[O iii] λ5007 AND X-RAY PROPERTIES OF A COMPLETE SAMPLE OF HARD X-RAY SELECTED AGNs IN THE LOCAL UNIVERSE

For our study, we utilize the Swift/BAT 9-month catalog (Tueller et al. 2008) to define a complete sample of hard X-ray selected AGNs in the local universe. The Tueller et al. (2008) catalog contains 137 AGNs in total, excluding blazars at a flux limit of 2 × 10⁻¹¹ erg cm⁻² s⁻¹ in the 14–195 keV band with detection significance above 4.8σ. To minimize the effects of extinction by Galactic interstellar medium, we limit the sample to those located at high Galactic latitudes of $| b| \gt 15^\circ$ for our optical spectral studies. We exclude Cen A, which is a very nearby object, and SWIFT J0350.1–5019, which likely is confused by two AGNs, PGC 13946 and ESO 201-IG 004 (C. Ricci et al. 2015, in preparation). These selections leave 103 AGNs that constitute our "parent" sample (hereafter "Sample A").

The Swift/BAT AGNs are extensively followed up by X-ray observatories covering below 10 keV, such as Swift/XRT, XMM-Newton, Suzaku, and Chandra. Key spectral parameters in our study are the absorption column density (${N}_{{\rm{H}}}$) and the fraction of scattered component (${f}_{{\rm{scat}}}$) for absorbed AGNs, which are often obtained by utilizing a partially covered absorber (or its equivalent) model. Because there can be other soft X-ray components that are spatially unresolved from the AGN emission, the ${f}_{{\rm{scat}}}$ value determined in this way is an upper limit to the true scattering fraction. Here we basically adopt the results of spectral analysis summarized in Table 1 of Ichikawa et al. (2012), which was largely based on Winter et al. (2009a) and was revised from (then) available Suzaku results for some targets. In our paper, we further revised their table by referring to later papers utilizing Suzaku data for more objects. Furthermore, for sources whose spectral parameters were not well constrained by using only the Swift/XRT data in Winter et al. (2009a), we update their spectral parameters according to C. Ricci et al. (2015, in preparation), who perform uniform broadband spectral analysis in the 0.3–150 keV band by including Swift/BAT spectra for the whole AGN sample of the Swift/BAT 70-month catalog. We also utilize the 70-month averaged, de-absorbed 2–10 keV flux of the primary continuum listed in the C. Ricci et al. (2015, in preparation) catalog, as well as the 9-month averaged 14–195 keV flux in the original Tueller et al. (2008) catalog.

We divide the sample into two types, X-ray unabsorbed AGNs (hereafter "X-ray type 1 AGNs") and absorbed AGNs ("X-ray type 2 AGNs"), which have absorptions of log ${N}_{{\rm{H}}}$ < 22 cm⁻² and log ${N}_{{\rm{H}}}$ ≥ 22, respectively. Among X-ray type 2 AGNs, we call those with ${f}_{{\rm{scat}}}$ < 0.5% the low scattering fraction AGNs (so-called new type AGNs), a putative population of AGNs deeply buried by geometrically thick tori. Owing to our revision of the X-ray spectral parameters in the original Swift/BAT 9-month catalog, the sample of low scattering fraction AGNs has also been updated¹⁴ from that originally defined in Ichikawa et al. (2012).

Table 1 lists the targets of Sample A with their basic X-ray properties: source number in Tueller et al. (2008) (** are attached to the low scattering fraction AGNs), source name, redshift, ${N}_{{\rm{H}}}$, ${f}_{{\rm{scat}}}$, observed luminosity in the 14–195 keV band (9-month average), absorption-corrected 2–10 keV luminosity (70-month average), and reference for the X-ray spectral parameters. Though not listed in Table 1, we also compile the information on the inclination angle of the host galaxy, ${i}_{{\rm{host}}}$, using the HyperLeda database, which is available for 98 AGNs.¹⁵ In addition, the black hole mass (and hence an estimate of Eddington ratio) is available for 99 AGNs¹⁶ from Winter et al. (2009a).

Table 1.  X-Ray and Optical Emission Line ([O iii] λ5007, Hα, Hβ) Properties of AGNs in the 9-Month Swift/BAT Catalog

No.	Object	z	$\mathrm{log}{L}_{14-195}$	$\mathrm{log}{L}_{2-10}$	$\mathrm{log}{N}_{{\rm{H}}}$	${f}_{{\rm{scat}}}$	$\mathrm{log}{L}_{{\rm{[O}}\;{\rm{III]}}}^{{\rm{cor}}}$	$\mathrm{log}{L}_{{\rm{[O}}\;{\rm{III]}}}$	$\mathrm{log}{F}_{{\rm{[O}}\;{\rm{III]}}}$	$\mathrm{log}{F}_{{\rm{H}}\alpha }$	$\mathrm{log}{F}_{{\rm{H}}\beta }$	References
			(erg s−1)	(erg s−1)	(cm−2)		(erg s−1)	(erg s−1)	(erg cm−2s−1)	(erg cm−2 s−1)	(erg cm−2 s−1)	(X-ray)	(Optical)
(1)	(2)	(3)	(4)	(5)	(6)	(7)	(8)	(9)	(10)	(11)	(12)	(13)	(14)
1 ${}^{**}$	NGC 235A	0.0222	43.56	43.22	23.50	0.003	42.17 ± 0.43	41.23 ± 0.10	−12.82 ± 0.10	−13.17 ± 0.10	−13.96 ± 0.10	(xa)	(oa)
2 ${}^{**}$	Mrk 348	0.0150	43.68	43.30	23.20	0.004	41.95 ± 0.01	41.07 ± 0.01	−12.64 ± 0.01	−13.02 ± 0.01	−13.80 ± 0.01	(xb)	(ob)
3	Mrk 352	0.0149	43.27	42.74	20.75	⋯	⋯	40.39 ± 0.01	−13.31 ± 0.01	⋯	⋯	(xc)	(oc)
4	NGC 454	0.0121	42.88	42.17	23.30	0.030	>40.87	40.41 ± 0.20	−13.11 ± 0.20	−13.13 ± 0.20	<−13.77	(xa)	(oa)
5	Fairall 9	0.0470	44.39	44.15	20.36	⋯	⋯	42.15 ± 0.20	−12.56 ± 0.20	⋯	⋯	(xc)	(oa)
6	NGC 526A	0.0191	43.63	43.07	22.18	⋯	41.44 ± 0.43	41.32 ± 0.10	−12.59 ± 0.10	−13.08 ± 0.10	−13.59 ± 0.10	(xd)	(oa)
7	NGC 612	0.0298	43.81	43.45	24.05	0.006	40.09 ± 0.92	40.09 ± 0.13	−14.22 ± 0.13	−13.96 ± 0.11	−14.41 ± 0.29	(xe)	(od)
8 ${}^{**}$	ESO 297-G018	0.0252	43.85	43.67	23.81	0.003	41.43 ± 1.02	40.84 ± 0.04	−13.33 ± 0.04	−13.44 ± 0.21	−14.12 ± 0.27	(xf)	(od)
9	NGC 788	0.0136	43.39	43.22	23.67	0.007	40.91 ± 0.49	40.91 ± 0.18	−12.71 ± 0.18	−13.29 ± 0.01	−13.03 ± 0.16	(xd)	(oe)
10	Mrk 1018	0.0424	44.17	43.61	0.00	⋯	41.71 ± 0.65	41.71 ± 0.09	−12.91 ± 0.09	−13.46 ± 0.20	−13.91 ± 0.09	(xa)	(oe)
12	Mrk 590	0.0264	43.77	42.71	20.43	⋯	42.07 ± 0.12	41.74 ± 0.04	−12.46 ± 0.04	−13.00 ± 0.01	−13.59 ± 0.04	(xc)	(oe)
15	NGC 931	0.0167	43.66	43.41	21.56	⋯	41.91 ± 0.03	41.14 ± 0.01	−12.66 ± 0.01	−13.07 ± 0.01	−13.81 ± 0.01	(xd)	(od)
16	NGC 985	0.0430	44.21	43.78	21.59	⋯	42.75 ± 1.99	41.92 ± 0.01	−12.72 ± 0.01	−13.00 ± 0.01	−13.75 ± 0.68	(xc)	(od)
17	ESO 416-G002	0.0592	44.42	43.57	20.43	⋯	⋯	41.63 ± 0.20	−13.29 ± 0.20	⋯	⋯	(xc)	(oa)
18	ESO 198-024	0.0455	44.27	43.50	21.00	⋯	⋯	41.22 ± 0.20	−13.46 ± 0.20	⋯	⋯	(xc)	(oa)
20 ${}^{**}$	NGC 1142	0.0289	44.17	43.88	23.80	0.003	42.03 ± 0.03	41.07 ± 0.01	−13.21 ± 0.01	−13.30 ± 0.01	−14.10 ± 0.01	(xf)	(od)
23	PKS 0326–288	0.1080	44.84	44.42	23.82	0.009	43.58 ± 0.85	42.52 ± 0.20	−12.95 ± 0.20	−13.15 ± 0.20	−13.98 ± 0.20	(xa)	(oa)
24	NGC 1365	0.0055	42.67	42.03	24.65	0.260	40.98 ± 0.01	39.62 ± 0.01	−13.21 ± 0.01	(8.7)	⋯	(xg)	(of)
25	ESO 548-G081	0.0145	43.19	42.91	20.00	⋯	41.94 ± 0.08	40.99 ± 0.01	−12.69 ± 0.01	−12.65 ± 0.01	−13.45 ± 0.03	(xd)	(od)
28	2MASX J03565655–4041453	0.0747	44.51	43.70	22.66	0.032	42.56 ± 0.43	42.18 ± 0.10	−12.95 ± 0.10	−13.29 ± 0.10	−13.90 ± 0.10	(xa)	(oa)
29 ${}^{**}$	3C 105	0.0890	44.83	44.32	23.75	0.003	42.79 ± 0.10	41.59 ± 0.01	−13.70 ± 0.01	−14.10 ± 0.01	−14.99 ± 0.03	(xa)	(oe)
31	1H 0419–577	0.1040	44.91	44.60	24.31	⋯	43.55 ± 0.43	43.16 ± 0.10	−12.28 ± 0.10	−12.64 ± 0.10	−13.26 ± 0.10	(xh)	(oa)
32	3C 120	0.0330	44.45	43.98	21.20	⋯	⋯	41.86 ± 0.01	−12.54 ± 0.01	⋯	⋯	(xd)	(og)
34	MCG –01-13-025	0.0159	43.41	42.54	19.60	⋯	40.95 ± 0.03	40.73 ± 0.01	−13.02 ± 0.01	−13.21 ± 0.01	−13.76 ± 0.01	(xc)	(oe)
36	XSS J05054–2348	0.0350	44.24	43.49	23.47	0.009	41.65 ± 0.43	41.42 ± 0.10	−13.03 ± 0.10	−13.15 ± 0.10	−13.71 ± 0.10	(xf)	(oa)
38	Ark 120	0.0323	44.11	43.79	20.30	⋯	⋯	41.35 ± 0.01	−13.03 ± 0.01	⋯	⋯	(xd)	(oc)
39	ESO 362-G018	0.0126	43.26	42.88	23.43	0.087	⋯	40.88 ± 0.20	−12.67 ± 0.20	⋯	⋯	(xd)	(oa)
40	PICTOR A	0.0351	43.80	43.45	20.78	⋯	⋯	41.57 ± 0.20	−12.89 ± 0.20	⋯	⋯	(xd)	(oa)
45	NGC 2110	0.0078	43.54	43.17	22.45	0.048	41.64 ± 0.01	40.36 ± 0.01	−12.77 ± 0.01	−12.66 ± 0.01	−13.57 ± 0.01	(xd)	(ob)
47	EXO 055620–3820.2	0.0339	44.14	43.07	22.41	0.034	⋯	41.46 ± 0.10	−12.97 ± 0.10	⋯	⋯	(xd)	(oa)
49 ${}^{**}$	ESO 005-G004	0.0062	42.56	41.93	24.06	0.003	⋯	<38.63	<−14.30	−13.80 ± 0.01	<−15.39	(xf)	(oh)
50	Mrk 3	0.0135	43.61	43.35	24.04	0.009	43.33 ± 0.01	42.31 ± 0.01	−11.31 ± 0.01	−11.65 ± 0.01	−12.47 ± 0.01	(xi)	(oe)
51 ${}^{**}$	ESO 121-IG028	0.0403	44.03	43.63	23.31	0.004	>40.86	40.86 ± 0.24	−13.72 ± 0.24	−13.69 ± 0.10	<−13.76	(xa)	(oa)
53	2MASX J06403799–4321211	0.0610	44.40	43.51	23.00	<0.011	>41.74	41.74 ± 0.20	−13.21 ± 0.20	−13.19 ± 0.20	<−13.33	(xa)	(oa)
55	Mrk 6	0.0188	43.72	43.09	20.76	⋯	42.38 ± 1.56	42.38 ± 0.01	−11.52 ± 0.01	−11.97 ± 0.01	−12.41 ± 0.53	(xa)	(oe)
56	Mrk 79	0.0222	43.72	43.19	19.78	⋯	41.96 ± 0.12	41.96 ± 0.02	−12.09 ± 0.02	−12.66 ± 0.01	−13.11 ± 0.04	(xc)	(oe)
60	Mrk 18	0.0111	42.93	41.82	23.26	0.030	40.56 ± 0.28	40.56 ± 0.11	−12.88 ± 0.11	−12.72 ± 0.01	−12.98 ± 0.09	(xd)	(oe)
61	2MASX J09043699+5536025	0.0370	44.03	43.31	20.78	⋯	41.91 ± 0.09	41.63 ± 0.03	−12.87 ± 0.03	−12.95 ± 0.01	−13.52 ± 0.03	(xd)	(oe)
62	2MASX J09112999+4528060	0.0268	43.69	43.16	23.52	0.006	40.98 ± 0.11	39.68 ± 0.01	−14.54 ± 0.01	−14.49 ± 0.01	−15.41 ± 0.04	(xd)	(oe)
64	2MASX J09180027+0425066	0.1560	45.31	⋯	23.05	0.013	42.59 ± 0.01	42.22 ± 0.01	−13.60 ± 0.01	−14.08 ± 0.01	−14.68 ± 0.01	(xd)	(oe)
65	MCG –01-24-012	0.0196	43.60	43.24	22.81	0.005	41.10 ± 0.15	41.10 ± 0.08	−12.83 ± 0.08	−13.14 ± 0.01	−13.46 ± 0.05	(xa)	(oe)
66	MCG +04-22-042	0.0323	43.99	43.46	20.59	⋯	42.12 ± 0.62	42.12 ± 0.20	−12.26 ± 0.20	−12.57 ± 0.01	−12.72 ± 0.20	(xc)	(oe)
67	Mrk 110	0.0353	44.19	43.86	20.20	⋯	42.71 ± 0.59	42.29 ± 0.19	−12.17 ± 0.19	−12.47 ± 0.01	−13.09 ± 0.19	(xd)	(oe)
68	NGC 2992	0.0077	42.94	41.93	22.08	0.524	42.51 ± 0.43	40.76 ± 0.10	−12.36 ± 0.10	−12.48 ± 0.10	−13.55 ± 0.10	(xd)	(oa)
69	MCG –05-23-016	0.0085	43.55	43.21	22.20	⋯	>41.18	40.64 ± 0.10	−12.56 ± 0.10	−12.85 ± 0.10	<−13.50	(xd)	(oa)
70	NGC 3081	0.0080	43.09	42.96	23.99	0.006	41.66 ± 0.43	41.32 ± 0.10	−11.83 ± 0.10	−12.38 ± 0.10	−12.97 ± 0.10	(xe)	(oa)
71	NGC 3227	0.0039	42.63	42.05	22.24	0.148	41.26 ± 0.01	40.44 ± 0.01	−12.09 ± 0.01	−12.54 ± 0.01	−13.29 ± 0.01	(xd)	(ob)
72	NGC 3281	0.0107	43.27	42.69	23.94	0.019	41.42 ± 0.85	41.05 ± 0.20	−12.36 ± 0.20	−12.71 ± 0.20	−13.31 ± 0.20	(xd)	(oa)
75 ${}^{**}$	Mrk 417	0.0328	43.95	43.73	23.93	0.002	41.19 ± 0.05	40.92 ± 0.01	−13.48 ± 0.01	−13.76 ± 0.01	−14.33 ± 0.02	(xd)	(oe)
77	NGC 3516	0.0088	43.26	42.72	21.55	⋯	41.54 ± 0.01	40.92 ± 0.01	−12.32 ± 0.01	(4.9)	⋯	(xd)	(oc)
78	RX J1127.2+1909	0.1055	44.79	43.84	0.00	⋯	43.02 ± 0.14	43.02 ± 0.11	−12.43 ± 0.11	−13.03 ± 0.01	−13.42 ± 0.03	(xa)	(oe)
79	NGC 3783	0.0097	43.53	43.30	21.76	0.278	⋯	41.47 ± 0.10	−11.85 ± 0.10	⋯	⋯	(xd)	(oa)
80	SBS 1136+594	0.0601	44.33	43.82	0.00	⋯	42.70 ± 0.06	42.55 ± 0.01	−12.39 ± 0.01	−12.92 ± 0.01	−13.45 ± 0.02	(xa)	(oe)
81	UGC 06728	0.0065	42.54	41.94	20.00	⋯	40.22 ± 0.01	40.22 ± 0.01	−12.75 ± 0.01	−12.34 ± 0.01	−12.80 ± 0.01	(xd)	(oe)
82	2MASX J11454045–1827149	0.0330	43.98	43.64	0.00	⋯	42.19 ± 0.43	42.19 ± 0.10	−12.21 ± 0.10	−12.91 ± 0.10	−13.36 ± 0.10	(xa)	(oa)
83	CGCG 041–020	0.0360	43.88	43.39	23.03	0.009	41.17 ± 0.07	40.38 ± 0.01	−14.10 ± 0.01	−13.96 ± 0.01	−14.70 ± 0.02	(xd)	(oe)
85	NGC 4051	0.0023	41.74	41.33	20.46	⋯	40.41 ± 0.56	40.20 ± 0.18	−11.87 ± 0.18	−11.97 ± 0.01	−12.52 ± 0.18	(xc)	(oe)
86	Ark 347	0.0224	43.42	42.90	23.36	0.016	41.53 ± 0.22	41.53 ± 0.19	−12.52 ± 0.19	−13.08 ± 0.01	−13.31 ± 0.04	(xa)	(oe)
87	NGC 4102	0.0028	41.62	41.41	24.30	⋯	40.72 ± 0.02	38.78 ± 0.01	−13.46 ± 0.01	−12.48 ± 0.01	−13.62 ± 0.01	(xj)	(od)
88	NGC 4138	0.0030	41.62	41.23	22.90	0.012	38.95 ± 0.01	38.95 ± 0.01	−13.35 ± 0.01	−13.33 ± 0.01	−13.72 ± 0.01	(xd)	(od)
89	NGC 4151	0.0033	42.96	42.58	22.73	0.041	41.87 ± 0.24	41.87 ± 0.01	−10.51 ± 0.01	−11.12 ± 0.01	−11.40 ± 0.08	(xd)	(oe)
90	Mrk 766	0.0129	42.94	42.67	21.72	⋯	42.03 ± 0.13	41.79 ± 0.01	−11.78 ± 0.01	−12.10 ± 0.04	−12.66 ± 0.02	(xc)	(oe)
91	NGC 4388	0.0084	43.60	43.17	23.53	0.011	41.29 ± 0.13	41.29 ± 0.11	−11.90 ± 0.11	−12.33 ± 0.01	−12.80 ± 0.03	(xk)	(oe)
92	NGC 4395	0.0011	40.81	40.64	22.52	0.322	39.02 ± 0.01	38.93 ± 0.01	−12.49 ± 0.01	−12.81 ± 0.01	−13.32 ± 0.01	(xd)	(oe)
94	NGC 4507	0.0118	43.78	43.54	23.54	0.029	42.12 ± 0.43	41.76 ± 0.10	−11.73 ± 0.10	−12.10 ± 0.10	−12.70 ± 0.10	(xd)	(oa)
95 ${}^{**}$	ESO 506-G027	0.0250	44.28	43.95	23.92	0.002	>41.14	40.96 ± 0.20	−13.19 ± 0.20	−13.67 ± 0.20	<−14.20	(xl)	(oa)
96	XSS J12389–1614	0.0366	44.26	43.35	22.63	0.043	41.88 ± 0.85	41.87 ± 0.20	−12.63 ± 0.20	−12.81 ± 0.20	−13.29 ± 0.20	(xa)	(oa)
97	NGC 4593	0.0090	43.21	42.81	20.49	⋯	41.04 ± 0.01	40.50 ± 0.01	−12.76 ± 0.01	(4.6)	⋯	(xd)	(oc)
100	SBS 1301+540	0.0299	43.72	43.09	20.60	⋯	41.25 ± 0.21	41.25 ± 0.01	−13.06 ± 0.01	−13.54 ± 0.07	−13.87 ± 0.01	(xc)	(oe)
102 ${}^{**}$	NGC 4992	0.0251	43.83	43.17	23.75	<0.003	39.85 ± 0.16	39.85 ± 0.01	−14.30 ± 0.01	−14.50 ± 0.02	−14.82 ± 0.05	(xm)	(od)
103	MCG –03-34-064	0.0165	43.46	43.95	23.61	0.039	42.55 ± 0.43	42.21 ± 0.10	−11.57 ± 0.10	−12.13 ± 0.10	−12.73 ± 0.10	(xd)	(oa)
105	MCG –06-30-015	0.0077	43.00	42.74	21.28	⋯	⋯	40.23 ± 0.20	−12.89 ± 0.20	⋯	⋯	(xd)	(oa)
106	NGC 5252	0.0230	43.90	43.39	22.64	0.038	40.79 ± 0.01	40.37 ± 0.01	−13.71 ± 0.01	−13.32 ± 0.01	−13.94 ± 0.01	(xd)	(oe)
108	IC 4329A	0.0160	44.24	43.84	21.79	⋯	41.62 ± 0.19	41.02 ± 0.04	−12.74 ± 0.04	−13.09 ± 0.04	−13.77 ± 0.04	(xd)	(od)
109	Mrk 279	0.0304	43.97	43.42	20.11	⋯	41.76 ± 0.21	41.68 ± 0.01	−12.65 ± 0.01	−13.01 ± 0.01	−13.51 ± 0.07	(xd)	(od)
110	NGC 5506	0.0062	43.30	42.91	22.44	0.011	41.89 ± 0.11	41.03 ± 0.08	−11.90 ± 0.08	−11.99 ± 0.01	−12.76 ± 0.03	(xd)	(oe)
112	NGC 5548	0.0172	43.59	43.14	20.85	⋯	42.13 ± 0.06	42.13 ± 0.02	−11.69 ± 0.02	−12.37 ± 0.01	−12.74 ± 0.02	(xd)	(oe)
113	ESO 511-G030	0.0224	43.73	43.41	20.99	⋯	⋯	40.62 ± 0.20	−13.44 ± 0.20	⋯	⋯	(xd)	(oa)
115	NGC 5728	0.0093	43.23	43.03	24.14	0.007	41.96 ± 0.43	41.52 ± 0.10	−11.76 ± 0.10	−12.23 ± 0.10	−12.85 ± 0.10	(xm)	(oa)
116	Mrk 841	0.0364	44.20	43.87	21.34	⋯	41.64 ± 0.04	41.64 ± 0.01	−12.85 ± 0.01	−13.37 ± 0.01	−13.79 ± 0.01	(xc)	(od)
117	Mrk 290	0.0296	43.79	42.93	21.18	⋯	41.71 ± 0.50	41.59 ± 0.01	−12.71 ± 0.01	−13.20 ± 0.01	−13.72 ± 0.17	(xd)	(od)
118	Mrk 1498	0.0547	44.50	44.05	23.10	0.016	42.43 ± 0.31	42.43 ± 0.20	−12.42 ± 0.20	−13.18 ± 0.01	−13.05 ± 0.08	(xf)	(oe)
120 ${}^{**}$	NGC 6240	0.0245	43.81	44.16	24.25	<0.005	42.12 ± 0.12	40.71 ± 0.02	−13.43 ± 0.02	−12.73 ± 0.01	−13.69 ± 0.04	(xa)	(oe)
124	1RXS J174538.1+290823	0.1113	45.09	44.37	0.00	⋯	⋯	42.75 ± 0.03	−12.75 ± 0.03	⋯	⋯	(xa)	(oe)
125	3C 382	0.0579	44.81	44.67	20.11	⋯	42.31 ± 0.08	41.70 ± 0.01	−13.20 ± 0.01	−13.36 ± 0.01	−14.04 ± 0.03	(xc)	(od)
126 ${}^{**}$	ESO 103-035	0.0133	43.58	43.38	23.33	0.001	42.20 ± 0.85	40.87 ± 0.20	−12.73 ± 0.20	−12.80 ± 0.20	−13.73 ± 0.20	(xd)	(oa)
127	3C 390.3	0.0561	44.88	44.52	21.08	⋯	42.96 ± 0.38	42.72 ± 0.01	−12.15 ± 0.01	−12.51 ± 0.13	−13.06 ± 0.02	(xd)	(od)
129	NGC 6814	0.0052	42.57	42.22	20.76	⋯	⋯	40.17 ± 0.10	−12.61 ± 0.10	⋯	⋯	(xc)	(oa)
133	NGC 6860	0.0149	43.39	42.89	21.00	⋯	⋯	40.93 ± 0.10	−12.76 ± 0.10	⋯	⋯	(xn)	(oa)
136	4C +74.26	0.1040	45.14	44.87	21.25	⋯	43.37 ± 0.61	43.13 ± 0.17	−12.31 ± 0.17	−12.83 ± 0.01	−13.39 ± 0.20	(xo)	(oe)
137	Mrk 509	0.0344	44.43	44.08	20.18	⋯	⋯	42.17 ± 0.20	−12.26 ± 0.20	⋯	⋯	(xd)	(oa)
138	IC 5063	0.0114	43.31	43.08	23.40	0.009	42.26 ± 0.43	41.58 ± 0.10	−11.88 ± 0.10	−12.20 ± 0.10	−12.91 ± 0.10	(xh)	(oa)
139	2MASX J21140128+8204483	0.0840	44.80	44.35	0.00	⋯	>43.60	42.89 ± 0.01	−12.35 ± 0.01	−12.54 ± 0.04	<−13.25	(xa)	(od)
144	UGC 11871	0.0266	43.80	43.26	22.32	0.016	42.56 ± 0.01	41.45 ± 0.01	−12.76 ± 0.01	−12.25 ± 0.01	−13.10 ± 0.01	(xa)	(oe)
145 ${}^{**}$	NGC 7172	0.0087	43.32	42.90	22.91	0.001	39.94 ± 0.85	39.94 ± 0.20	−13.28 ± 0.20	−13.55 ± 0.20	−13.75 ± 0.20	(xd)	(oa)
146	NGC 7213	0.0058	42.59	41.85	20.40	⋯	40.23 ± 0.85	40.23 ± 0.20	−12.64 ± 0.20	−12.39 ± 0.20	−12.83 ± 0.20	(xd)	(oa)
147	NGC 7314	0.0048	42.45	41.96	21.60	⋯	40.25 ± 0.43	39.78 ± 0.10	−12.93 ± 0.10	−13.21 ± 0.10	−13.85 ± 0.10	(xd)	(oa)
148 ${}^{**}$	NGC 7319	0.0225	43.68	43.40	23.82	0.004	40.72 ± 0.20	40.72 ± 0.03	−13.34 ± 0.03	−13.68 ± 0.01	−13.86 ± 0.07	(xa)	(oe)
149	3C 452	0.0811	44.73	43.83	23.36	0.064	42.09 ± 0.09	40.98 ± 0.01	−14.23 ± 0.01	−14.22 ± 0.01	−15.07 ± 0.03	(xd)	(oe)
151	MR 2251–178	0.0640	45.03	44.60	21.45	⋯	42.87 ± 0.43	42.33 ± 0.10	−12.66 ± 0.10	−12.89 ± 0.10	−13.55 ± 0.10	(xd)	(oa)
152	NGC 7469	0.0163	43.70	42.97	20.61	⋯	43.14 ± 0.01	41.70 ± 0.01	−12.08 ± 0.01	−11.89 ± 0.01	−12.85 ± 0.01	(xc)	(ob)
153	Mrk 926	0.0469	44.45	44.19	20.54	⋯	42.66 ± 0.03	42.66 ± 0.01	−12.05 ± 0.01	−12.68 ± 0.01	−13.05 ± 0.01	(xd)	(oe)
154	NGC 7582	0.0052	42.61	42.65	23.80	0.033	41.26 ± 0.43	40.38 ± 0.10	−12.39 ± 0.10	−12.10 ± 0.10	−12.88 ± 0.10	(xp)	(oa)

Note. This table summarizes X-ray and optical emission line ([O iii] λ5007, Hα, Hβ) properties of 95 Swift/BAT 9-month AGNs in Tueller et al. (2008) excluding Cen A, blazars, and those at low Galactic latitudes ($| b| \lt 15^\circ$). Columns: (1) Source no. in Tueller et al. (2008); (2) object name; (3) redshift; (4) 9-month averaged 14–195 keV luminosity calculated from the observed flux; (5) absorption-corrected 2–10 keV luminosity of the transmitted component averaged for 70 months; (6) X-ray absorption hydrogen column density (0.00 means N_H = 0); (7) soft X-ray scattering fraction; (8) [O iii] λ5007 luminosity corrected for extinction based on the Balmer decrement; (9) observed [O iii] λ5007 luminosity (with no extinction correction); (10) [O iii] λ5007 flux; (11) narrow Hα flux (number in parentheses refers to Hα/ Hβ flux ratio); (12) narrow Hβ flux; (13) reference for the X-ray spectra (columns (6) and (7)): (xa) C. Ricci et al. (2015, in preparation), (xb) Noguchi et al. (2010), (xc) Tueller et al. (2008), (xd) Winter et al. (2009a), (xe) Eguchi et al. (2011), (xf) Eguchi et al. (2009), (xg) Risaliti et al. (2009), (xh) Turner et al. (2009), (xi) Awaki et al. (2008), (xj) González-Martín et al. (2011), (xk) Shirai et al. (2008), (xl) Winter et al. (2009b), (xm) Comastri et al. (2010), (xn) Winter & Mushotzky (2010), (xo) Ballantyne (2005), (xp) Bianchi et al. (2009); (14) reference for the optical line fluxes (columns (8)–(12)): (oa) this work, (ob) Dahari & De Robertis (1988), (oc) Mulchay et al. (1994), (od) M. Koss et al. (2015, in preparation), (oe) Winter et al. (2010), (of) Bassani et al. (1999), (og) Xu et al. (1999), (oh) Landi et al. (2007). Columns (1)–(4) are taken from Tueller et al. (2008) except for the revised redshift of no. 53 (2MASX J06403799–4321211). Column (5) is taken from C. Ricci et al. (2015, in preparation). All luminosities are calculated from the redshift given in column (3) with (H₀, Ω_m, Ω_λ) = (70 km s⁻¹ Mpc⁻¹, 0.3, 0.7).

Download table as:  ASCIITypeset images: 1 2 3

Figure 1 plots the host inclination against log ${N}_{{\rm{H}}}$ for Sample A. In all plots of our paper, the diagonal crosses correspond to X-ray type 1 (X-ray unabsorbed) AGNs and the filled circles to X-ray type 2 (X-ray absorbed) AGNs, among which the open circles denote those with low scattering fractions. As noticed, 9 objects out of 10 with ${i}_{{\rm{host}}}$ > 85° are X-ray type 2 AGNs, rejecting the null hypothesis that X-ray absorption is independent of the host inclination at >98% confidence level. This is expected as galactic interstellar matter could produce an X-ray absorption of log ${N}_{{\rm{H}}}$ > 22 when viewed edge-on, and it is in agreement with the deficiency of nearly edge-on Seyfert 1 galaxies reported by Keel (1980). Except for that, there is no correlation between ${i}_{{\rm{host}}}$ and ${N}_{{\rm{H}}}$. These results are consistent with the random distribution of the orientation angle of the torus (or the accretion disk) with respect to that of the galactic plane, confirming previous findings (Schmitt et al. 2001). We note that the ${i}_{{\rm{host}}}$ distribution of the AGNs with low scattering fractions in our sample is not concentrated at large values; more than half of this population is free from absorption by interstellar matter along the galactic disk, which therefore cannot account for their observed low scattering fractions. Indeed, a K-S test for the ${i}_{{\rm{host}}}$ distribution between the low scattering fraction AGNs and the rest of X-ray type 2 AGNs in Sample A yields a matching probability of 0.53. By considering the small sample size, this does not necessarily contradict the statistical result by Hönig et al. (2014); they obtain a more edge-on dominant ${i}_{{\rm{host}}}$ distribution of the same population based on a slightly larger sample collected from the literature, although some of their sample is revised in our paper (see Section 2.1). We can conclude that there are at least two origins for their low scattering fractions: (1) intrinsic nature of the nucleus and (2) interstellar absorption in the host galaxy.

Zoom In Zoom Out Reset image size

We performed optical spectroscopic observations of Swift/BAT AGNs visible in the southern sky (δ < −10°) by using the SAAO 1.9 m telescope with the Cassegrain spectrograph during four observation runs: 2007 July, 2008 January, 2008 August, and 2009 February, each consisting of roughly 14 nights. In this paper, we focus on sources in the Swift/BAT 9-month catalog, although our observation targets at the SAAO also include those in the Swift/BAT 22-month catalog, whose results will be reported in M. Koss et al. (2015, in preparation). In total, the spectra of 38 AGNs have been analyzed in this work.

We used the 300 line mm⁻¹ grating, blazed at 6000 Å, covering about 4400–7600 Å, with a 2'' slit width placed on the center of each galaxy, producing a spectral resolution of ≈5 Å. The integration is split into a series of 150 s exposures, added up to a total integration time ranging from 750 to 3600 s. Wavelength calibration of the spectra was obtained from CuAr arc lamp exposures taken during the same night. A flux calibration was obtained from long-slit (with 6'' slit width) observations of spectrophotometric standard stars. To derive the sensitivity curve, we fit the observed spectral energy distribution of the standard stars with a low-order polynomial.

The spectral line flux of [O iii] λ5007 and those of narrow components of Hα and Hβ were measured using the IRAF task splot from the co-added, dispersion-corrected, and flux-calibrated spectra. If the lines were not significantly detected, we then estimated their upper limits (3σ) from the fluctuation of the noise level. The line fluxes are corrected for reddening from the Milky Way, by using the $E(B-V)$ map by Schlafly & Finkbeiner (2011) and the reddening curve by Cardelli et al. (1989) with R_V = 3.1. Finally, we approximately corrected these fluxes for the slit loss in the following way. For each dispersed spectrum, we projected the 4500–5500 Å region onto the spatial axis and measured its spread by fitting with a Gaussian. We then calculated the fraction contained within the slit by assuming that the image is axisymmetric. By comparing the results of the same target taken on different days when available, we estimate that the flux uncertainties are typically of 0.1–0.2 dex, depending on the quality of the spectrum. This is similar to general errors in the [O iii] λ5007 fluxes reported by Whittle (1992) when they are measured with small (2–4'') apertures. In some cases of broad-line AGNs, we were unable to reliably measure the fluxes (or the upper limits) of the narrow components of Hα and Hβ lines by separating them from the broad components.

To complement the results from the SAAO observations, we gather [O iii] λ5007, Hα, and Hβ fluxes in the literature for AGNs in the northern sky. We mainly adopt the results summarized by Winter et al. (2010), except for those with too large uncertainties, and refer to other references (Mulchay et al. 1994; Bassani et al. 1999; Xu et al. 1999; Landi et al. 2007; M. Koss et al. 2015, in preparation) for the rest. We obtain constraints on the [O iii] λ5007 flux for all 103 AGNs (48 X-ray type 1 and 55 X-ray type 2 AGNs) of the parent sample defined in Section 2.1 (Sample A), where one object (No. 49) does not show detectable [O iii] λ5007 emission and hence has only an upper limit. Among them, 77 objects (31 X-ray type 1 and 46 X-ray type 2 AGNs) have reliable flux measurements (not upper limits) of both narrow Hα and Hβ emission lines, F_Hα and F_Hβ, or their flux ratios, constituting "Sample B."

Table 1 lists the observed [O iii] λ5007 luminosity (${L}_{[{\rm{O}}\;{\rm{III}}]\;}$) along with the fluxes of [O iii] λ5007, narrow Hα, and narrow Hβ lines for Sample A with the reference of the optical spectroscopic data. For Sample B, we also calculate an extinction-corrected luminosity of [O iii] λ5007 (${L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}$) from the Balmer decrement as

$$\begin{eqnarray*}&&{L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}={L}_{[{\rm{O}}\;{\rm{III}}]}\;{(\displaystyle \frac{{F}_{{\rm{H}}\alpha }/{F}_{{\rm{H}}\beta }}{3.0})}^{2.94},\end{eqnarray*}$$

following Bassani et al. (1999). When the F_Hα/F_Hβ ratio is smaller than 3.0, we do not apply any correction. As discussed in Hao et al. (2005), however, the intrinsic flux ratio between Hα and Hβ in the NLR of an AGN could be different from the value assumed here, being subject to the gas density and radiative transfer effects. Also, there is an uncertainty in the correction because the spatial distributions of the [O iii] λ5007 and Balmer line emitting regions may not be the same owing to the clumpiness of the NLR (see Section 4.1). Thus, we should regard these corrections only as an approximation.

Figure 2 plots the correlation of the observed [O iii] λ5007 luminosity ($\mathrm{log}{L}_{[{\rm{O}}\;{\rm{III}}]}$) against (a) the luminosity in the 14–195 keV band (logL_14–195) or (b) that in the 2–10 keV band (logL_2–10), using Sample A. Figure 3 shows the same but for the [O iii] λ5007 luminosity corrected for extinction ($\mathrm{log}{L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}$), using Sample B. For each plot, we evaluate the strength of the luminosity–luminosity and flux–flux correlations separately for X-ray type 1, X-ray type 2, and all (X-ray type 1 + type 2) AGNs; the resultant Spearman's rank coefficients and Student's t-null significance levels are summarized in Table 2. We also calculate the ordinary least-squares bisector regression lines of the luminosity–luminosity correlation with the form of $Y=a+{bX}$, where Y is either $\mathrm{log}{L}_{[{\rm{O}}\;{\rm{III}}]}$ or $\mathrm{log}{L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}$ and X is either log L_14–195 or log L_2–10. The parameters and their 1σ errors are listed in Table 2. The best-fit lines obtained from all AGNs are plotted in Figures 2 and 3. When Sample A is used, we ignore the two objects whose [O iii] λ5007 luminosities are upper limits and restrict the luminosity range above log L_2–10 > 41.

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Table 2.  Correlation Properties between Different Luminosities

Y	X	Sample	N	${\rho }_{L}$	${\rho }_{f}$	PL	Pf	a	b
(1)	(2)	(3)	(4)	(5)	(6)	(7)	(8)	(9)	(10)
		All	101	0.672	0.419	$1.4\times {10}^{-14}$	$1.3\times {10}^{-5}$	−11.0 ± 2.9	1.20 ± 0.07
$\mathrm{log}{L}_{{\rm{[O\ III]}}}$	$\mathrm{log}{L}_{14-195}$	Type 1	48	0.781	0.240	$5.9\times {10}^{-11}$	$1.0\times {10}^{-1}$	−8.0 ± 3.8	1.13 ± 0.09
		Type 2	53	0.461	0.548	$5.1\times {10}^{-4}$	$2.2\times {10}^{-5}$	−10.3 ± 3.7	1.18 ± 0.09
		All	100	0.630	0.375	$2.1\times {10}^{-12}$	$1.2\times {10}^{-4}$	−10.0 ± 2.9	1.18 ± 0.07
$\mathrm{log}{L}_{{\rm{[O\ III]}}}$	$\mathrm{log}{L}_{2-10}$	Type 1	48	0.761	0.239	$3.4\times {10}^{-10}$	$1.0\times {10}^{-1}$	−5.5 ± 3.6	1.08 ± 0.09
		Type 2	52	0.447	0.488	$9.1\times {10}^{-4}$	$2.4\times {10}^{-4}$	−10.2 ± 4.0	1.19 ± 0.10
		All	76	0.594	0.475	$1.5\times {10}^{-8}$	$1.5\times {10}^{-5}$	−9.8 ± 4.0	1.18 ± 0.09
$\mathrm{log}{L}_{{\rm{[O\ III]}}}^{{\rm{cor}}}$	$\mathrm{log}{L}_{14-195}$	Type 1	31	0.705	0.341	$9.5\times {10}^{-6}$	$6.0\times {10}^{-2}$	−5.4 ± 4.2	1.08 ± 0.10
		Type 2	45	0.421	0.571	$4.0\times {10}^{-3}$	$4.2\times {10}^{-5}$	−12.0 ± 5.5	1.23 ± 0.13
		All	75	0.619	0.516	$3.2\times {10}^{-9}$	$2.2\times {10}^{-6}$	−8.5 ± 3.7	1.16 ± 0.09
$\mathrm{log}{L}_{{\rm{[O\ III]}}}^{{\rm{cor}}}$	$\mathrm{log}{L}_{2-10}$	Type 1	31	0.716	0.463	$5.9\times {10}^{-6}$	$8.8\times {10}^{-3}$	−1.6 ± 3.5	1.01 ± 0.09
		Type 2	44	0.503	0.552	$5.0\times {10}^{-4}$	$1.0\times {10}^{-4}$	−12.8 ± 5.3	1.26 ± 0.13
		All	82	0.608	0.352	$1.4\times {10}^{-9}$	$1.2\times {10}^{-3}$	−3.0 ± 3.1	1.02 ± 0.08
$\mathrm{log}{L}_{{\rm{H}}\alpha }$	$\mathrm{log}{L}_{2-10}$	Type 1	32	0.739	0.128	$1.3\times {10}^{-6}$	$4.8\times {10}^{-1}$	3.9 ± 3.9	0.87 ± 0.09
		Type 2	50	0.542	0.514	$4.9\times {10}^{-5}$	$1.3\times {10}^{-4}$	−4.0 ± 4.2	1.04 ± 0.10
		All	82	0.934	0.851	$2.1\times {10}^{-37}$	$4.6\times {10}^{-24}$	−7.6 ± 2.1	1.19 ± 0.05
$\mathrm{log}{L}_{{\rm{[O\ III]}}}$	$\mathrm{log}{L}_{{\rm{H}}\alpha }$	Type 1	32	0.937	0.799	$3.1\times {10}^{-15}$	$4.1\times {10}^{-8}$	−6.2 ± 2.7	1.16 ± 0.07
		Type 2	50	0.886	0.877	$1.2\times {10}^{-17}$	$6.4\times {10}^{-17}$	−7.6 ± 3.1	1.03 ± 0.11

Note. Columns: (1) Y variable; (2) X variable; (3) AGN type; (4) number of sample; (5) Spearman's rank coefficient for luminosity–luminosity correlation (ρ_L); (6) Spearman's rank coefficient for flux–flux correlation (ρ_f); (7) Student's t-null significance level for luminosity–luminosity correlation (P_L); (8) Student's t-null significance level for flux–flux correlation (P_f); (9) regression intercept (a) and its 1σ uncertainty; (10) slope (b) and its 1σ uncertainty. Equation is represented as $Y=a+{bX}.$

Download table as:  ASCIITypeset image

As shown in Table 2, we find significant correlations at confidence levels of >99% between all combinations of the [O iii] λ5007 and X-ray luminosities for any AGN types. The flux–flux correlations are weaker but significant at >90% confidence levels; relatively weak correlation is obtained for the X-ray type 1 AGN sample, most probably as a result of the narrow X-ray flux range (F_X ≃ 2 × 10⁻¹¹–3 × 10⁻¹⁰ erg cm⁻² s⁻¹ in the 14–195 keV band).

From the luminosity correlations for the entire AGN sample, we obtain b ≈ 1.2 in the regression line, which is significantly (>1σ) different from 1. This result is confirmed by the recent work based on a larger but less complete sample of Swift/BAT AGNs at z > 0.01 by Berney et al. (2015). We find that the slope for the X-ray type 1 AGNs is smaller than that for X-ray type 2 AGNs, although consistent within errors, given the large scatter of the correlations. The correlations with respect to L_14–195 and those to L_2–10 are found to be similar except for the normalizations. This is expected because absorption has a small effect on the observed hard X-ray luminosity (L_14–195) except for heavily Compton-thick AGNs, and L_2–10 is corrected for absorption through the X-ray spectral analysis.

We compare our results on the ${L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}$–L_2–10 correlation obtained from X-ray type 2 AGNs with previous works. The slope we obtain, b = 1.26 ± 0.13 (i.e., ${L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}$ $\propto \;{L}_{2-10}^{1.26\pm 0.13}$), is somewhat larger than that of Lamastra et al. (2009), who derive b = 0.98 ± 0.06 from a sample consisting of X-ray and optically selected Seyfert 2 galaxies. To check the effects of sample incompleteness of Sample B, we perform regression analysis with the asurv software (Isobe et al. 1986), by considering the lower limits of ${L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}$ to be ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ for the objects excluded in Sample B. We find that the slope b changes only by ∼0.01 compared with the case obtained from Sample B. Hence, the sample incompleteness cannot explain the difference of our result from Lamastra et al. (2009). The reason behind the discrepancy is not clear, but could be due to the different sample selections and luminosity ranges. Lamastra et al. (2009) include a sample compiled by Panessa et al. (2006) from the Palomar optical spectroscopic survey, which covers a lower luminosity range (L_2–10 < 10⁴² erg s⁻¹) than our sample. In fact, Panessa et al. (2006) obtain a much smaller slope, b = 0.75 ± 0.09, from their Seyfert 2 sample including Compton-thick AGNs, whose intrinsic X-ray luminosities are simply estimated by multiplying by a constant factor. The slope flatter than ours would be explained if contamination of [O iii] λ5007 from star formation in the host galaxy is more significant in lower-luminosity AGNs (see Section 4.1). Another possibility is enhanced past activity in the low-luminosity AGNs, which are left with a higher [O iii] λ5007 luminosity with respect to the current low X-ray activity.

We calculate the error-weighted mean value of the [O iii] λ5007 to X-ray luminosity ratio and its standard deviation for different AGN types. Here we consider a systematic error of 0.2 dex in $\mathrm{log}{L}_{[{\rm{O}}\;{\rm{III}}]}$ and 0.5 dex in $\mathrm{log}{L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}$ in addition to the errors listed in Table 1. The results are summarized in Table 3. Although we find that the best-fit regression line is not linear (b > 1), its effect can be checked by calculating an averaged $\mathrm{log}{L}_{{\rm{X}}}$ value in each sample, which is also listed in Table 3. In fact, we confirm that it little affects the following discussions.

Table 3.  Summary of Luminosity Ratios

Luminosity	Sample	N	$\lt r\gt$	σ	$\lt {L}_{{\rm{X}}}\gt$
Ratio	(1)	(2)	(3)	(4)	(5)
	All	102	−2.45 ± 0.06	0.60 ± 0.05	43.71
	Type 1	48	−2.27 ± 0.07	0.47 ± 0.05	43.85
$\mathrm{log}({L}_{{\rm{[O\ III]}}}/{L}_{14-195})$	Type 2	54	−2.62 ± 0.10	0.66 ± 0.07	43.60
	Type 2 (${f}_{{\rm{scat}}}\lt 0.5\%$)	12	−3.13 ± 0.10	0.33 ± 0.08	43.92
	Type 2 (${i}_{{\rm{host}}}\gt 80^\circ$)	10	−2.43 ± 0.13	0.41 ± 0.10	43.24
	All	101	−1.99 ± 0.07	0.63 ± 0.05	43.24
	Type 1	48	−1.78 ± 0.08	0.51 ± 0.06	43.36
$\mathrm{log}({L}_{{\rm{[O\ III]}}}/{L}_{2-10})$	Type 2	53	−2.19 ± 0.10	0.67 ± 0.07	43.14
	Type 2 (${f}_{{\rm{scat}}}\lt 0.5\%$)	12	−2.85 ± 0.09	0.31 ± 0.07	43.62
	Type 2 (${i}_{{\rm{host}}}\gt 80^\circ$)	10	−2.01 ± 0.17	0.53 ± 0.13	42.81
	All	77	−1.94 ± 0.08	0.69 ± 0.06	43.65
	Type 1	31	−1.89 ± 0.09	0.48 ± 0.07	43.79
log $({L}_{{\rm{[O\ III]}}}^{{\rm{cor}}}/{L}_{14-195})$	Type 2	46	−1.98 ± 0.12	0.80 ± 0.09	43.56
	Type 2 (${f}_{{\rm{scat}}}\lt 0.5\%$)	9	−2.44 ± 0.28	0.82 ± 0.21	43.87
	Type 2 (${i}_{{\rm{host}}}\gt 80^\circ$)	7	−1.63 ± 0.23	0.59 ± 0.17	43.06
	All	76	−1.48 ± 0.08	0.69 ± 0.06	43.18
	Type 1	31	−1.38 ± 0.10	0.51 ± 0.07	43.28
log $({L}_{{\rm{[O\ III]}}}^{{\rm{cor}}}/{L}_{2-10})$	Type 2	45	−1.55 ± 0.12	0.78 ± 0.09	43.11
	Type 2 (${f}_{{\rm{scat}}}\lt 0.5\%$)	9	−2.16 ± 0.24	0.71 ± 0.18	43.58
	Type 2 (${i}_{{\rm{host}}}\gt 80^\circ$)	7	−1.20 ± 0.29	0.76 ± 0.22	42.62
	All	83	−2.18 ± 0.07	0.61 ± 0.05	43.22
	Type 1	32	−1.95 ± 0.10	0.54 ± 0.07	43.36
$\mathrm{log}({L}_{{\rm{H}}\alpha }/{L}_{2-10})$	Type 2	51	−2.35 ± 0.09	0.60 ± 0.07	43.14
	Type 2 (${f}_{{\rm{scat}}}\lt 0.5\%$)	13	−2.97 ± 0.08	0.28 ± 0.06	43.49
	Type 2 (${i}_{{\rm{host}}}\gt 80^\circ$)	11	−2.28 ± 0.18	0.57 ± 0.13	42.73
	All	83	0.23 ± 0.04	0.32 ± 0.03	41.06
	Type 1	32	0.32 ± 0.05	0.27 ± 0.04	41.40
$\mathrm{log}({L}_{{\rm{[O\ III]}}}/{L}_{{\rm{H}}\alpha })$	Type 2	51	0.16 ± 0.05	0.34 ± 0.04	40.85
	Type 2 (${f}_{{\rm{scat}}}\lt 0.5\%$)	12	0.15 ± 0.09	0.30 ± 0.07	40.63
	Type 2 (${i}_{{\rm{host}}}\gt 80^\circ$)	10	0.26 ± 0.06	0.17 ± 0.05	40.52

Note. Columns: (1) luminosty ratio; (2) AGN type; (3) number of objects; (4) average; (5) standard deviation; (6) mean luminosity value of the denominator in the sample.

Download table as:  ASCIITypeset image

As noticed from Table 3, we find that the mean ratio of observed [O iii] λ5007 luminosity (${L}_{[{\rm{O}}\;{\rm{III}}]\;}$) to X-ray luminosity is significantly smaller in X-ray type 2 AGNs than in X-ray type 1 AGNs, by ≈0.4 dex, using Sample A. This trend remains the same for the extinction-corrected [O iii] λ5007 luminosity (${L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}$) obtained from Sample B, although the difference between X-ray type 1 and X-ray type 2 AGNs is reduced to 0.1–0.2 dex. The "reduction" is consistent with the fact that the mean extinction correction factor is larger in X-ray type 2 AGNs ($\langle$ ${L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}$/${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ $\rangle$ = 0.59 ± 0.09) than in X-ray type 1 AGNs ($\langle$ ${L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}$/${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ $\rangle$ = 0.29 ± 0.11). This indicates a higher degree of obscuration toward the NLR in X-ray type 2 AGNs, consistent with previous results (e.g., Dahari & De Robertis 1988; Mulchay et al. 1994; Meléndez et al. 2008). It may be explained if the torus, or its extended structure such as dusty outflow (Hönig et al. 2012), is large enough to block a part of the NLR.

[To investigate the nature of AGNs with low scattering fractions, we calculate the mean [O iii] λ5007 to X-ray luminosity ratios for two subsamples of X-ray type 2 AGNs: (1) those with ${f}_{{\rm{scat}}}$ < 0.5%, and (2) those hosted by edge-on galaxies (${i}_{{\rm{host}}}$ > 80°). The results are also listed in Table 3. We find that the mean extinction-corrected [O iii] λ5007 to X-ray luminosity ratio of the low scattering fraction AGNs is much smaller than that of the total X-ray type 2 AGN sample, while that of the edge-on galaxies does not differ from it within uncertainties. A simple χ² test shows that the difference of the mean value of ${L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}/{L}_{2-10}$ between the low scattering fraction AGNs and the other X-ray type 2 AGNs is significant at >99.9% confidence level. This is also noticeable from Figure 4(a), where we plot the ${L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}/{L}_{2-10}$ ratio against log ${N}_{{\rm{H}}}$ for Sample B.

Zoom In Zoom Out Reset image size

These results suggest that a significant fraction of low scattering fraction AGNs are indeed buried in a torus with very small opening angles as originally proposed by Ueda et al. (2007). This population of AGNs could contribute to reducing the averaged ${L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}/{L}_{{\rm{X}}}$ ratio in the total X-ray type 2 AGN sample compared with that of X-ray type 1 AGNs, because they are predominantly identified as X-ray type 2 AGNs owing to the large covering fraction by the torus. We can rule out the possibility that their low scattering fractions are merely the result of deficiency of scattering gas in the NLR. If this were the case, we should observe a similar fraction of low ${L}_{[{\rm{O}}\;{\rm{III}}]}/{L}_{{\rm{X}}}$ objects among the X-ray type 1 AGN sample. Figure 4(b) plots the ${L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}/{L}_{2-10}$ ratio against "X-ray Eddington ratio" (the 2–10 keV luminosity divided by the Eddington luminosity) using objects with available black hole masses in Winter et al. (2009a). No clear correlation is noticeable for the whole sample. The low scattering fraction AGNs do not always have high Eddington ratios, while Noguchi et al. (2010) report a possible negative correlation between the scattering fraction and Eddington ratio. Theoretically, deeply buried AGNs would be expected in the early growth phases of SMBHs with relatively small masses (hence with low luminosities). Thus, to further investigate the natures of this population, we need a larger sample of low-luminosity AGNs.

In AGNs, intense narrow Hα and Hβ lines are also produced from the NLR. Hence, we also perform regression analysis between the narrow Hα and X-ray luminosities, and that between the narrow Hα and [O iii] λ5007 luminosities, in the same way as done in Section 3.1. For each analysis, we utilize objects in Sample A that have available flux measurements of Hα or [O iii] λ5007. The correlation plots are displayed in Figures 5 and 6, respectively, together with the best-fit linear regression forms, which are given in Table 2. We also calculate the mean and standard deviation of the log(L_Hα/L_2–10) ratio and the log(L_{[O iii]}/L_Hα) ratio, which are summarized in Table 3.

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

We find that, for all AGNs, (1) L_Hα $\propto \;{L}_{2-10}^{1.02\pm 0.08}$ with a similarly large scatter (≈0.6 dex) to that seen in the ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ versus L_2–10 correlation, and (2) ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ $\propto \;{L}_{{\rm{H}}\alpha }^{1.19\pm 0.05}$ with a much smaller scatter (≈0.3 dex). The slope of the L_Hα versus L_2–10 correlation obtained from the X-ray type 2 AGNs, 1.04 ± 0.10, is larger than that obtained by Panessa et al. (2006) from their sample of 34 Seyfert 2s, 0.78 ± 0.09, which covers a lower luminosity range (L_2–10 < 10⁴² erg s⁻¹) than ours. Even though here we use only the luminosity of the "narrow" component of Hα, the regression slope and scatter between L_Hα and L_2–10 are similar to those found between "total" L_Hα (i.e., that including the broad component) and L_2–10 (e.g., Ho 2008).

Using so far the largest (N > 100) statistically complete sample of hard X-ray (E > 14 keV) selected AGNs in the local universe, we determine the statistical properties between [O iii] λ5007 and hard X-ray luminosities with the best accuracy. The linear regression form of ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ $\propto \;{L}_{2-10}^{1.18\pm 0.07}$ (${L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}$ $\propto \;{L}_{2-10}^{1.16\pm 0.09}$) is obtained from the whole sample (see Table 2). These results can be used as the reference for AGNs in the luminosity range of log L_2–10 = 41–46. We also find that the mean luminosity ratio between ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ and L_2–10 of X-ray type 2 AGNs is significantly smaller than that of X-ray type 1 AGNs. The difference is largely contributed by a population of low scattering fraction AGNs. Another important result is the very large variance in the log(L_{[O iii]}/L_2–10) ratio, corresponding to its standard deviation of ∼0.5 in X-ray type 1 AGNs and ∼0.7 in X-ray type 2 AGNs (see Table 3).

The nonlinear correlation (i.e., $b\ne 1$) between ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ and ${L}_{{\rm{X}}}$ may be explained by a combination of multiple effects. The first effect is the luminosity dependence of the AGN spectral energy distribution. The luminosity of the narrow lines is predominantly determined by the continuum flux of ultraviolet photons responsible for photoionization of the NLR gas, rather than the X-ray flux. Thus, if the spectral slope α (for the flux density ${F}_{\nu }\propto {\nu }^{-\alpha }$) between UV and hard X-rays above 2 keV is larger in more luminous AGNs as suggested by Scott & Stewart (2014), it works to make the [O iii] λ5007 to X-ray luminosity correlation steeper. Second, according to the luminosity-dependent unification model (Ueda et al. 2003; Ricci et al. 2013), the opening angle of a torus increases with luminosity, thus making the angular spread of the "NLR cone" larger in more luminous AGNs. This also leads to increased b. The third effect is due to the luminosity dependence of the NLR size in the radial direction, which is proportional to L^0.33±0.04 (Schmitt et al. 2003). Thus, the actual size of the NLR might be saturated in very luminous AGNs if the outer radius exceeds the scale height of the host galaxy (Netzer et al. 2004). The fourth effect is the contamination of ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ from star formation in low-luminosity AGNs, as mentioned in Section 3.1. The last two effects make the regression slope flatter than unity.

To better understand the origin of the observed luminosity correlations and scatters between ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ (or ${L}_{[{\rm{O}}\;{\rm{III}}]}^{{\rm{cor}}}$) and ${L}_{{\rm{X}}}$, comparison with the L_Hα and ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ correlation is useful. The Balmer lines are emitted by recombination as the result of photoionization, whereas the [O iii] λ5007 line is emitted via collisional excitation in the heated gas. Thus, the intensity ratio between Hα and [O iii] λ5007 depends on the physical parameters, such as the ionization parameter and density (Ferland & Netzer 1983). Also, the [O iii] λ5007 line comes preferentially from gas with a density of ∼10⁶ cm⁻³ unlike the Balmer lines, which come from a wide range of densities. In fact, detailed images of the NLR with Hubble Space Telescope for a few low-z objects (e.g., Evans et al. 1991; Fischer et al. 2013) show that much of the [O iii] λ5007 flux comes from clumpy structures. The effects of dust extinction inside the NLR, which may not be correctly measured with the Balmer decrement, make it even more complex. Hence, depending on how the NLR gas and dust is distributed, nonlinear correlation, as well as a significant scatter in the flux ratio between the Balmer lines and [O iii] λ5007 line, would be also expected.

The results for all AGNs, L_Hα $\propto \;{L}_{2-10}^{1.02\pm 0.08}$ and ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ $\propto \;{L}_{{\rm{H}}\alpha }^{1.19\pm 0.05},$ show that the observed nonlinear correlation (b ≈ 1.2) between ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ and ${L}_{{\rm{X}}}$ cannot be simply explained by a single reason. In addition to the four possibilities listed above, it is found that the nonlinear correlation between ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ and L_Hα, which is determined by plasma physics, also plays a role. The fact that the slope between L_Hα and L_2–10 is close to unity (b = 1.02 ± 0.08) indicates that the third effect (luminosity dependence of the NLR physical size) and/or the fourth effect (contamination by star formation) must work to cancel the first and second effects. The fact that flatter slopes are found from the X-ray type 1 AGNs, which are dominant in the largest luminosity range, suggests that the third effect is more important.

The large variation between ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ and ${L}_{{\rm{X}}}$ may be explained because the optical emission lines from the NLR are a secondary indicator of the intrinsic AGN luminosity in that they do not directly come from near the SMBH and have strong dependence on the geometry and size of the NLR, its averaged density, clumpiness, and amount of dust. In fact, a significant scatter of ∼0.4 dex between the [O iii] λ5007 luminosity and the continuum luminosity at 5100 Å is also reported in the Sloan Digital Sky Survey (SDSS) quasar sample (Shen et al. 2011). The presence of the low scattering fraction AGNs accounts for the larger scatter of the ${L}_{[{\rm{O}}\;{\rm{III}}]}/{L}_{{\rm{X}}}$ ratio in X-ray type 2 AGNs (∼0.7 dex) than in X-ray type 1 AGNs (∼0.5 dex), which could be understood in terms of variation in the geometry (cone angle) of the NLR. Since the correlation between ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ and L_Hα is found to be tighter than that between ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ and ${L}_{{\rm{X}}}$, the clumpiness of the NLR gas and dust extinction effects would not be the prime cause of the ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$–${L}_{{\rm{X}}}$ scatter. Another effect could be time variability; even though we utilize "70-month" averaged hard X-ray fluxes, the emission from the NLR reflects the past AGN power averaged over >10² yr.

The LF is one of the most important statistical properties of AGNs. Utilizing an AGN sample selected from the SDSS, Hao et al. (2005) determined [O iii] λ5007 and Hα LFs of AGNs at z ≤ 0.15 (they adopt emission-line luminosities not corrected for extinction, and we follow the same procedure below). Heckman et al. (2005) then compared the SDSS [O iii] λ5007 LF with an X-ray LF in the 3–20 keV band derived from the RXTE Slew Survey by Sazonov & Revnivtsev (2004). They found that the X-ray LF significantly underpredicts the [O iii] λ5007 LF when the mean luminosity ratio between ${L}_{{\rm{X}}}$ and ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ obtained from the RXTE AGN sample is assumed without considering the scatter. On the basis of this result, they argue that X-ray surveys seem to miss a significant fraction of AGNs, particularly Compton-thick AGNs.

Recently, Ueda et al. (2014) determined the X-ray LF of AGNs, including Compton-thick AGNs over a redshift range of z = 0–5, using a highly complete sample of X-ray selected AGNs. The local AGN sample from the Swift/BAT survey is also utilized. Detection biases against (mildly) Compton-thick AGNs are taken into account to correctly estimate their intrinsic number. Because heavily Compton-thick AGNs with log ${N}_{{\rm{H}}}$ = 25–26 are difficult to detect even in the E > 10 keV hard X-ray band, they assume that the fraction of AGNs with log ${N}_{{\rm{H}}}$ = 25–26 is the same as those with log ${N}_{{\rm{H}}}$ = 24–25. The X-ray LF and absorption distribution function are used as the basis of a standard population synthesis model of the X-ray background (Ueda et al. 2014). We note that the X-ray AGN LF by Sazonov & Revnivtsev (2004) may not be appropriate to adopt for direct comparison with LFs in other wavelengths because (1) the original X-ray LF by Sazonov & Revnivtsev (2004) was unfortunately affected by an error in the count-rate-to-flux conversion (by a factor of 1.4; see Sazonov et al. 2008; Ueda et al. 2011); (2) even after correcting for that error, it significantly underestimates other X-ray LFs of Compton-thin AGNs (Ueda et al. 2011); and (3) Compton-thick AGNs, which are difficult to detect in the 3–20 keV band, are not included.

Thus, it is very interesting to make comparison with the [O iii] λ5007 and Hα LFs with the most up-to-date X-ray LF of local AGNs, including Compton-thick AGNs, in order to understand the completeness and cleanness of AGN selections in these different wavelengths. The red curve in Figures 7(a) and (b) represent the best-fit [O iii] λ5007 and narrow Hα LFs in Hao et al. (2005) (two-power-law model, the sum of Seyfert 1s and 2s), after correcting both luminosity and space density for the difference of the adopted Hubble constant, from H₀ = 100 km s⁻¹ Mpc⁻¹ (Hao et al. 2005) to H₀ = 70 km s⁻¹ Mpc⁻¹ (our paper). The black curve in each figure is a prediction for the [O iii] λ5007 (or Hα) LF calculated from the Ueda et al. (2014) X-ray LF at z = 0. Here we convert L_2–10 into ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ (or Hα) with the best-fit linear regression form (Table 2) separately for X-ray type 1 and X-ray type 2 AGNs, and we also consider the scatter around it by assuming a Gaussian distribution with the standard deviation listed in Table 3. As the Hao et al. (2005) result is obtained from AGNs at z ≤ 0.15, we then multiply luminosity-dependent density evolution factors (Ueda et al. 2014) at the mean redshift. The black dashed curves denote the boundaries when both errors (1σ) in the mean and standard deviation of log(L_{[O iii]}/L_2–10) (or log(L_Hα/L_2–10)) are taken into account. For comparison, we also plot the case when the standard deviation is set to be zero (i.e., no scatter is considered) with the blue, dot-dashed curve.

Zoom In Zoom Out Reset image size

As noticed from Figure 7(a), the [O iii] λ5007 (red) and X-ray (black) LFs are roughly consistent with each other within a factor of ∼2 when we take into account the uncertainties in the ${L}_{{\rm{X}}}$ to ${L}_{[{\rm{O}}\;{\rm{III}}]\;}$ conversion. Thus, the systematic (≈4) underestimate of the [O iii] λ5007 LF by the X-ray LF over a wide range of luminosity reported by Heckman et al. (2005) is now resolved. Rather, at $\mathrm{log}{L}_{[{\rm{O}}\;{\rm{III}}]}$ ≳ 40, the X-ray LF outnumbers the [O iii] λ5007 LF, while statistical uncertainties in the [O iii] λ5007 LF are large (a factor of >2) at $\mathrm{log}{L}_{[{\rm{O}}\;{\rm{III}}]}$ ≳ 41.6 owing to the limited sample size in the SDSS. Figure 7(b) shows even better agreement between the Hα and X-ray LFs over a wider luminosity range, although a similar discrepancy is noticed at log L_Hα ≳ 41. We note that it is important to consider the scatter between the two luminosities when making the comparison of LFs, as seen in the difference between the black solid curve (with scatter) and blue dot-dashed curve (without scatter).

These results confirm that hard X-ray (>10 keV) observations are a very powerful tool to find AGNs with high completeness, not missing a dominant portion of the entire AGN population, once biases against Compton-thick AGNs are properly corrected (see, e.g., Malizia et al. 2009). For the correction, however, it is essential to obtain the broadband X-ray spectra covering up to, at least, a few tens of keV, with sufficiently good sensitivities. If the discrepancy between the [O iii] λ5007 (or Hα) and X-ray LFs at the high-luminosity range is true, this instead implies that the optical selection would miss some AGN populations. The selection based on emission-line diagrams could be incomplete for AGNs significantly contaminated by star formation; indeed, Winter et al. (2010) show that a nonnegligible fraction of hard X-ray selected AGNs could be optically classified as H ii galaxies, even though they are truly AGNs. Other candidates of "optically missing" AGNs are those deeply embedded in tori with almost spherical geometry, in which no or little NLR is formed. They may be similar to some of the low scattering fraction AGNs in our sample whose [O iii] λ5007 fluxes are very weak. If many of the heavily Compton-thick AGNs assumed in the Ueda et al. (2014) model correspond to this population, it would partially account for the mismatch between the optical and X-ray LFs.