071$ (2.69σ), respectively.

Figure 8 illustrates not only the RFe versus FWHM(Hβ) correlation through the symbol size, but also the dependence of the two identified correlations on both parameters, RFe and FWHM(Hβ). The corresponding correlation coefficient matrix is listed in Table 2. Each correlation coefficient and the corresponding probability of the null correlations shown in brackets are calculated through the Spearman rank-order test. One can learn from the table that there is a significant relationship between ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ and the EI space, which implies that the EI space can be well reproduced in the current X-ray-selected type I AGN sample, though the EI space is found to be marginally (moderately) correlated with the line asymmetry index $\delta \upsilon$ (bulk velocity shift ${\rm{\Delta }}\upsilon$). A significant anti-correlation between hard-X-ray photon index Γ and FWHM of broad Hβ emission line has been firmly established in previous studies (e.g., Brandt et al. 1997; Leighly 1999; Reeves & Turner 2000; Shemmer et al. 2006, 2008; Zhou & Zhang 2010; Jin et al. 2012).

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4.3. Role of SMBH Mass and Eddington Ratio

Both SMBH mass (${M}_{\mathrm{BH}}$) and $L/{L}_{\mathrm{Edd}}$ are critical parameters describing AGN phenomena (e.g., Shen & Ho 2014). In fact, BG92 first argued that the EI space is potentially driven by $L/{L}_{\mathrm{Edd}}$, which was then confirmed by various authors (e.g., Sulentic et al. 2000; Boroson 2002; Marziani et al. 2003; Xu et al. 2003, 2012) since the great progress made in the reverberation mapping technique (e.g., Kaspi et al. 2000, 2005; Peterson & Bentz 2006; and see Marziani & Sulentic 2012 and Peterson 2014 for recent reviews). To investigate the role played by the two basic parameters in the AGN feedback process and to explore the physical origin of the two new identified correlations, we first estimate the ${M}_{\mathrm{BH}}$ and $L/{L}_{\mathrm{Edd}}$ for the used hard-X-ray selected type-I AGNs from their optical spectra, and then study the statistical properties of the estimated ${M}_{\mathrm{BH}}$ and $L/{L}_{\mathrm{Edd}}$.

4.3.1. Derivation of ${M}_{\mathrm{BH}}$ and $L/{L}_{\mathrm{Edd}}$

The calculated ${M}_{\mathrm{BH}}$ and $L/{L}_{\mathrm{Edd}}$ are shown in columns (8) and (9) in Table 2 for each of the X-ray selected type-I AGNs, respectively. We estimate ${M}_{\mathrm{BH}}$ from the modeled broad Hα line emission for all the 47 X-ray selected type-I AGNs, except for SDSS J015950.24+002340.8, according to the calibrated relationship provided in Greene & Ho (2007), and references therein)

$$\begin{eqnarray}{M}_{\mathrm{BH}} & = & ({3.0}_{-0.5}^{+0.6})\times {10}^{6}{\left(\displaystyle \frac{{L}_{{\rm{H}}\alpha }}{{10}^{42}\mathrm{erg}{{\rm{s}}}^{-1}}\right)}^{0.45\pm 0.03}\\ & & \times {\left[\displaystyle \frac{\mathrm{FWHM}({\rm{H}}\alpha )}{{10}^{3}\mathrm{km}{{\rm{s}}}^{-1}}\right]}^{2.06\pm 0.06}{M}_{\odot }\end{eqnarray} \tag{ 6 }$$

where ${L}_{{\rm{H}}\alpha }$ is the intrinsic luminosity of the Hα broad component corrected for local extinction and $\mathrm{FWHM}({\rm{H}}\alpha )$ is the line width of broad Hα emission that resulted from our line profile modeling (Section 3.2.2). With the estimate ${M}_{\mathrm{BH}}$ (i.e., the Eddington luminosity), the Eddington Ratio ${L}_{\mathrm{bol}}/{L}_{\mathrm{Edd}}$ is inferred from a combination of the ${L}_{5100\mathring{\rm A} }$–${L}_{{\rm{H}}\alpha }$ relation⁸ (Greene & Ho 2005)

$$\begin{eqnarray}&&{L}_{5100\mathring{\rm A} }=2.4\times {10}^{43}{\left(\displaystyle \frac{{L}_{{\rm{H}}\alpha }}{{10}^{42}\mathrm{erg}{{\rm{s}}}^{-1}}\right)}^{0.86}\ \mathrm{erg}\;{{\rm{s}}}^{-1}\end{eqnarray} \tag{ 7 }$$

and the bolometric correction of ${L}_{\mathrm{bol}}=9\lambda {L}_{\lambda }(5100\;\mathring{\rm A} )$ (Kaspi et al. 2000).

Because of its poorly observed Hα line profile, the parameters ${M}_{\mathrm{BH}}$ and ${L}_{\mathrm{bol}}/{L}_{\mathrm{Edd}}$ of SDSS J015950.24+002340.8 are estimated from its broad Hβ emission based on a combination of the calibrations of (Greene & Ho 2005; Vestergaard & Peterson 2006):

$$\begin{eqnarray}&&{M}_{\mathrm{BH}}={10}^{6.67}{\left(\displaystyle \frac{{L}_{{\rm{H}}\beta }}{{10}^{42}\mathrm{erg}{{\rm{s}}}^{-1}}\right)}^{0.63}{\left[\displaystyle \frac{\mathrm{FWHM}({\rm{H}}\beta )}{{10}^{3}\mathrm{km}{{\rm{s}}}^{-1}}\right]}^{2}{M}_{\odot }\end{eqnarray} \tag{ 8 }$$

and

$$\begin{eqnarray}&&{L}_{5100\mathring{\rm A} }=7.31\times {10}^{43}{\left(\displaystyle \frac{{L}_{{\rm{H}}\beta }}{{10}^{42}\mathrm{erg}{{\rm{s}}}^{-1}}\right)}^{0.883}\ \mathrm{erg}\;{{\rm{s}}}^{-1}\end{eqnarray} \tag{ 9 }$$

where the luminosity of its broad Hβ emission ${L}_{{\rm{H}}\beta }$ is corrected for the local extinction estimated from a combination of a Balmer decrement for the standard case B recombination and a Galactic extinction curve with R_V = 3.1. Its bolometric luminosity is again obtained from ${L}_{5100\mathring{\rm A} }$ through the bolometric correction of ${L}_{\mathrm{bol}}=9\lambda {L}_{\lambda }(5100\;\mathring{\rm A} )$ (Kaspi et al. 2000).

4.3.2. Statistics

With the estimated ${M}_{\mathrm{BH}}$ and $L/{L}_{\mathrm{Edd}}$, Figure 9 illustrates the role of the two parameters on the identified ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ versus [O iii]λ5007 line profile correlations. The related correlation coefficient matrix, which is again based on the Spearman rank-order test, is listed in Table 2. On the one hand, the statistics show that both ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ and $\delta \upsilon$ are strongly correlated with the estimated $L/{L}_{\mathrm{Edd}}$. The Spearman rank-order tests indicate that the null probabilities of both correlations are smaller than 0.05, which corresponds to a significance level larger than 2σ. Specifically speaking, the significance levels are estimated to be $\gt 3.89\sigma$ and 2.89σ for the $L/{L}_{\mathrm{Edd}}$–${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ and $L/{L}_{\mathrm{Edd}}$–$\delta \upsilon$ correlations, respectively. For the ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ versus $\delta \upsilon$ correlation, AGNs associated with high $L/{L}_{\mathrm{Edd}}$ tend to occupy the soft X-ray spectrum end with strong [O iii] blue asymmetry, and ones with low $L/{L}_{\mathrm{Edd}}$ the hard-X-ray spectrum end with weak [O iii] blue asymmetry. This tendency implies that $L/{L}_{\mathrm{Edd}}$ is a potential physical driver of the ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ versus $\delta \upsilon$ correlation. On the other hand, in the current sample, ${\rm{\Delta }}\upsilon$ is found to be much better correlated with ${M}_{\mathrm{BH}}$ than $L/{L}_{\mathrm{Edd}}$.

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In this paper, we study the origin of AGN outflow in its NLR by focusing on the relationship between the [O iii]λ5007 line profile and hard-X-ray emission from the central SMBH. A joint spectral analysis in both optics and hard-X-ray allows us to reveal a moderate correlation between the hard-X-ray spectral photon index and [O iii] line asymmetry in a sample of 47 local ($z\lt 0.2$) hard-X-ray selected type-I AGNs at a significance level of $2.78\sigma$. It is noted that the results and implications presented here are only relevant for the AGNs that are most luminous in the local universe.

5.1.  ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}\mbox{--}\delta \upsilon$ Correlation: A Connection between the Accretion Disk and Outflow in NLR

5.1.1.  ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$: an Assessment of SMBH Accretion

We argue that the identified hard-X-ray spectral photon index versus the [O iii]$\lambda 5007$ line asymmetry correlation (i.e., the ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}\mbox{--}\delta \upsilon$ correlation) provides moderate evidence that the commonly observed AGN outflow in its NLR (at a radial distance of the order of 0.1–1 kpc from the central SMBH, e.g., Osterbrock & Ferland 2006; Heckman & Best 2014) is related with the accretion process occurring around the central SMBH (i.e., at a distance scale of $\sim {10}^{1-2}{R}_{s}$ from the central SMBH, where ${R}_{s}=2{{GM}}_{\mathrm{BH}}/{c}^{2}$ is the Schwarzschild radius).

The parameter ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ is believed to be closely linked with the accretion process around central SMBH. It is generally believed that the hard-X-ray emission of AGNs is produced in the region very close to the SMBH. A commonly accepted scenario of the hard-X-ray emission is the accretion disk-corona model in which a fraction of soft photons from the cold accretion disk is transformed to the hard-X-ray band through the inverse Compton scattering of the hot electrons with a temperature of ∼10⁹ K. These electrons are likely accelerated in the corona above the disk by the reconnection of the magnetic fields (e.g., Haardt & Maraschi 1991, 1993; Svensson & Zdziarski 1994; Kawaguchi et al. 2001; Liu et al. 2002, 2003; Cao 2009). This model can successfully explain the observed $L/{L}_{\mathrm{Edd}}\mbox{--}{\rm{\Gamma }}$ dependence (e.g., Lu & Yu 1999; Gierlinski & Done 2004; Grupe 2004; Porquet et al. 2004; Wang et al. 2004; Bian et al. 2005; Shemmer et al. 2006, 2008; Desroches et al. 2009; Risaliti et al. 2009; Zhou & Zhang 2010; Jin et al. 2012) as follows. The scattering can efficiently cool the corona above the accretion disk when the disk flux irradiating the corona increases, which finally results in a soft, steep X-ray spectrum at a high $L/{L}_{\mathrm{Edd}}$ state (e.g., Pounds et al. 1995).

5.1.2.  $L/{L}_{\mathrm{Edd}}$ as a Physical Driver

We further argue that $L/{L}_{\mathrm{Edd}}$ is a potential physical driver of the ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}\mbox{--}\delta \upsilon$ correlation. In the current X-ray type-I AGN sample, the dependence of ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ on $L/{L}_{\mathrm{Edd}}$ can be indirectly learned from the fact that ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ is found to be correlated with both RFe and FWHM of Hβ (see Table 2), which are basic parameters defining the EI space that is widely believed to be physically driven by $L/{L}_{\mathrm{Edd}}$. The line strength from BLR is determined by the vertical structure of the accretion disk, governed by $L/{L}_{\mathrm{Edd}}$, in a way in which a large $L/{L}_{\mathrm{Edd}}$ results in a large X-ray-heated volume that generates strong Fe ii complex emission (e.g., BG92). A direct relationship between ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$ and $L/{L}_{\mathrm{Edd}}$ can be further identified in the current sample from Table 2 and the bottom panel in Figure 10.

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If the above discussion on the physical driver of ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}\mbox{--}\delta \upsilon$ correlation is correct, a dependence of $\delta \upsilon$ on $L/{L}_{\mathrm{Edd}}$ (or EI space) is expected in the current sample. As shown in Table 2 and the top panel of Figure 10, this dependence at a significance level of 2.89σ can be virtually identified in the current sample. The dependence is not hard to be understood because it is generally believed that high $L/{L}_{\mathrm{Edd}}$ drives accretion disk winds (e.g., Proga & Kallman 2002). In fact, a correlation between [O iii] line profile asymmetry and $L/{L}_{\mathrm{Edd}}$ has been frequently reported in previous studies (e.g., Bian et al. 2005; Boroson 2005; Wang et al. 2011; Zhang et al. 2011; Wang 2015). A deficit of the extended emission-line region in the AGNs with high $L/{L}_{\mathrm{Edd}}$ is revealed by Matsuoka (2012). The deficit could be explained either by the AGN's outflow that blows the gas around central SMBH away or by galaxy minor merger, which produces radio-loud AGNs that are usually associated with an inefficient accretion.

5.1.3. Disk Wind Scenario

So far, two major types of feedback have been proposed for AGNs (Fabian 2012). One is known as the wind/radiation (quasar) mode (e.g., Murray et al. 1995; Proga et al. 2000; Crenshaw et al. 2003; King & Pounds 2003; Pounds et al. 2003; King 2003, 2005; Ganguly et al. 2007; Reeves et al. 2009; Alexander et al. 2010; Dunn et al. 2010 King et al. 2011; Zubovas & King 2012), and the other is the kinetic (radio) mode (e.g., Morganti et al. 2005, 2007; Rosario et al. 2010; Holt et al. 2011; Mahony et al. 2013). As recently reviewed in Fabian (2012), there is a current consensus that the two types of feedback occur in different AGN types. The wind/radiation mode dominantly operates in the luminous quasar phase, while the kinetic mode operates in the less luminous AGNs with low $L/{L}_{\mathrm{Edd}}$.

The linkage between the SMBH accretion disk and outflow in NLR that is diagnosed by the [O iii] line blue wing suggests that the observed feedback in NLR rationally originated from the disk winds.⁹ A likely scenario is that the blue wing of the [O iii] line is likely produced in the inner NLR region (e.g., Bian et al. 2005; Wang et al. 2005) in which the kinematics of the emitting gas is dominated by an acceleration caused by the wind/radiation pressure. The disk wind model has been successfully applied to explain the observed broad ultraviolet absorption lines in a fraction of $\sim 20\%$ quasars and the ultra-fast outflows identified from the blueshifted X-ray Fe xxv and Fe xxvi absorptions in a few local AGNs (e.g., Tombesi et al. 2012; Higginbottom et al. 2014). In the wind/radiation mode, a wind can be launched from the inner accretion disk where the ultraviolet photons are emitted from (e.g., Murray et al. 1995). The hydrodynamic outflow model calculated by Proga et al. (2008) indicates that the wind launched from the accretion disk can extended into the inner NLR, though the specific launch mechanism is still under debate. Possible mechanisms include radiation/line-driven (e.g., Proga et al. 1998, 2000; Laor & Brandt 2002; Proga & Kallman 2004; Nomura et al. 2013; Higginbottom et al. 2014; Hagino et al. 2015), thermally driven (e.g., Begelman et al. 1983; Krolik & Kriss 2001), magnetically driven (e.g., Blandford & Payne 1982; Ferreira 1997; Fukumura et al. 2014; Stepanovs & Fendt 2014), and hybrid models (e.g., Proga 2003; Everett 2005).

5.2. Radio Emission and [O iii] Line Profile

Some previous studies argued that the outflow in NLR is potentially driven by the interaction between the radio jet and the interstellar medium (e.g., Heckman et al. 1984; Whittle 1985, Brotherton 1996; Whittle & Wilson 2004; Morganti et al. 2007; Holt et al. 2008, 2011; Nesvadba et al. 2008; Guillard et al. 2012; Mahony et al. 2013). This argument was recently reinforced by the studies in Mullaney et al. (2013) and Zakamska & Greene (2014). Both studies claim a relation between [O iii] line asymmetry and radio luminosity in both type I and II AGNs observed by SDSS.

The role of radio emission played in the [O iii] line profile is a test for the 47 X-ray selected type I AGNs used in this study. The 47 X-ray AGNs are cross-matched with the FIRST survey catalog (Becker et al. 2003). This environmental cross-match returns only 20 radio objects with detected radio flux exceeding the FIRST limiting flux density (5σ) of 1mJy. The upper limit of radio flux is also taken from the FIRST survey for the each of the other 27 sources, basing upon the reported detection limit at the corresponding celestial position. The luminosity at 1.4 GHz (rest frame) of each radio source at a given redshift z is calculated from the observed integrated flux density at 1.4 GHz f_ν through ${L}_{1.4\mathrm{GHz}}=4\pi {d}_{L}^{2}{f}_{\nu }{(1+z)}^{-1-\alpha }$, where d_L is the luminosity distance, and $\alpha =-0.8$ (e.g., Ker et al. 2012) is the spectral slope defined as ${f}_{\nu }\propto {\nu }^{\alpha }$. The calculated ${L}_{1.4\mathrm{GHz}}$ ranges from 10²² to ${10}^{26}\ {\rm{W}}\ {\mathrm{Hz}}^{-1}$, and is plotted against ${{\rm{\Gamma }}}_{2-10\mathrm{keV}}$, $\delta \upsilon$, and ${\rm{\Delta }}\upsilon$ in the left, middle, and right panels in Figure 11, respectively. The solid blue points denote the sources with a detected radio flux, and the open red points the ones with a flux upper limit. The corresponding statistics based on Kendall's τ is tabulated in Table 2. The values tabulated in line (6) are based on the 20 sources with a detected radio flux, and the ones in line (7) on all the 47 sources through the survival analysis with non-parametric model (e.g., Isobe et al. 1986). Our statistics shows that in the current sample there is no evidence that both blue wing and bulk velocity blueshift of the [O iii] line are driven by the radio emission, i.e., no direct relationship can be identified between ${L}_{1.4\mathrm{GHz}}$ and $\delta \upsilon$ (${\rm{\Delta }}\upsilon$) in the current sample.

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We argue that this result does not mean an inevitable disagreement on the previous studies. First, one should bear in mind that the relationships related to radio luminosity are difficult to firmly test in the current study because its sample size (47 sources) is significantly less than those (with hundreds to thousands of sources) in Mullaney et al. (2013) and Zakamska & Greene (2014). Second, the middle panel in Figure 11 shows that the most significant [O iii] blue asymmetry tends to occur in the objects with a radio luminosity at 1.4 GHz of 10²²–10²⁴ W Hz⁻¹, which is close to that observed in Mullaney et al. (2013). Zakamska & Greene (2014) recently proposed that the radio emission in radio-quite AGNs might be produced by the accelerated particles in the interstellar medium of the host galaxy that is shocked by the accretion disk wind.

5.3. Evolution of Feedback in AGNs

We study the origin of AGN outflow occurring in its NLR by focusing on the relationship between the [O iii]λ5007 line profile and hard-X-ray emission from the central SMBH in a sample of 47 local X-ray selected type I AGNs ($z\lt 0.2$). These luminous AGNs are extracted from the 2XMMi/SDSS-DR7 catalog, and have X-ray luminosities of 2–10 keV in a range from 10⁴² to ${10}^{45}\ \mathrm{erg}\ {{\rm{s}}}^{-1}$. A joint spectral analysis in both optics and hard X-ray on the sample allows us to identify a moderate correlation with a significance level of 2.78σ, in which luminous AGNs with more significant [O iii] blue asymmetry tend to be associated with steeper X-ray spectra. Our statistics show that the correlation is related to $L/{L}_{\mathrm{Edd}}$ at a 2–3σ significance level, which suggests that the AGN outflow in its NLR is likely driven by the accretion process occurring around the central SMBH.

The authors thank the anonymous referee for his/her careful review and helpful suggestions for improving the manuscript. We thank Dr. X. L. Zhou for help with the X-ray spectral analysis. This study uses the SDSS archive data that was created and distributed by the Alfred P. Sloan Foundation. This work is based on observations obtained with XMM-Newton, an ESA science mission with instruments and contributions directly funded by ESA Member States and the USA (NASA). The study is supported by the National Basic Research Program of China (grant 2009CB824800) and by the National Natural Science Foundation of China under grants 11473036 and 11273027.