Discovery and analysis of a ULX nebula in NGC 3521

TL;DR

This study combines optical and X-ray observations to analyze a ULX in NGC 3521, revealing its nebular environment, emission mechanisms, and updated X-ray luminosity, contributing to understanding ULX surroundings and properties.

Contribution

It provides the first detailed spectral analysis of the ULX nebula and refines the X-ray luminosity measurement, highlighting shock and photoionisation effects in the nebula.

Findings

ULX nebula is spatially extended (~34 pc) and connected to an H II region.

Emission lines indicate shock and photoionisation in low metallicity gas.

X-ray luminosity is approximately four times higher than previous estimates.

Abstract

We present Very Large Telescope/X-shooter and Chandra X-ray observatory/ACIS observations of the ULX [SST2011] J110545.62+000016.2 in the galaxy NGC 3521. The source identified as a candidate near-infrared counterpart to the ULX in our previous study shows an emission line spectrum of numerous recombination and forbidden lines in the visible and near-infrared spectral regime. The emission from the candidate counterpart is spatially extended ($\sim$ 34 pc) and appears to be connected with an adjacent H II region, located $\sim$ 138 pc to the NE. The measured velocities of the emission lines confirm that both the candidate counterpart and H II region reside in NGC 3521. The intensity ratios of the emission lines from the ULX counterpart show that the line emission originates from the combined effect of shock and photoionisation of low metallicity (12 + log (O/H) = 8.19 $\pm$ 0.11) gas. Unfortunately, there is no identifiable spectral signature directly related to the photosphere of the mass-donor star in our spectrum. From the archival Chandra data, we derive the X-ray luminosity of the source in the 0.3-7 keV range to be (1.9 $\pm$ 0.8) $\times$ 10$^{40}$ erg cm$^{-2}$ s$^{-1}$, almost a factor of four higher than what is previously reported.

Tables

Table 1. Table 1 : Selected properties of the host galaxy of the observed ULX counterpart.

Galaxy	Morphological	Distance	Axis ratio	Inclination	ULX distance to
	type	(Mpc)		(deg)	galaxy nucleus (kpc)
NGC 3521	SAB(rs)bca	14.19 $\pm$ 2.84b	0.51c	73d	10.46 $\pm$ 2.09
References: aBagetakos et al. (2011), bWalter et al. (2008), cTully et al. (2013) and dJarrett et al. (2003).

Table 2. Table 2 : Theoretical and restframe wavelengths, integrated flux and intrinsic FWHM of the emission lines in the UVB, VIS and NIR data for both source B and J1105.

	Theoretical	Restframe		Integrated flux		FWHMint
Ion	wavelength	wavelength (Å)$†$		($\times {10}^{-17}$ erg cm-2 s-1)		(Å)
	(Å) *	Source B	J1105	Source B	J1105	Source B	J1105
$[$O ii$]$	3726.032 $\pm$ 0.010	3726.21 $\pm$ 0.04	3726.11 $\pm$ 0.03	(1.63 $\pm$ 0.01)$\times {10}^2$	11.3 $\pm$ 0.1	0.59 $\pm$ 0.03	0.84 $\pm$ 0.03
$[$O ii$]$	3728.815 $\pm$ 0.010	3728.88 $\pm$ 0.03	3728.91 $\pm$ 0.03	(2.29 $\pm$ 0.03)$\times {10}^2$	16.9 $\pm$ 0.1	0.63 $\pm$ 0.03	0.77 $\pm$ 0.03
$[$Fe i$]$	3744.10236 $\pm$ 0.0002	-	3744.22 $\pm$ 0.03	-	7.81 $\pm$ 0.72	-	0.65 $\pm$ 0.03
$[$Ne iii$]$	3868.76 $\pm$ 0.10	3868.79 $\pm$ 0.01	3868.8 $\pm$ 0.1	4.83 $\pm$ 0.01	3.1 $\pm$ 0.1	0.7 $\pm$ 0.1	0.7 $\pm$ 0.1
$[$He i$]$	3888.41607 $\pm$ 0.00021	3888.54 $\pm$ 0.04	-	16.2 $\pm$ 0.1	-	0.86 $\pm$ 0.03	-
H$ϵ$	3970.0788 $\pm$ 0.0022	3970.14 $\pm$ 0.03	-	19.4 $\pm$ 0.1	-	0.68 $\pm$ 0.03	-
H$\delta$	4101.7415 $\pm$ 0.0024	4101.81 $\pm$ 0.03	-	39.9 $\pm$ 0.1	-	0.72 $\pm$ 0.03	-
H$\gamma$	4340.471 $\pm$ 0.003	4340.49 $\pm$ 0.03	4340.24 $\pm$ 0.03	77.6 $\pm$ 0.3	2.59 $\pm$ 1.09	0.77 $\pm$ 0.03	1.04 $\pm$ 0.03
$[$Fe ii$]$	4436.9204 $\pm$ 0.0003	-	4436.7 $\pm$ 0.1	-	0.55 $\pm$ 0.04	-	0.7 $\pm$ 0.1
$[$Fe ii$]$	4452.09585 $\pm$ 0.00012	-	4452.1 $\pm$ 0.1	-	1.23 $\pm$ 0.01	-	0.7 $\pm$ 0.1
H$\beta$	4861.333 $\pm$ 0.003	4861.39 $\pm$ 0.03	4861.54 $\pm$ 0.04	(1.95 $\pm$ 0.01)$\times {10}^2$	8.86 $\pm$ 0.08	0.87 $\pm$ 0.03	0.87 $\pm$ 0.03
$[$O iii$]$	4958.911 $\pm$ 0.010	4959.01 $\pm$ 0.03	4958.88 $\pm$ 0.03	60.0 $\pm$ 0.2	6.71 $\pm$ 0.19	0.83 $\pm$ 0.03	1.12 $\pm$ 0.03
$[$O iii$]$	5006.843 $\pm$ 0.010	5006.81 $\pm$ 0.03	5006.88 $\pm$ 0.03	(1.78 $\pm$ 0.02)$\times {10}^2$	18.4 $\pm$ 0.4	0.83 $\pm$ 0.03	0.83 $\pm$ 0.03
$[$O i$]$	6300.304 $\pm$ 0.010	6300.5 $\pm$ 0.2	6300.46 $\pm$ 0.01	7.11 $\pm$ 0.40	5.83 $\pm$ 0.03	0.8 $\pm$ 0.5	0.79 $\pm$ 0.04
$[$S iii$]$	6311.21 $\pm$ 0.10	6311.6 $\pm$ 0.1	-	1.86 $\pm$ 0.64	-	1.5 $\pm$ 0.5	-
$[$N ii$]$	6548.05 $\pm$ 0.10	6548.11 $\pm$ 0.01	6547.87 $\pm$ 0.01	74.6 $\pm$ 1.0	7.03 $\pm$ 0.07	0.76 $\pm$ 0.03	0.91 $\pm$ 0.03
H$\alpha$	6562.819 $\pm$ 0.007	6562.85 $\pm$ 0.01	6562.75 $\pm$ 0.03	(7.57 $\pm$ 0.07)$\times {10}^2$	36.0 $\pm$ 0.1	0.88 $\pm$ 0.03	0.96 $\pm$ 0.03
$[$N ii$]$	6583.45 $\pm$ 0.10	6583.48 $\pm$ 0.01	6583.38 $\pm$ 0.03	(2.22 $\pm$ 0.05)$\times {10}^2$	24.7 $\pm$ 0.1	0.85 $\pm$ 0.03	0.86 $\pm$ 0.03
$[$S ii$]$	6716.44 $\pm$ 0.01	6716.62 $\pm$ 0.01	6716.63 $\pm$ 0.01	76.0 $\pm$ 0.7	15.0 $\pm$ 0.3	0.79 $\pm$ 0.04	0.98 $\pm$ 0.01
$[$S ii$]$	6730.816 $\pm$ 0.010	6730.92 $\pm$ 0.01	6730.83 $\pm$ 0.01	55.7 $\pm$ 0.5	8.51 $\pm$ 0.04	0.81 $\pm$ 0.04	0.97 $\pm$ 0.04
He i	6678.1517 $\pm$ 0.0003	6678.11 $\pm$ 0.03	-	7.4 $\pm$ 0.1	-	0.81 $\pm$ 0.03	-
$[$Ar iii$]$	7135.79 $\pm$ 0.04	7135.84 $\pm$ 0.03	-	18.6 $\pm$ 0.1	-	0.79 $\pm$ 0.03	-
$[$O ii$]$	7319.99	7320.11 $\pm$ 0.01	-	4.2 $\pm$ 0.1	-	0.94 $\pm$ 0.03	-
?	-	7325.71 $\pm$ 0.03	-	4.46 $\pm$ 0.03	-	0.64 $\pm$ 0.03	-
$[$O ii$]$	7329.67 $\pm$ 0.1	7329.59 $\pm$ 0.01	-	2.25 $\pm$ 0.02	-	0.65 $\pm$ 0.03	-
$[$O ii$]$	7330.73 $\pm$ 0.1	7330.84 $\pm$ 0.03	-	2.28 $\pm$ 0.01	-	0.83 $\pm$ 0.03	-
$[$Ar iii$]$	7751.11 $\pm$ 0.05	7751.14 $\pm$ 0.03	-	4.31 $\pm$ 0.01	-	0.82 $\pm$ 0.03	-
Pa10	9014.910 $\pm$ 0.007	9015.03 $\pm$ 0.03	-	6.69 $\pm$ 0.65	-	1.15 $\pm$ 0.04	-
$[$S iii$]$	9069.0 $\pm$ 1	9068.84 $\pm$ 0.02	9069.74 $\pm$ 0.03	46.2 $\pm$ 1.3	2.38 $\pm$ 0.22	1.05 $\pm$ 0.01	2.1 $\pm$ 0.1
Pa9	9229.014 $\pm$ 0.007	9229.04 $\pm$ 0.01	-	7.99 $\pm$ 0.96	-	1.18 $\pm$ 0.01	-
$[$S iii$]$	9530.6 $\pm$ 1	9530.81 $\pm$ 0.02	9530.64 $\pm$ 0.01	(1.21 $\pm$ 0.25)$\times {10}^2$	8.21 $\pm$ 0.19	1.22 $\pm$ 0.03	1.69 $\pm$ 0.04
Pa$ϵ$	9545.969 $\pm$ 0.006	9545.47 $\pm$ 0.06	-	13.7 $\pm$ 2.0	-	1.3 $\pm$ 0.1	-
$[$Fe i$]$	10048.5845 $\pm$ 0.0014	-	10048.0 $\pm$ 0.1	-	3.9 $\pm$ 2.5	-	0.7 $\pm$ 0.3
Pa$\delta$	10049.369 $\pm$ 0.007	10049.9 $\pm$ 0.1	-	19 $\pm$ 2	-	1.6 $\pm$ 0.1	-
He i	10830.33977 $\pm$ 0.00005	10830.4 $\pm$ 0.1	10830.4 $\pm$ 0.1	70.3 $\pm$ 0.3	3.65 $\pm$ 0.30	1.9 $\pm$ 0.1	2.1 $\pm$ 0.1
Pa$\gamma$	10938.086 $\pm$ 0.01	10938.2 $\pm$ 0.1	-	34.7 $\pm$ 0.1	-	1.7 $\pm$ 0.1	-
$[$Fe ii$]$	12566.7688 $\pm$ 0.0009	12567.3 $\pm$ 0.4	-	8.28 $\pm$ 1.94	-	2.5 $\pm$ 0.6	-
3He i	12785.50005 $\pm$ 0.00006	12785.9 $\pm$ 0.5	12785.1 $\pm$ 0.1	2.85 $\pm$ 0.01	0.75 $\pm$ 0.07	1.6 $\pm$ 1.2	1.8 $\pm$ 0.1
$[$N iii$]$	12809 $\pm$ 1	12809.5 $\pm$ 0.6	-	3.063 $\pm$ 0.001	-	1.5 $\pm$ 1.6	-
Pa$\beta$	12818.07 $\pm$ 0.01	12818.2 $\pm$ 0.1	12818.1 $\pm$ 0.1	70.6 $\pm$ 0.1	2.22 $\pm$ 0.01	1.8 $\pm$ 0.1	2.1 $\pm$ 0.1
Br19	15260.555 $\pm$ 0.009	15259.6 $\pm$ 0.1	-	1.12 $\pm$ 0.05	-	3.5 $\pm$ 0.1	-
Br16	15556.466 $\pm$ 0.015	15556.0 $\pm$ 0.1	-	1.37 $\pm$ 0.07	-	3.1 $\pm$ 0.1	-
Br13	16109.330 $\pm$ 0.013	16109.3 $\pm$ 0.1	-	2.39 $\pm$ 0.09	-	3.4 $\pm$ 0.1	-
Br12	16407.210 $\pm$ 0.016	16406.3 $\pm$ 0.1	-	2.07 $\pm$ 0.21	-	3.1 $\pm$ 0.2	-
$[$Fe ii$]{}^{‡}$	16435.5271 $\pm$ 0.0019	16434.3 $\pm$ 0.1	16436.3 $\pm$ 0.1	5.70 $\pm$ 0.87	0.92 $\pm$ 0.10	5.7 $\pm$ 0.1	2.1 $\pm$ 0.1
Br11	16806.538 $\pm$ 0.017	16807.0 $\pm$ 0.1	-	4.09 $\pm$ 0.09	-	2.3 $\pm$ 0.1	-
Br10	17362.127 $\pm$ 0.018	17363.1 $\pm$ 0.1	-	4.47 $\pm$ 0.16	-	2.9 $\pm$ 0.1	-
Br9	18174.141 $\pm$ 0.020	18173.7 $\pm$ 0.1	-	5.35 $\pm$ 1.64	-	2.9 $\pm$ 0.1	-
Pa$\alpha$	18750.976 $\pm$ 0.025	18751.4 $\pm$ 0.1	18751.1 $\pm$ 0.1	91.8 $\pm$ 3.1	3.99 $\pm$ 0.01	2.7 $\pm$ 0.1	2.6 $\pm$ 0.1
*Wavelengths are indicated in air, not vacuum, and the reference values are taken from Kramida et al. (2018).
†Corrected for heliocentric motion and peculiar and radial velocity of NGC 3521.
‡The properties of this line are possibly contaminated by the presence of a sky emission line close to the observed wavelength.

Table 3. Table 3 : Electron temperature and density, and ion abundances for source B and J1105

Parameter	Source B	J1105
Te	8130 $\pm$ 1213 K	$>$ 5270 K
$n_e[$O ii$]$	50 $\pm$ 15 cm-3	5 $\pm$ 5 cm-3
$n_e[$S ii$]$	62 $\pm$ 24 cm-3	-
12 + $\log$ (O ii/H ii)	7.93 $\pm$ 0.14	-
12 + $\log$ (O iii/H ii)	7.83 $\pm$ 0.19	-
12 + $\log$ (O/H)	8.19 $\pm$ 0.11	-
12 + $\log$ (S ii/H ii)	6.43 $\pm$ 0.14	-
12 + $\log$ (S iii/H ii)	6.69 $\pm$ 0.15	-
12 + $\log$ (S/H)	6.97 $\pm$ 0.12	-
$\log$ R23	0.52 $\pm$ 0.03	-
$\log$ S23	0.182 $\pm$ 0.001	-
$\log U$	$-$2.84 $\pm$ 0.08	-
$\log {\eta }^{\prime }$	0.30 $\pm$ 0.01	-

Table 4. Table 4 : Observed and dereddened emission line flux ratios and Balmer decrements for source B and J1105. For the dereddening we used E ​ ( B − V ) = 0.28 ± 0.01 𝐸 𝐵 𝑉 plus-or-minus 0.28 0.01 E(B-V)=0.28\pm 0.01 and E ​ ( B − V ) = 0.27 ± 0.03 𝐸 𝐵 𝑉 plus-or-minus 0.27 0.03 E(B-V)=0.27\pm 0.03 for source B and J1105, respectively (see text).

Line	Source B		J1105
flux ratio	Observed	Dereddened	Observed	Dereddened
$\log \left(\right[$O iii$]\lambda 5007/$H$\beta ){}^a$	$-$0.04 $\pm$ 0.02	$-$0.06 $\pm$ 0.02	0.32 $\pm$ 0.01	0.30 $\pm$ 0.01
$\log \left(\right[$O i$]\lambda 6300/$H$\alpha ){}^a$	$-$2.03 $\pm$ 0.02	$-$2.02 $\pm$ 0.02	$-$0.79 $\pm$ 0.01	$-$0.78 $\pm$ 0.0
$\log \left(\right[$N ii$]\lambda 6583/$H$\alpha ){}^a$	$-$0.53 $\pm$ 0.02	$-$0.54 $\pm$ 0.02	$-$0.16 $\pm$ 0.01	$-$0.17 $\pm$ 0.01
$\log \left(\right[$S ii$]\lambda 6716+\lambda 6731/$H$\alpha ){}^{a,b,c}$	$-$0.76 $\pm$ 0.02	$-$0.77 $\pm$ 0.02	$-$0.185 $\pm$ 0.004	$-$0.194 $\pm$ 0.004
$\log \left(\right[$N ii$]\lambda 6548+\lambda 6583$/H$\alpha ){}^c$	$-$0.41 $\pm$ 0.02	$-$0.41 $\pm$ 0.02	$-$0.055 $\pm$ 0.003	$-$0.056 $\pm$ 0.003
$\log \left(\right[$S iii$]\lambda 9069+\lambda 9530$/H$\alpha ){}^c$	$-$0.7 $\pm$ 0.2	$-$0.8 $\pm$ 0.2	$-$0.53 $\pm$ 0.03	$-$0.65 $\pm$ 0.03
$\log \left(\right[$O ii$]\lambda 3727$/H$\beta ){}^d$	N/A	N/A	0.28 $\pm$ 0.01	0.39 $\pm$ 0.01
$\log \left(\right[$N ii$]\lambda 6583$/H$\beta ){}^d$	N/A	N/A	0.45 $\pm$ 0.01	0.31 $\pm$ 0.01
$\log \left(\right[$S ii$]\lambda 6731$/H$\beta ){}^d$	N/A	N/A	$-$0.02 $\pm$ 0.01	$-$0.16 $\pm$ 0.01
H$\alpha /$H$\beta$	3.88 $\pm$ 0.01	2.81 $\pm$ 0.01	4.06 $\pm$ 0.12	2.97 $\pm$ 0.12
H$\gamma /$H$\beta$	0.398 $\pm$ 0.003	0.452 $\pm$ 0.003	0.29 $\pm$ 0.03	0.33 $\pm$ 0.03
H$\delta /$H$\beta$	0.205 $\pm$ 0.001	0.246 $\pm$ 0.00	-	-
H$ϵ/$H$\beta$	0.10 $\pm$ 0.01	0.12 $\pm$ 0.01	-	-
Notes: aDiagnostic line flux ratios in the BPT diagram, to distinguish AGNs, LINERs and H ii regions
(Kauffmann et al., 2003; Kewley et al., 2006). bLine ratio that distinguishes between shock and photoionised
regions (Dodorico et al., 1980). cDiagnostic line flux ratios to distinguish supernova remnants with shocks
from H ii regions (Levenson et al., 1995). dLine flux ratios indicative of shock velocities, from the diagrams
Dopita et al. (1984). N/A = not applicable.

Equations

$$m_{H}=-2.5\log\bigg{(}\frac{F_{H}}{F_{H,0}}\bigg{)}$$

Full text

††pubyear: 2019††pagerange: Discovery and analysis of a ULX nebula in NGC 3521††thanks: Based on observations made with ESO Telescopes at the La Silla Paranal Observatory under programme 0100.D-0660(B).–A

1 Introduction

An ultraluminous X-ray source (ULX) is defined as a point-like, off-nuclear source with an X-ray luminosity $L_{X}$ larger than the Eddington luminosity for a 10 M*⊙* black hole, i.e. $L_{X}>10^{39}$ erg s*-1* (Kaaret et al., 2017). To explain these high luminosities, different possibilities on the nature of the accretors powering the ULX are considered. The first possibility is the ULX being powered by a stellar mass compact object with either emitting anisotropically (King et al., 2001), or accreting at super-Eddington rates (Begelman, 2002; Moon et al., 2003; Gladstone et al., 2009). Though a bona-fide black hole (BH) ULX is yet to be confirmed, several neutron star ULXs have been discovered in the past years (e.g. Bachetti et al. 2014; Fürst et al. 2016; Israel et al. 2016, 2017; Carpano et al. 2018), in line with the proposed super-Eddington accretion.

The second possibility is that the accretor is a BH more massive than 10M*⊙* (e.g. Zampieri & Roberts 2009), i.e. with masses similar to the BHs whose merger produced gravitational waves (e.g. Abbott et al. 2016). The third possiblity is that the ULX harbors a BH much more massive than the systems in the first two scenarios, but less massive than the BHs in the center of most galaxies. These type of BH would accrete at sub-Eddington rates, have masses between 102 and 105M*⊙* and are called intermediate mass black holes (IMBHs) (e.g. Farrell et al. 2011; Mezcua et al. 2013; Earnshaw 2016).

A reliable way to identify which accretor powers the ULXs is by dynamical mass measurements. To date, a mass constraint is available for the neutron star ULX M82-X2 (Bachetti et al., 2014), estimated through the detection of pulsations in the source. That technique can only be used for ULXs powered by neutron stars, so in the absence of pulsations, astronomers have focused on detecting the donor star (e.g. Patruno & Zampieri 2008; Heida et al. 2016; Vinokurov et al. 2018; Qiu et al. 2019). If detected, spectroscopic observations can be used to constrain the radial velocity amplitude, and hence, provided the orbital period is determined as well from those data, the black hole mass function (e.g. Motch et al. 2011).

On the other hand, other studies have focused on the analysis of the environment in which a ULX is embedded. Thanks to their high luminosities, ULXs can have a strong effect on their surroundings, i.e. ionising the gas around them. This is done either by photoionisation due to the high X-ray and UV luminosity of the ULX, or by shock ionisation caused by jets, or outflows or disc winds. Examples of ULXs photoionising the nebula surrounding them are Holmberg II X-1 (Pakull & Mirioni, 2002; Kaaret et al., 2004; Moon et al., 2011) and NGC 5408 X-1 (Soria et al., 2006); while ULXs responsible of shock ionising the material around them are IC 342 X-1 (Grisé et al., 2006b; Cseh et al., 2012), Holmberg IX X-1 (Miller, 1995; Grisé et al., 2006a; Moon et al., 2011) and M51 ULX-1 (Urquhart et al., 2018).

To distinguish whether a nebula surrounding the ULX is photoionised or shock ionised, one should look at the spectra. A high $[$S ii$]/H\alpha$ line ratio, i.e. $>0.4$, is a sign of shock ionisation (Dodorico et al., 1980). The electron temperature $T_{e}$ in a shock ionised region is higher than in a photoionised region (Urquhart et al., 2018). The shock ionisation models from Allen et al. (2008) show that a high Balmer decrement can be caused by ionising shocks with velocities $>100$ km s*-1*. Furthermore, the diagnostic diagrams from Dopita et al. (1984) that use the $[$O ii$]\lambda 3727$/H$\beta$, $[$O iii$]\lambda 5007$/H$\beta$, $[$N ii$]\lambda 6584$/H$\beta$, and $[$S ii$]\lambda 6731$/H$\beta$ ratios, could shed some light on the shock velocities of the ionised gas. There is, sometimes, the presence of collimated jets (e.g. Pakull et al. 2010; Soria et al. 2010; Soria et al. 2014) in shock ionised nebulae, evident in the elongated morphology of the nebula with symmetrical lobes. Nonetheless, in a large number of cases the nebula has both photoionised and shock ionised gas present (e.g. Roberts et al. 2003; Abolmasov et al. 2007a, b; Urquhart et al. 2018).

We performed a systematic search in the NIR to identify possible counterparts to nearby ULXs (López et al., 2017), where we detected several red supergiant candidates, some of which we followed-up spectroscopically. In this manuscript, we present the results for one of these sources in particular, [SST2011] J110545.62+000016.2 (hereafter J1105) in the galaxy NGC 3521, which turned out to be a nebula. We describe the target in Section 2 and the observations and data reduction of the spectra in Section 3. Our results are presented in Section 4 and discussed in detail in Section 5. We conclude in Section 6.

2 Target

We present the analysis of VLT/X-Shooter spectra of the NIR counterpart to J1105 in NGC 3521. López et al. (2017) detected the NIR counterpart for J1105, with an absolute magnitude $H=-10.93\pm 0.93$ (consistent with it being a red supergiant candidate), assuming a distance to NGC 3521 of 14.19 $\pm$ 2.84 Mpc (Walter et al., 2008) (see Table 1). J1105 is located on the northern side of the spiral arms of NGC 3521, approximately 10.9 kpc (152″) from its nucleus. According to Coccato et al. (2018), the radial velocity of NGC 3521 at 30″ (2.1 kpc) from its center is 175 $\pm$ 50 km s*-1*. Bagetakos et al. (2011) estimated that at 23.8 kpc (459″) the radial velocity of Hi is 220 $\pm$ 10 km s*-1*. It is worth noting that Bagetakos et al. (2011) assumed a distance to NGC 3521 of 10.7 Mpc. If we assume a distance of 14.19 $\pm$ 2.84 Mpc, 459″ corresponds to 31.6 kpc. Since we do not have any more information about the radial velocity of NGC 3521 at other positions, only the two mentioned before, we adopt the galactic radial velocity at the location of J1105 as 200 $\pm$ 25 km s*-1*.

3 Observations and data reduction

3.1 X-Shooter data

We obtained spectra of the counterpart to J1105 with X-Shooter (Vernet et al., 2011) on the Very Large Telescope (VLT) UT3 at Cerro Paranal under programme ID 0100.D-0660(B) on 2018 March 31. The spectra were taken in service mode.

X-Shooter has three spectroscopic arms that together provide spectral coverage from the near-UV to the near-IR. We used a slit width of 0.8″ in the UVB arm, 0.7″ in the VIS arm and 0.6″ in the NIR arm, resulting in resolutions of $\sim$ 6200, $\sim$ 11400 and $\sim$ 8100, respectively. Five ABBA nodding sequences with a 5″ nodding throw were used. The integration times for individual exposures for the UVB, VIS and NIR arms were 260, 210 and 285 s, respectively. The total exposure times were 5200 s for the UVB arm, 4200 s for the VIS arm and 5700 s for the NIR arm. The airmass during our observations was $\sim$ 1.1, and the seeing at the start and end of our observations was 0.39″ and 0.4$\arcsec$, respectively. The slit was aligned to the parallactic angle, which at the beginning of the observation was 206*∘* and at the end of the observation, 180*∘*. During our observations, the night was between clear and photometric.

We reduce the data with the Reflex environment (Freudling et al., 2013), processing the observations of J1105 and a standard star observed at similar airmass and close in time to the target. In the Reflex environment, the workflow designed for X-Shooter data produces flat-fielded, sky-subtracted, wavelength- and flux- calibrated 1D and 2D spectra. The flux calibration is done using the instrument response curve, derived from the observation of the standard star (see Moehler et al. 2014). The rms of the wavelength correction is 0.03 Å for the UVB and VIS regions, and 0.1 Å  for the NIR region. Since the spectrum for J1105 is an emission line spectrum (see below), we do not perform telluric corrections. We estimate the slit-losses assuming the slit transmission factor used by ESO111https://www.eso.org/observing/etc/doc/formulabook/node18.html and a constant seeing, which yields losses of 8$\%$, 4$\%$ and 2$\%$, in the UVB, VIS and NIR regions, respectively (see e.g. Spanò et al. 2006).

To calculate the flux, the full width half maximum (FWHM), velocity dispersion $\sigma$ and radial velocity of the emission lines in the three X-Shooter spectral regions, we fit Gaussian curves to them. To estimate the uncertainties of these fits we perform a Monte Carlo simulation. We assume that the noise of the spectrum is Gaussian with a mean equal to zero and a standard deviation equal to the value of the rms in each spectral region. We then make 1000 copies of the spectrum of the object, where each copy is the result of adding a value of this Gaussian distribution of the noise to every data point. Finally, we fit the emission lines in the resulting spectrum. The standard deviation of each parameter distribution is assumed to be its uncertainty. We add statistical uncertainties in quadrature, i.e. the rms of the wavelength correction plus the fit uncertainty. All values taken from the literature are presented with their published uncertainties.

3.2 $Chandra$ data

We calculate the column of neutral hydrogen $N_{H}$ and the X-ray flux by analizing the Chandra/ACIS archival observation ID 9552 of J1105. This observation was performed on 2008 January 28, with an exposure time of 71.5 ks. We reprocess the event files with the latest calibration files (CALDB version 4.8.1) using the CIAO 4.10 software (Fruscione et al., 2006). We extract a source spectrum with the ciao task specextract, using the X-ray position of J1105 (Liu, 2011). We then create a source region of 2″ centered on J1105, and a background region with with inner and outer radii of 10 and 30″, respectively, both centered on J1105. Using the ciao task pileup_map we derive that pile-up is less than 4% at the position of the source. We rebin the extracted source spectrum such that each channel has at least 30 X-ray photons. We fit the extracted spectra using the HEASOFT XSPEC tool version 12.10.1 (Arnaud, 1996). We exclude photons detected outside the range 0.3–7 keV as this energy interval is the best calibrated and most sensitive range for Chandra. We find that a fit-function consisting of an absorbed power law describes the data well (phabs $\times$ pegpwr in XSPEC; $\chi^{2}=51.45$ for 59 degrees of freedom, see Figure 2). The best–fitting power law index is 2.58 $\pm$ 0.09 for an N${}_{H}=($0.44 $\pm$ 0.04$)\times 10^{22}$ cm*-2*.

In order to test potential systematic effects related to the fit-function used to describe the X-ray spectrum, we tried a few other well motivated models such as a phabs $\times$ (diskbb + pegpwr) to describe the ULX states and a diskpbb to determine the range of values obtained for NH given these different continuum models. Typically, a value of 0.3$\times 10^{22}$ cm*-2* was obtained using these two fit-functions. Therefore, we conclude that the best-fit NH of 0.4$\times 10^{22}$ cm*-2* that was obtained from our simple absorbed power law model might have been on the high-end of possible values found when more complex fit-functions were used to describe the X-ray spectrum, and we add this systematic effect to have a final value of N${}_{H}=(0.44\pm 0.14)\times 10^{22}$ cm*-2*. The unabsorbed source 0.3–7 keV flux is then $($8.3 $\pm$ 1.0$)\times 10^{-13}$ erg cm*-2* s*-1*, which for the distance of NGC 3521, yields $L_{X}=(1.9\pm 0.8)\times 10^{40}$ erg s*-1*. In contrast, Liu (2011) reported a flux of 1.92$\times 10^{-13}$ erg cm*-2* s*-1*, which for the distance of NGC 3521, translates to $L_{X}=4.6\times 10^{39}$ erg s*-1*, a value almost four times lower than our estimate, using a phabs $\times$ pegpwr fit with a power law index of 1.7. It is worth noting that the analysis from Liu (2011) is based on the Chandra observation ID 4694, with an exposure time of 9.4 ks, executed on 2004 March 11. Since this observation was performed four years earlier than the observation we analyse in this manuscript, we can interpret this difference in X-ray luminosities as variability in the source, which is not uncommon in ULXs (see Liu et al. 2019 and references therein).

4 Results

The finding chart for J1105 is shown in Figure 1 where it can be seen that, due to the orientation of the slit, there are two sources inside it. The NIR counterpart identified in López et al. (2017) falls near the center of the slit, whereas a brighter extended source is located 2″(137.6 pc) away (hereafter source B). We extract the spectra from both these sources and analyse them below. In Figures 5 through 10 we provide plots of the regions around the most prominent emission lines.

There are several emission lines present in the spectra of both sources (see Table 2). We measure the radial velocity of the emission line region in the three X-Shooter spectral regions separately by fitting Gaussian curves to the lines detected. The average radial velocities found are 238 $\pm$ 10 km s*-1* for J1105 and 235 $\pm$ 4 km s*-1* for source B, which are both consistent with the radial velocity of NGC 3521 of 200 $\pm$ 25 km s*-1*, confirming that both objects are not foreground objects.

We calculate the intrinsic FWHM (FWHMint) of the emission lines. In order to do this, we estimate the instrumental FWHM (FWHMins), i.e. the width that a delta function line would have due to the instrument, given by the resolution of X-Shooter in the respective arms. As indicated in Section  3.1, the resolution R in the UVB, VIS and NIR arms is 6200, 11400 and 8100, respectively, and it is defined as R = $\frac{\lambda}{\Delta\lambda}$, where $\Delta\lambda$ is equivalent to FWHMins. Note that this value is wavelength dependent. We then calculate the FWHMint via the equation FWHM${}_{\rm int}=($FWHM${}_{\rm obs}^{2}-$ FWHM${}_{\rm ins}^{2})^{1/2}$, where FWHMobs is the observed FWHM, i.e. the FWHM we derive by fitting Gaussians to the emission lines. We present the FWHMint in Table 2.

4.1 Nebular source B

4.1.1 UVB spectra

Faint continuum emission is detected in the UVB and VIS regions only of the spectrum from source B. We find the $[$O ii$]\lambda\lambda$3727, 3729 doublet, H$\beta$, $[$Ne iii$]\lambda$3868 and $[$O iii$]\lambda\lambda$4956, 5007 lines. We also detect the $[$Fe i$]\lambda$3744 line, although the Gaussian fit to this feature did not converge at the position of source B. So, we deem this a marginal detection for source B.

We also find the $[$He i$]\lambda$3888, H$\delta$, H$\epsilon$ and H$\gamma$ lines. The He ii $\lambda$4686 emission line, observed in several ULX nebulae (e.g. Pakull & Mirioni 2002; Kaaret et al. 2004; Gutiérrez & Moon 2014) is not detected with a 2–$\sigma$ upper limit of 2.26$\times 10^{-19}$ erg cm*-2* s*-1*.

4.1.2 VIS spectra

We detect $[$O i$]\lambda$6300, $[$N ii$]\lambda$6548, H$\alpha$, $[$N ii$]\lambda$6583, the $[$S ii$]\lambda\lambda$6716, 6731 doublet, $[$S iii$]\lambda\lambda$9069, 9530 and Pa$\delta$ emission. Additionally, we detect emission from $[$He i$]\lambda$6678, $[$Ar iii$]\lambda\lambda$7135, 7751, $[$O ii$]\lambda\lambda\lambda$7320, 7329, 7330, Pa10, Pa9, Pa$\epsilon$ and the temperature sensitive auroral $[$S iii$]\lambda$6312. The latter allows us to calculate the electron temperature $T_{e}$ and density $n_{e}$, and ion abundances.

Additionally, we detected an emission line only in the spectrum from source B, at the wavelength 7325 Å , which we were not able to identify.

4.1.3 NIR spectra

We detect He i, 3He i, Pa$\beta$, Pa$\alpha$, Pa$\gamma$, $[$Fe ii$]\lambda$12566, $[$Fe ii$]\lambda$16435, $[$N iii$]$, Br19, Br16, Br13, Br12, Br11, Br10 and Br9 emission.

4.1.4 Emission line analysis

Using the results from the Gaussian fits to the emission lines, we calculate the line flux ratios $\log([$O iii$]/$H$\beta)$, $\log([$N ii$]/$H$\alpha)$, $\log([$S ii$]/$H$\alpha)$ and $\log([$O i$]/$H$\alpha)$. As can be seen in Figure 3, the flux line rations for source B are consistent with those of an Hii region (Kauffmann et al., 2003; Kewley et al., 2006). Moreover, in Figure  4 we show that the emission lines associated with source B are extended.

The $n_{e}$ for source B was calculated using both the $[$O ii$]$ and $[$S ii$]$ line ratios (Aller et al., 1949; Copetti & Writzl, 2002; Osterbrock & Ferland, 2006) using the TEMDEN task from the ANALYSIS package in Iraf. The calculations give $n_{e}[$O ii$]$ = 50 $\pm$ 15 cm*-3* and $n_{e}[$S ii$]$ = 62 $\pm$ 24 cm*-3*, consistent with each other. These $n_{e}$ values suggest that the Hii region associated with source B is a low-density region (Osterbrock & Ferland, 2006). With these values and the $[$S iii$]\lambda$9069, 9530/$[$S iii$]\lambda$6312 intensity ratio (e.g. Osterbrock & Ferland 2006) we calculated a value of $T_{e}[$S iii$]$ = 8130 $\pm$ 1213 K. This method is applicable to densities up to 105 cm*-3* (i.e. negligible collisional deexcitations) and assuming an isothermal environment (see Osterbrock & Ferland 2006).

We were also able to calculate the $[$O ii$]$, $[$O iii$]$, $[$S ii$]$ and $[$S iii$]$ ion abundances, following the equations from Pagel et al. (1992), adequate for low to moderate $n_{e}$ (i.e. up to 103 - 104 cm*-3*, Osterbrock & Ferland 2006). Additionally, we determined the empirical oxygen and sulfur abundance indicators, $R_{23}$ and $S_{23}$, respectively (Pagel et al., 1979; Vilchez & Esteban, 1996), the ionisation parameter u (Diaz et al., 1991) and the radiation softness parameter $\eta^{\prime}$ (Vilchez & Pagel, 1988). The latter is an indicator on the temperature of the ionizing stars in an Hii region. The derived values are given in Table 3.

4.2 The ULX counterpart

4.2.1 UVB spectra

We detect the $[$O ii$]\lambda\lambda$3727, 3729 doublet, $[$Fe i$]\lambda$3744, H$\beta$, $[$Ne iii$]\lambda$3868 and $[$O iii$]\lambda\lambda$4956, 5007 lines. Addtionally, we detect emission lines from $[$Fe ii$]\lambda\lambda$4436, 4452. For this source we also do not detect He ii $\lambda$4686 emission, with a 2–$\sigma$ upper limit of 2.12$\times 10^{-19}$ erg cm*-2* s*-1*.

4.2.2 VIS spectra

We detect $[$O i$]\lambda$6300, $[$N ii$]\lambda$6548, H$\alpha$, $[$N ii$]\lambda$6583, the $[$S ii$]\lambda\lambda$6716, 6731 doublet, $[$S iii$]\lambda\lambda$9069, 9530 and $[$Fe i$]\lambda$10048 emission. The auroral $[$S iii$]\lambda$6312 line was not detected at the position of J1105 with a 2–$\sigma$ upper limit of 2.58$\times 10^{-18}$ erg cm*-2* s*-1*. This implies that we can only place limits on the $T_{e}$ and $n_{e}$ parameters.

4.2.3 NIR spectra

The lines present in the NIR data for J1105 are He i, 3He i, Pa$\beta$, $[$Fe ii$]\lambda$16435 and Pa$\alpha$. Faint continuum emission is detected only for J1105. However, the spectrum is too faint to allow us determine if the continuum is stellar.

To verify whether the flux detected with X-Shooter is consistent with the apparent magnitudes given by López et al. (2017) in the $H$–band, we estimate the apparent magnitudes from the flux calibrated spectra using the equation:

[TABLE]

where $m_{H}$ is the apparent magnitude in the $H$-band, $F_{H}$ is the flux of the emission lines detected in the wavelength range $1500-1750$ nm (corresponding to the $H$-band) and $F_{H,0}$ is the reference flux ($1.883\pm 0.038\times 10^{-6}$ erg cm*-2* s*-1*, Cohen et al. 2003). We assume $R_{V}=3.1$ (Fitzpatrick, 1999a) and $R_{H}=0.46\pm 0.01$ (Yuan et al., 2013) to correct the observed flux for extinction.

The total integrated flux of J1105 in the $H$–band is 1.78$\times 10^{-15}$ erg cm*-2* s*-1*, which is equivalent to an apparent magnitude of 22.54 $\pm$ 2.01, consistent within 2-$\sigma$ with our reported value of $H$ = 19.83 $\pm$ 0.06 mag (López et al., 2017).

4.2.4 Emission line analysis

We also calculate the line flux ratios $\log([$O iii$]/$H$\beta)$, $\log([$N ii$]/$H$\alpha)$, $\log([$S ii$]/$H$\alpha)$ and $\log([$O i$]/$H$\alpha)$ for J1105 and, as can be seen in Figure 3, they are consistent with those of Low-ionisation nuclear emission-line regions (LINERs) and Active Galactic Nuclei (AGNs) (Kauffmann et al., 2003; Kewley et al., 2006).

We were not able to calculate ion abundance for J1105, as they are dependent on $T_{e}$-, $n_{e}$- and $[$S iii$]\lambda$6312, and the latter was not detected for J1105. We can, however, obtain a lower limit on the temperature at the location of J1105, assuming that it is a photoionised region, like source B. We follow the same procedure as for source B but we found that the values for the $[$O ii$]$ and $[$S ii$]$ line ratios are at the limit of $n_{e}\rightarrow 0$. In fact, the $[$S ii$]$ line ratio is outside the range allowed by photoionisation models. We derive a value for $n_{e}[$O ii$]$ of 5 $\pm$ 5 cm*-3*, that, combined with the upper limit of the $[$S iii$]\lambda$9069, 9530/$[$S iii$]\lambda$6312 intensity ratio, yields $T_{e}[$S iii$]>$ 5270 K.

5 Discussion

5.1 The ULX counterpart

López et al. (2017) identified a possible counterpart to J1105 as a potential red supergiant based on its photometry. However, using VLT/X-Shooter spectroscopy we find that it has an emission line spectrum, whose line ratios place it in the LINER region of the Baldwin-Phillips-Terlevich (BPT) diagram. These line ratios are different from those of a nearby region called source B (see Figure 3), and could be interpreted as due to the X-ray emission ionising the part of the nebular source B that surrounds the ULX. The effects of the bright ULX on its environment has been seen in other Hii regions surrounding (or close to) a ULX (e.g. Mucciarelli et al. 2005; Heida et al. 2015). Since the presence of the ULX seems to affect the line flux ratios of J1105, and its radial velocity is consistent with the radial velocity of NGC 3521 at that position, we deem it likely that the ULX and the nebular source are physically related. Therefore, we refer to J1105 as the counterpart of the ULX although the mass donor star has not been detected so far.

We calculate the extinction at the position of J1105 using the NH we derived by analising the Chandra archival observations, with observation ID 9552, of the ULX. We obtained N${}_{H}=(0.44\pm 0.14)\times 10^{22}$ atoms cm*-2*, which for a typical value of $R_{V}=3.1$, translates to $E(B-V)=0.76\pm 0.24$. If we deredden the H$\alpha$ and H$\beta$ line fluxes with this extinction and then calculate the Balmer decrement, we get unphysical values, i.e. H${\alpha}/$H${\beta}<2$ (values only seen in environments with $n_{e}>10^{9}$ cm*-3*, Adams & Petrosian 1974). This can be explained if the X-ray emission from the ULX and the emission from the ionised gas comes from different regions, implying that there is more neutral gas in the line of sight towards the X-ray source than towards the gas responsible for the H$\alpha$ and H$\beta$ emission lines.

In light of this, we decide to calculate the extinction using the Balmer decrement for J1105 (Table 4). Assuming a standard Milky Way gas-to-dust ratio ($R_{V}=3.1$, Fitzpatrick 1999b) and an intrinsic Balmer decrement of H${\alpha}/$H${\beta}=3.05$ (valid for a typical optically thick nebula in the low density limit, with $T_{e}\sim$ 5,000 K), we derive $E(B-V)=0.27\pm 0.03$. The NH corresponding to this extinction is N${}_{H}=(0.15\pm 0.02)\times 10^{22}$ atoms cm*-2*. We then deredden the key fluxes with this two values of extinction and report them in Table 4, where it can be seen that they still place J1105 in the LINER region of the BPT diagram.

From the abundances derived for the extended source B nearby (Table 3), we conclude that J1105 resides in a low-metallicity environment. This is in line with previous findings that ULXs are preferentially found in low-metallicity environments (e.g. Kaaret et al. 2011; Walton et al. 2011; Prestwich et al. 2013; Basu-Zych et al. 2016). As stated above, the dereddened line flux ratios of J1105 place it in the LINER region of the BPT diagram. Though the power source behind the emission of a LINER is under debate, several explanations have been considered, e.g. photoionisation by hot young stars (Terlevich & Melnick, 1985; Shields, 1992), low-luminosity active galactic nuclei (Halpern & Steiner, 1983; Ferland & Netzer, 1983), planetary nebulae (Taniguchi et al., 2000) or exposed cores of evolved stars (Singh et al., 2013; Belfiore et al., 2016), or shock-ionization (Heckman, 1980; Dopita et al., 1996) We further investigate whether the latter is the case for J1105. We first analyse the $[$S ii$]$/H$\alpha$ ratio, which distinguishes between shock ionised and photoionised regions: shock ionised regions have $[$S ii$]$/H$\alpha>$ 0.4 (Dodorico et al., 1980). The dereddened ratio for source B is 0.17, a typical value for Hii regions (Levenson et al., 1995; Payne et al., 2008), whereas J1105 has a dereddened ratio of 0.64. Moreover, in the line flux ratio diagrams often used to distinguish between supernova remnants with shocks and Hii regions by Levenson et al. (1995), J1105 falls in the supernova remnants region while source B is consistent with being an Hii region (Table 4). Another clue on whether J1105 is shock ionised could come from the $[$O ii$]\lambda 3727$/H$\beta$, $[$N ii$]\lambda 6584$/H$\beta$, and $[$S ii$]\lambda 6731$/H$\beta$ ratios. Our dereddened values (Table 4) are consistent with shock velocities between 50 and 80 km s*-1* (Dopita et al., 1984). When we compare with the FWHMs from J1105 (see Table 2), we find that the average FWHM for this source is 50 $\pm$ 1 km s*-1*, whereas for source B is 45 $\pm$ 1 km s*-1*. Hence, our results are consistent with the diagrams from Dopita et al. (1984).

We estimate the mechanical power $P_{{\rm jet}}$ of J1105. According to Weaver et al. (1977) and Cseh et al. (2012), for a nebula which is shock ionised, $P_{{\rm jet}}\sim 2.85L_{{\rm rad}}$, where $P_{{\rm jet}}$ is the power that inflates the nebula bubble and $L_{{\rm rad}}$ is the radiative luminosity of the bubble. Following Dopita & Sutherland (1996); Abolmasov et al. (2007b); Allen et al. (2008) and Cseh et al. (2012) we know that $L_{{\rm H}\beta}=6.53\times 10^{-3}v_{100}^{-0.59}L_{{\rm rad}}$ where $L_{{\rm H}\beta}$ is the luminosity from H$\beta$ emission and $v_{100}$ is the shock velocity in units of 100 km s*-1*. This means that $P_{{\rm jet}}=437v_{100}^{0.59}L_{{\rm H}\beta}$. To calculate $L_{{\rm H}\beta}$, we deredden our $H\beta$ flux using $E(B-V)=0.27\pm 0.03$, which at the distance of NGC 3521, is equivalent to $L_{{\rm H}\beta}=5.25\times 10^{36}$ erg s*-1*. Hence, for shock velocities between 50 and 80 km s*-1*, $P_{{\rm jet}}$ is $(1.52-2.01)\times 10^{39}$, i.e. almost one-tenth of our derived X-ray luminosity of the ULX and 0.3–0.4 times the X-ray luminosity derived by Liu (2011). In comparison, for ULX IC 342-X1, $P_{{\rm jet}}\sim\frac{1}{20}L_{X}$ (Cseh et al., 2012), for M51 ULX-1, $P_{{\rm jet}}\sim L_{X}$ (Urquhart et al., 2018) and for S26, $P_{{\rm jet}}\sim 10^{4}L_{X}$ (Pakull et al., 2010).

As we have seen in other nebulae surrounding ULXs (e.g. Lehmann et al. 2005; Fabrika et al. 2006; Abolmasov et al. 2007a, b), both photoionised and shock ionised gas is present. We consider the possibility that shocks are not the only ionisation source and the X-ray emission from J1105 is photoionising the region as well. To investigate this scenario further, we calculate the Strömgren radius $R_{S}$ (Strömgren, 1939), i.e. the radius of the photoionised region, if any, expected to be created by J1105 (see e.g. Roberts et al. 2003). First we need to calculate the hydrogen ionising flux $Q$, which correlates with $L_{{\rm H}\beta}$ as $L_{{\rm H}\beta}=4.8\times 10^{-13}Q$ erg $s^{-1}$ (Osterbrock & Ferland, 2006). This gives a value of $Q=1.1\times 10^{49}$ photons s*-1*, and assuming a filling factor $\epsilon$ between 0.01 and 0.5, a typical range for most Hii regions (Osterbrock & Ferland, 2006) and $T_{e}=10000$ K, we get 30 $<R_{S}<$ 111 pc. In comparison, assuming spherical symmetry, we estimate a radius of $\sim 34$ pc for J1105 from its appearance in our $H$-band image. So, for a $\epsilon<0.35$, it would be possible for J1105 to be photoionising the entirety of the region. However, we do not need the entirety of the region to be photoionised, since we are inferring that the total ionisation is a combination of X-ray photoionisation and shock ionisation. Alternatively, a wind blown off the accretion disc of the ULX can also cause shock ionisation (Abolmasov et al., 2007a; Kaaret et al., 2017). In fact, powerful disc winds have been observed in a few ULXs (e.g. Walton et al. 2016; Pinto et al. 2016; Pinto et al. 2017; Kosec et al. 2018).

5.2 Nebular source B

The nebular ion abundances for source B are consistent with those seen in other Hii regions (see e.g. Pérez-Montero et al. 2003; Pérez-Montero & Díaz 2005). The $n_{e}$ value is at the low density end, whereas typical $n_{e}$ in Hii regions are of the order of 102 cm*-3* (Osterbrock & Ferland, 2006). The density that we obtain, within uncertainties, is seen in nebulae like NGC 281 and NGC 7000 (Roger & Pedlar, 1981; Osterbrock & Ferland, 2006).

Moustakas et al. (2010) studied the metallicity of NGC 3521 in an area enclosing $\sim$ 49$\%$ of the integrated light, up to 3.9 kpc from the nucleus. The oxygen abundance they calculated is 12 + $\log$(O/H) = 9.08 $\pm$ 0.06, which could be consistent with our value of 8.19 $\pm$ 0.11, since metallicity in spiral galaxies decreases with distance to the nucleus (Pagel & Edmunds, 1981; Deharveng et al., 2000; Magrini et al., 2016), and source B is located 10.5 kpc from the nucleus. This is further confirmed with the $\log R_{23}$ parameter, i.e. the oxygen vs. H$\beta$ abundance indicator. Moustakas et al. (2010) get $\log R_{23}=0.32\pm 0.20$ and we get $\log R_{23}=0.52\pm 0.03$, indicative of a decrease in the oxygen abundance (see Pérez-Montero & Díaz 2005). The value we calculate for $\log S_{23}$ (the sulfur vs. H$\beta$ abundance indicator) is also consistent with our value for 12 + $\log$ (O/H). Coccato et al. (2018) measured the metallicity of NGC 3521 between its nucleus and 30″ from the nucleus. At the latter position, they get a value of [Z/H] = $-$0.3. In comparison, we calculate values of [O/H] = $-$0.65 and [S/H] = $-$0.31222Using $\log O_{\odot}=8.83$ and $\log S_{\odot}=7.27$ (e.g. Bowey 2003)., consistent with a metallicity decrease.

Moustakas et al. (2010) also calculated the ionisation parameter for NGC 3521 (in an area enclosing 49$\%$ of the integrated light), getting a value of $\log U=-2.89\pm 0.25$, consistent with our measurement, since $\log U$ was found to maintain a practically constant value with distance to the nucleus in some galaxies (i.e. M33, Vilchez et al. 1988). Additionally, they found a Balmer decrement of H${\alpha}/$H${\beta}=5.78\pm 0.50$, whereas we found a value of 3.88, indicative of lower extinction (Mathis, 1983; Diaz et al., 1991; Pérez-Montero & Díaz, 2005) at the position of source B than at 3.9 kpc from the nucleus and consistent with the extinction decreasing as a function of distance to the nucleus of a galaxy (Cikota et al., 2016).

With the Balmer decrement we could estimate the extinction at the position of source B, assuming a standard Milky Way gas-to-dust ratio ($R_{V}=3.1$, Fitzpatrick 1999b). We assume an intrinsic H${\alpha}/$H${\beta}=2.87$, valid for a typical optically thick nebula (recombination case B, low density limit, and $T_{e}\sim$ 10,000 K Osterbrock & Ferland 2006) and we obtain $E(B-V)=0.28\pm 0.01$. We deredden the emission line fluxes from source B and calculate again the line ratios, finding that they are still consistent with source B being an H ii region (see Table 4).

The $T_{e}$ we find for the Hii region source B is also within the range of observed temperatures in Hii regions, i.e. 7000–10000 K (e.g. Osterbrock et al. 1992; Garnett et al. 1997; Esteban et al. 1998; Kennicutt et al. 2003). Lastly, our value of the radiation softness parameter $\eta^{\prime}$ along with 12 + $\log$ (O/H) correspond to an Hii region ionised by a young hot stellar population, with an effective temperature $T_{eff}=40,000K$ (for $\log U=-2.6$, Bresolin et al. 1999; Viegas et al. 2000). This is in line with previous findings that ULXs often reside close to young OB associations (e.g. Grisé et al. 2008; Grisé et al. 2011; Voss et al. 2011; Jonker et al. 2012).

6 Conclusions

We took X-Shooter spectra of a candidate NIR counterpart to the ULX J1105 in NGC 3521. We obtained spectra of both the counterpart of the ULX and a nearby HII region. We detect several emission lines from these two spectra. The line flux ratios from the extended source B classify it as an Hii region, while the line flux ratios from J1105 place it in the LINER region of the BPT diagram. Source B has the characteristics of a typical low density Hii region, with a $T_{e}=8130\pm 1213$ K and $n_{e}=50\pm 15$ cm*-3*. It is a low metallicity region, i.e. 12 + $\log$ (O/H) = 8.19 $\pm$ 0.11. This is consistent with the fact that several ULXs are found primarily in low metallicity environments and close to young OB associations. J1105 is a different object. The fact that the line ratios at the position of this source are different than source B is interpreted as being caused by the ULX. So we establish that the ULX is located in NGC 3521 although we do not detect the mass donor star. We calculate $L_{X}$ from the $Chandra$ observation of the ULX and derive a value significantly higher than the one reported before by Liu (2011), implying source variability. Additionally, we derive a lower limit of $T_{e}\sim 5270$ K for J1105 and its line ratios indicate that this source could be shock ionised. The diagnostic diagrams from Dopita et al. (1984) suggest that the shocks have low velocity, i.e. between 50 and 80 km s*-1*. This is further confirmed by the average FWHM from J1105 (50 $\pm$ 1 km s*-1*). However, the luminosity of the H$\beta$ line seems to indicate that the ULX is photoionising the nebula. Hence, we argue that the nebula is not only shock ionised but also photoionised by the X-ray radiation from the ULX. The shock ionisation could be driven by either a disc wind or a jet.

Acknowledgements

We would like to thank the anonymous referee whose valuable comments helped improve the quality of this manuscript. KML would like to thank Davide Lena for valuable discussions. PGJ and KML acknowledge funding from the European Research Council under ERC Consolidator Grant agreement no 647208. TPR acknowledges funding from STFC as part of the consolidated grant ST/L00075X/1. MAPT acknowledges support via a Ramón y Cajal Fellowship (RYC-2015-17854). MAPT also acknowledges support by the Spanish Ministry of Economy, Industry and Competitiveness under grant AYA2017-83216-P. DJW acknowledges support from an STFC Ernest Rutherford Fellowship. KML We have made use of the SIMBAD database, operated at CDS, Strasbourg, France; of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration; and of data obtained from the Chandra Data Archive and the Chandra Source Catalog, and software provided by the Chandra X-ray Center (CXC) in the application packages CIAO, ChIPS, and Sherpa.

Appendix A Spectra