The Origin of the Vanishing Soft X-Ray Excess in the Changing-look Active Galactic Nucleus Mrk 590

XMM-Newton observed Mrk 590 in 2002 January 1 and then in 2004 July 4. The details of the observations and the short IDs are mentioned in Table 1. Archival data from the EPIC, RGS, and OM instruments are available. We preferred the EPIC-pn (Strüder et al. 2001) over MOS data owing to their better signal-to-noise ratio, which is critical for the broadband spectral study of our source. For both observations (obs1 and obs2), the EPIC-pn camera operated in the small-window mode. The EPIC-pn data were reprocessed with V18.0.0 of the Science Analysis Software (SAS; Gabriel et al. 2004) using the task epchain. We created the filtered event list after screening for flaring background due to high-energy particles. Circular regions of 40'' centered on the centroid of the source, were used to extract the source counts, whereas a 40'' circular region away from the source but located on the same CCD was selected to estimate the background counts. The SAS task epatplot was used to estimate the pileup in our observations. We found that both obs1 and obs2 are free of any pileup. The corresponding response matrix function (RMF) and auxiliary response function (ARF) for each observation were created employing the SAS tasks arfgen and rmfgen. We used the command specgroup to group the XMM-Newton spectra by a minimum of 20 counts per channel and a maximum of three resolution elements required for the χ² minimization technique. The task omichain was used to reduce the data from the Optical Monitor (OM) for the six active filters (V, B, U, UVW1, UVM2, and UVW2). We used the task om2pha to create the necessary files to be analyzed with XSPEC, together with the simultaneous X-ray data. We corrected the observed UV fluxes for the Galactic reddening assuming (Fitzpatrick 1999) a reddening law with R _v = 3.1. We fixed the color excess parameter of the reddened component at E(B − V) = 0.0306, and the Galactic extinction coefficient value used was 0.257 (Mrk 590).

Suzaku started observing Mrk 590 in 2011 January 23; however, it was interrupted owing to a Target of Opportunity trigger. The observation continued on 2011 January 26, making them two separate observations (eee Table 1). Suzaku has three X-ray Imaging Spectrometers (XISs; Koyama et al. 2007) and the Hard X-ray Detector (HXD; Takahashi et al. 2007) on board that cover a broad energy band of 0.2–50 keV. There are, however, no simultaneous optical–UV observations. In both observations (obs3 and obs4), the XIS data were obtained in both 3 × 3 and 5 × 5 data modes and XIS nominal position. We reprocessed the data using the Suzaku pipeline with the screening criteria recommended in the Suzaku Data Reduction Guide. All extractions were done using HEASOFT (V6.27.2) software and the recent calibration les. For HXD/PIN, which is a nonimaging instrument, we used appropriately tuned background les provided by the Suzaku team and available at the HEASARC website. We co-added the spectral data from the front-illuminated XIS instruments to enhance the signal-to-noise ratio. We used the tool grppha to group the XIS spectral data from both the observations to a minimum of 100 counts in each energy bin. We also grouped the HXD/PIN data to produce ∼60 energy bins with more than 20 counts per bin.

We have used four quasi-simultaneous NuSTAR and Swift observations of Mrk 590 (obs5, obs6, obs7, and obs8). See Table 1 for details. These are all the currently available archival data that are free of any technical issues reported in the NuSTAR Master Catalog. Years 2016 and 2021 both have two NuSTAR observations that have simultaneous Swift data. We selected those two NuSTAR observations, one each from 2016 and 2021, that have highest exposure in the Swift XRT instrument. This is essential for detecting the soft excess in the X-ray band. We reprocessed the NuSTAR FPM (Harrison et al.2013) and Swift XRT (Burrows et al. 2004) plus UVOT (Roming et al. 2005) data. For NuSTAR we produced the cleaned event files using the standard NUPIPELINE (v2.0.0) command, part of the HEASOFT V6.28 package, and instrumental responses from NuSTAR CALDB version V20210202. For light curves and spectra, we used a circular extraction region of 80'' centered on the source position and a 100'' radius for background, respectively. The NuSTAR spectra were grouped to a minimum count of 20 per energy bin, using the command grppha in the HEASOFT software.

The Swift XRT and UVOT data reprocessing and spectral extraction were carried out following the steps described in Ghosh et al. (2016, 2018). All the Swift XRT spectra were grouped by a minimum of 20 counts per channel. In both observations, Swift UVOT observed the source in all six filters, i.e., in the optical (V, B, U) bands and the near-UV (W1, M2, W2) bands. We used the UVOT2PHA tool to create the source and background spectra and used the response files provided by the Swift team.

We began our spectral fitting of the source spectra with a set of phenomenological models. This exercise helps us to characterize the source spectra at different epochs and determine the spectral features quantitatively. For Mrk 590, we used a power law representing the primary continuum emission, diskbb to model the soft excess, zgauss for the Fe line emission, and pexrav with a negative reflection fraction to model the Compton hump. The power-law, diskbb, and Gaussian model components are left free to vary during the spectral fitting of each observation to check the variability of the continuum and discrete spectral properties of the source. We fixed all pexrav model components except for the reflection fraction. We tied the photon index and model normalization of pexrav with that of the primary continuum. The abundance of iron and other heavier elements than He was fixed to solar values. We made the inclination of pexrav a free parameter for obs3 only. This particular observation was selected owing to its longest exposure among the Suzaku and NuSTAR data sets. We used this best-fit inclination value and fixed it for all other observations. For all the model components, we also report the improvement in the χ² values that indicates how significant these components are in the spectral fit. A constant component was added to take into account the relative normalization of Suzaku and NuSTAR instruments. In XSPEC the model reads as constant × tbabs × (po+diskbb+zgauss+pexrav). Below we discuss the different epochs of observations using different telescopes.

3.1.1. XMM-Newton

For obs1 and obs2, the two XMM-Newton observations, fitting the 2–5 keV energy band with an absorbed power-law model and extrapolating to the rest of the X-ray band, revealed a prominent soft excess below 2 keV (see Appendix A for details). The addition of the diskbb component improved the fit statistics considerably (Δχ² ∼ 20) for both these observations. The best-fit value of inner disk temperature is consistent with a best-fit value of ${0.23}_{-0.05}^{+0.05}\,\mathrm{keV}$ and ${0.20}_{-0.04}^{+0.03}\,\mathrm{keV}$ for obs1 and obs2, respectively. Next, we added a Gaussian component to the best-fit model. We could not constrain the line width σ for the Fe K line emission. We obtained an upper limit of 0.10 and 0.06 keV for obs1 and obs2, respectively. The best-fit power-law photon index values were consistent between observations. All best-fit parameter values, along with their fit statistics, are quoted in Table 2.

Table 2. Best-fit Parameters of the Baseline Phenomenological Models for the Observations of Mrk 590

Models	Parameter	obs1	obs2	obs3	obs4	obs5	obs6	obs7	obs8
Gal. abs.	NH (× 1020 cm−2)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)
power-law	Γ	${1.79}_{-0.03}^{+0.07}$	${1.76}_{-0.04}^{+0.04}$	${1.70}_{-0.01}^{+0.01}$	${1.70}_{-0.01}^{+0.01}$	${1.64}_{-0.09}^{+0.10}$	${1.66}_{-0.09}^{+0.10}$	${1.64}_{-0.04}^{+0.05}$	${1.61}_{-0.05}^{+0.05}$
	norm (10−3)	${1.19}_{-0.12}^{+0.09}$	${1.56}_{-0.06}^{+0.06}$	${1.78}_{-0.02}^{+0.02}$	${1.59}_{-0.02}^{+0.02}$	${0.62}_{-0.08}^{+0.10}$	${0.65}_{-0.09}^{+0.10}$	${2.86}_{-0.16}^{+0.16}$	${1.38}_{-0.09}^{+0.10}$
	ECutoff (keV)	⋯	⋯	⋯	⋯	⋯	⋯	>95	>79
diskbb	Tin (keV)	${0.23}_{-0.05}^{+0.05}$	${0.20}_{-0.04}^{+0.03}$	0.20(f)	0.20(f)	0.20(f)	0.20(f)	0.20(f)	0.20(f)
	norm	${22.4}_{-7.6}^{+7.1}$	${28.5}_{-10.7}^{+3.7}$	${7.1}_{-5.6}^{+5.5}$	${9.9}_{-6.6}^{+5.2}$	<11.9	<19.5	<5.1	<25.8
	a Δχ2/dof	21/2	26/2	4/1	9/1	1/1	1/1	1/1	3/1
Gaussian	E (keV)	${6.39}_{-0.06}^{+0.05}$	${6.41}_{-0.03}^{+0.03}$	${6.42}_{-0.03}^{+0.03}$	${6.42}_{-0.04}^{+0.05}$	${6.48}_{-0.16}^{+0.21}$	${6.48}_{-0.16}^{+0.22}$	${6.35}_{-0.08}^{+0.17}$	${6.32}_{-0.08}^{+0.08}$
	σ (keV)	<0.10	<0.06	<0.12	${0.09}_{-0.07}^{+0.09}$	<0.45	<0.44	${0.24}_{-0.16}^{+0.33}$	${0.19}_{-0.10}^{+0.14}$
	norm (10−5)	${1.05}_{-0.46}^{+0.51}$	${0.78}_{-0.24}^{+0.23}$	${1.23}_{-0.25}^{+0.26}$	${1.02}_{-0.30}^{+0.37}$	${0.73}_{-0.39}^{+0.46}$	${0.72}_{-0.39}^{+0.45}$	${2.61}_{-0.71}^{+1.35}$	${1.68}_{-0.43}^{+0.49}$
	a Δχ2/dof	17/3	33/3	99/3	53/3	14/3	15/3	75/3	54/3
Pexrav b	R	<0.57	0.31 (t)	<0.24	${0.35}_{-0.30}^{+0.35}$	<0.44	<0.53	⋯	⋯
	Incl (deg)	12(t)	12 (t)	12 c	12(t)	12(t)	12(t)	⋯	⋯
	a Δχ2/dof	4/2	7/2	2/2	4/2	3/2	4/2	⋯	⋯
Gaussian	EqW (eV)	174	123	150	132	227	169	182	229
χ2/dof		104/122	201/161	907/831	555/549	193/208	304/313	878/837	572/613

Notes.

^a The Δχ² improvement in statistics upon addition of the corresponding discrete component. ^b The model pexrav was used only for Suzaku and NuSTAR observation, as it had broadband spectra necessary for constraining the parameters. The values quoted for the XMM-Newton observations are from the simultaneous fit of all the data sets. R represents the reflection component only. The temperature at inner disk radius T_in (keV) for Suzaku and Swift + NuSTAR observations, when left free, was taking very low values and hence was fixed at 0.2 keV. ^c Indicates that parameters are not constrained.

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3.1.2. Suzaku

3.1.3. NuSTAR and Swift

3.1.4. Summary of Results

3.2.1. Ionized Disk Reflection

We used the relxill model, version 1.4.0 (García et al. 2014; Dauser et al. 2014), in our spectral fitting. This model assumes the origin of soft excess to be relativistic reflection from an ionized accretion disk or simply ionized reflection. We added the MyTorus model (Yaqoob et al. 2010; Yaqoob & Murphy 2011) to take into account the distant neutral reflection from the outer part of the disk or torus. In XSPEC, our model reads as constant×tbabs×(relxill+MyTorus). The relxill model describes the soft excess emission, the X-ray continuum, and the broad Fe Kα emission line. The distant neutral reflection, on the other hand, is modeled with the two MyTorus model components, first MyTorusL, which describes the iron Fe Kα and Kβ lines, and second MyTorusS, which models the scattered emission due to the reflection of primary power-law emission from the torus. The best-fit parameters of model components for all observations are quoted in Table 3, along with their χ² fit statistic.

Table 3. Best-fit Parameters for Observations of Mrk 590 with the First Set of Physical Models

Component	Parameter	obs1	obs2	obs3	obs4	obs5	obs6	obs7	obs8
Gal. abs.	NH (1020 cm−2)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)
relxill	AFe	1 (f)	1 (f)	1(f)	1 (f)	1(f)	1(f)	1(f)	1(f)
	$\mathrm{log}\xi \,(\mathrm{erg}\,\mathrm{cm}\,{{\rm{s}}}^{-1})$	<2.01	${0.52}_{-0.30}^{+0.77}$	${3.19}_{-0.84}^{+0.45}$	${2.72}_{-0.68}^{+0.18}$	${3.19}_{-0.47}^{+0.49}$	${3.30}_{-1.45}^{+0.91}$	${3.11}_{-0.80}^{+0.46}$	${3.30}_{-0.99}^{+0.55}$
	Γ	${1.88}_{-0.08}^{+0.02}$	${1.81}_{-0.05}^{+0.02}$	${1.68}_{-0.01}^{+0.02}$	${1.65}_{-0.01}^{+0.06}$	${1.62}_{-0.06}^{+0.10}$	${1.60}_{-0.03}^{+0.04}$	${1.60}_{-0.06}^{+0.06}$	${1.58}_{-0.03}^{+0.02}$
	Ecut (keV)	300 (f)	300(f)	>32	>73	>28	>65	${92}_{-25}^{+55}$	${60}_{-8}^{+10}$
	Ecut (keV)	300 (f)	300(f)	55(f)	90 (f)	52 (f)	70 (f)	${92}_{-25}^{+55}$	${60}_{-8}^{+10}$
	${n}_{\mathrm{rel}}{\,({10}^{-5})}^{a}$	${2.44}_{-0.12}^{+0.38}$	${3.71}_{-0.18}^{+0.55}$	${5.32}_{-0.75}^{+0.28}$	${5.28}_{-0.17}^{+0.17}$	${2.17}_{-0.43}^{+0.45}$	${7.27}_{-0.81}^{+0.97}$	${6.20}_{-1.20}^{+1.32}$	${3.11}_{-1.81}^{+1.11}$
	q1	3(f)	3(f)	3(f)	3(f)	3(f)	3(f)	3(f)	3(f)
	a	0(f)	0(f)	0(f)	0(f)	0(f)	0(f)	0(f)	0(f)
	Rfrac	${0.47}_{-0.40}^{+0.34}$	${0.46}_{-0.26}^{+0.16}$	${0.04}_{-0.02}^{+0.06}$	${0.08}_{-0.04}^{+0.05}$	${0.14}_{-0.10}^{+0.28}$	<0.14	${0.24}_{-0.13}^{+0.27}$	${0.24}_{-0.12}^{+0.32}$
	Rin (rg )	6(f)	6(f)	6(f)	6(f)	6(f)	6(f)	6(f)	6(f)
	Rbr (rg )	10(f)	10(f)	10(f)	10(f)	10(f)	10(f)	10(f)	10(f)
	Rout (rg )	400 (f)	400(f)	400(f)	400(f)	400(f)	400(f)	400(f)	400(f)
	i (deg)	45(f)	${45}_{-7}^{+5}$	45(f)	45(f)	45(f)	45(f)	45(f)	45(f)
MYTorusL	i (deg)	45(f)	45(f)	45(f)	45(f)	45(f)	45(f)	45(f)	45(f)
	norm (10−3)	${6.88}_{-2.69}^{+2.71}$	${4.07}_{-1.26}^{+1.27}$	${4.09}_{-0.70}^{+0.65}$	${2.80}_{-0.76}^{+0.70}$	${1.46}_{-0.82}^{+1.02}$	${4.58}_{-1.82}^{+2.03}$	${4.69}_{-1.52}^{+1.70}$	${3.01}_{-1.05}^{+1.33}$
MYTorusS	NH (1024 cm−2)	10(f)	10(f)	10(f)	10(f)	10(f)	10(t)	10(t)	10(t)
	norm (10−3)	6.88(f)	4.07(f)	<0.50	<0.62	<0.55	<0.75	<0.71	<0.50
	χ2/dof	108/127	206/163	899/833	551/550	191/209	302/314	867/836	578/612
With OM data
diskbb	Tin (eV)	${1.25}_{-0.04}^{+0.04}$	${0.94}_{-0.02}^{+0.02}$	⋯	⋯	${1.08}_{-0.04}^{+0.04}$	${1.75}_{-0.07}^{+0.09}$	${1.62}_{-0.07}^{+0.08}$	${1.41}_{-0.05}^{+0.06}$
	norm ( × 1012)	${0.30}_{-0.05}^{+0.06}$	${1.59}_{-0.16}^{+0.18}$	⋯	⋯	${0.60}_{-0.12}^{+0.15}$	${0.13}_{-0.03}^{+0.04}$	${0.15}_{-0.03}^{+0.04}$	${0.20}_{-0.04}^{+0.05}$
	χ2/dof	164/134	222/171	⋯	⋯	230/214	344/317	910/845	580/621

Note. In XSPEC, the models read as (constant × tbabs × (relxill+MYTorus)). Spectral fitting of all observations includes simultaneous optical–UV data except obs3 and obs4. (f) indicates a frozen parameter. (t) indicates a tied parameter between observations. (a) n_rel represents normalization for the model relxill.

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The relxill model assumes a lamp-post geometry of the corona where part of the hard X-ray continuum enters the accretion disk, ionizes it, and emits fluorescence lines. These emission lines then get blurred and distorted owing to the extreme gravity around the central SMBH and, along with scattered emission from the ionized accretion disk, produce the soft excess emission and a broad Fe emission line around 6 keV. The transition between this relativistic and Newtonian geometry is characterized by a breaking radius r_br. We fix the emissivity index of reflection from the disk outside this r_br (q2) at 3, as expected for a point source under Newtonian geometry at a large distance from the source, whereas the emissivity inside the radius r_br, q1, was allowed to vary, as this region falls under the relativistic high-gravity regime and previous studies (Dabrowski & Lasenby 2001; Miniutti et al. 2003; Wilkins & Fabian 2011) suggest a very steeply falling profile in the inner parts of the disk. In our spectral fits with phenomenological models we did not detect any broad Fe emission line, which would indicate a rotating black hole. To confirm the nondetection, we tested two extreme scenarios. First, we fixed the black hole spin parameter to a maximum value of 0.998 and the inner radius (r_in) to 1.24r _g , the lowest value allowed in the model, and in the second scenario, we fixed the spin to zero and inner radius to 6r _g . We found that the fit statistics is insensitive to both rotating and nonrotating scenarios for all our observations. Since the data are insensitive to the spin of the black hole, we continued with the nonrotating black hole scenario in all the spectral fits in our work. Accordingly, we fixed r_br to a larger value of 10r _g and q1, the emissivity index inside r_br, to 3.

The availability of hard X-ray data beyond 10 keV for obs3 to obs8 involving Suzaku, Swift, and NuSTAR provided us an opportunity to measure the high-energy cutoff in this CLAGN. We made the parameter E_cut free in all our fits. However, we could not constrain the value of E_cut for obs3 to obs6. Interestingly, we get a well-constrained high-energy cutoff of ${92}_{-25}^{+65}\,\mathrm{keV}$ and ${60}_{-08}^{+10}\,\mathrm{keV}$ for obs7 and obs8, respectively. To robustly identify the errors on the energy cutoff, we carried out some statistical tests. Using the steppar command in XSPEC, we determined the confidence intervals for the parameter. In the top panel of Figure 1 we show the example of obs4 and obs6, where E_cut could not be constrained. The bottom panel shows the confidence intervals of obs7 and obs8, where E_cut is well constrained. Obs1 and obs2 involving XMM-Newton do not cover the energy band above 10 keV. We tested with different E_cut values (60, 90, and 300 keV) and noted that the fit is insensitive to the parameter, and hence we fixed the parameter to 300 keV for these two observations.

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The iron abundance of the material in the accretion disk is represented by the parameter A_Fe. We first tied this parameter between obs1 and obs2 and only got an upper limit ( <0.78); hence, we fixed it to solar values. For other observations, when left free, we found it to be pegged at the highest allowed value of 5. Hence, we fixed it to a solar abundance value for obs3 to obs8, and it did not affect the fit statistics (Δχ² = −3).

The inclination angles (in degrees) found for XMM-Newton observations are consistent with a best-fit value of $45{^\circ }_{-7}^{+5}$. The inclination angle is not supposed to change during the human timescale, and we froze this parameter to 45° for all the other observations.

The ionization of the accretion disk is characterized by the parameter $\mathrm{log}\xi$, which varied significantly between observations. Due to better data quality, we were able to constrain $\mathrm{log}\xi$ to ${0.52}_{-0.30}^{+0.77}\,\mathrm{erg}\,\mathrm{cm}\,{{\rm{s}}}^{-1}$ for obs2 and got an upper limit of 2.01 for obs1. For the rest of the observations we got a highly ionized accretion disk with the best-fit parameter value ranging between ${2.72}_{-0.68}^{+0.18}\,\mathrm{erg}\,\mathrm{cm}\,{{\rm{s}}}^{-1}$ in obs4 and ${3.30}_{-1.45}^{+0.91}\,\mathrm{erg}\,\mathrm{cm}\,{{\rm{s}}}^{-1}$ in obs8 but remained consistent within errors.

The reflection fraction (R) determines the ratio of photons emitted toward the disk compared to escaping to infinity. We obtained relatively higher values (∼0.46) for obs1 and obs2. However, we note that the best-fit values of the reflection fraction are poorly constrained and are within the 3σ limit.

For the model MyTorus we fixed the inclination angle to 45° and tied the column density of both MyTorusL and MyTorusS in all our observations. Due to the lack of data beyond 10 keV, in XMM-Newton observations, MyTorus normalization was tied between MyTorusL and MyTorusS. We were unable to constrain the column density and found it to be pegged at 10 × 10²⁴ cm⁻² for all observations. Hence, we fixed this value to 10 × 10²⁴ cm⁻² for all our observations and made flux normalization the only free parameter during our spectral fitting.

The optical/UV data from simultaneous XMM-Newton OM and Swift UVOT instruments were modeled with diskbb. To avoid the effects of host galaxy and starburst contribution, we did not consider the V and B bands in our spectral fits, which are more likely to be affected by these phenomenon. Further, the optical/UV flux may be contaminated owing to emission from the broad-line region (BLR)/narrow-line region. Unfortunately, the contamination cannot be quantified accurately from the XMM-Newton OM data. Using HST data, Mathur et al. (2018) measured the BLR continuum flux to be ∼7%−10% of the UV continuum flux. Following this measurement, we corrected the count rates in the UV and derived the intrinsic count rates of the source. We wrote these count rates in an OGIP compliant spectral file generated using om2pha and uvot2pha tasks for XMM-Newton and Swift, respectively. We also introduced a typical 5% systematic uncertainty (Laha et al. 2013; Ghosh et al. 2018) to the optical–UV data sets for each epoch to account for the intrinsic galactic extinction and the host galaxy contribution. Here, we froze the relxill and MyTorus model parameters to their best-fit values obtained from the 0.3–50 keV X-ray spectral fitting. We included a REDDEN component to account for the interstellar extinction. The diskbb model describes the optical/UV band well and provides a satisfactory fit for all the observations (see Appendix B for details).

3.2.2. Warm Comptonization

In our work, we used optxagnf as the warm Comptonization model to describe the soft excess emission. Optxagnf (Done et al. 2012) is an intrinsic thermal Comptonization model that describes (a) the optical/UV spectra of AGNs as multicolor blackbody emission from a color-temperature-corrected disk; (b) the soft excess emission as thermal Comptonization of disk seed photons from an optically thick, low-temperature plasma; and (c) the power-law continuum as thermal Comptonization of disk photons from an optically thin, hot (fixed at 100 keV) plasma. All three components are powered by the gravitational energy released owing to accretion. Parameter r_corona determines the inner radius below which the gravitational energy cannot completely thermalize and is distributed among the soft excess and the power-law components. This ratio is determined by the f_pl parameter. The electron temperature (kT_e ) and optical depth (τ) represent the warm corona responsible for the soft excess emission. The model flux is determined by four parameters, the black hole mass (M_BH), the Eddington ratio (λ_Edd), the comoving distance (D in Mpc), and the dimensionless black hole spin (a). Hence, the model normalization is fixed at unity during spectral fitting. Similar to relxill, we included the two MyTorus model components, MyTorusL and MyTorusS, in our set of models to account for the neutral reflection of hard X-ray continuum from the outer part of the disk.

We fixed the black hole mass of Mrk 590 to 4.75 × 10⁷ M_⊙, determined using the reverberation mapping (Peterson et al. 2004), and the cosmological distance to 112.88 Mpc (Mould et al. 2000). The Optxagnf model needs optical–UV data to constrain the multicolor blackbody emission from the disk. Hence, the model resulted in a poor constraint in parameters when we fitted only the X-ray band with this set of models. To get a better constraint, we included the simultaneous optical–UV data from XMM-Newton and Swift telescopes for all observations except for obs3 and obs4, for which we did not have simultaneous data.

Similar to previous sets of physical models, the fit statistic was insensitive to the black hole spin for all observations. Hence, we fixed the spin parameter to zero and allowed the Eddington ratio (L/L_Edd), the optical depth (τ), the electron temperature (kT_e ), the photon index (Γ), and the f_pl parameter to vary freely. For the MyTorus model components we fixed the inclination angle to 45° and allowed the MyTorusL and MyTorusS normalization parameters to vary freely except for obs1 and obs2, where we tied them together. Similar to the case of the relxill model, we found the MyTorus column density to be pegged at 10 × 10²⁴ cm⁻² for all observations and hence fixed this value to 10 × 10²⁴ cm⁻² for all the spectral analyses. The best-fit parameters are quoted in Table 4.

Table 4. Best-fit Parameters for Observations of Mrk 590 with the Second Set of Physical Models

Component	Parameter	obs1	obs2	obs3	obs4	obs5	obs6	obs7	obs8
Gal. abs.	NH (1020 cm−2)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)	2.77 (f)
diskbb	Tin (eV)	${1.22}_{-0.20}^{+0.12}$	${0.78}_{-0.05}^{+0.03}$	⋯	⋯	${0.25}_{-0.02}^{+0.05}$	⋯	⋯	⋯
	norm (×1012)	${0.25}_{-0.08}^{+0.22}$	${2.92}_{-0.56}^{+1.25}$	⋯	⋯	${880}_{-595}^{+112}$	⋯	⋯	⋯
	A Δχ2/dof	91/2	311/2	⋯	⋯	140/2	⋯	⋯	⋯
optxagnf	${M}_{\mathrm{BH}}^{a}$	4.75(f)	4.75(f)	4.75(f)	4.75(f)	4.75(f)	4.75(f)	4.75(f)	4.75(f)
	d (Mpc)	113(f)	113(f)	113(f)	113(f)	113(f)	113(f)	113(f)	113(f)
	$(\tfrac{L}{{L}_{E}})$	${0.006}_{-0.001}^{+0.003}$	${0.009}_{-0.002}^{+0.001}$	${0.22}_{-0.01}^{+0.01}$	${1.20}_{-1.05}^{+0.38}$	${0.005}_{-0.001}^{+0.003}$	${0.030}_{-0.002}^{+0.002}$	${0.020}_{-0.001}^{+0.001}$	${0.012}_{-0.001}^{+0.001}$
	kTe (keV)	${0.18}_{-0.08}^{+0.14}$	${0.17}_{-0.04}^{+0.05}$	${0.05}_{-0.01}^{+0.27}$	0.05(t)	>0.48	${0.03}_{-0.01}^{+0.01}$	>0.57	>0.55
	τ	>16	>23	>9	99(t)	>6	>71	>40	>31
	rcor (rg )	>63	>93	${9.8}_{-0.1}^{+0.1}$	>7.2	${55}_{-6}^{+10}$	>89	${78.8}_{-0.1}^{+0.1}$	${64.7}_{-1.3}^{+0.9}$
	a	0(f)	0(f)	0(f)	0(f)	0(f)	0(f)	0(f)	0(f)
	fpl	${0.97}_{-0.11}^{+0.02}$	${0.97}_{-0.03}^{+0.01}$	${0.62}_{-0.01}^{+0.37}$	>0.52	${0.98}_{-0.03}^{+0.01}$	${0.50}_{-0.04}^{+0.06}$	${0.99}_{-0.01}^{+0.01}$	${0.98}_{-0.01}^{+0.01}$
	Γ	${1.78}_{-0.16}^{+0.11}$	${1.74}_{-0.05}^{+0.05}$	${1.71}_{-0.01}^{+0.01}$	${1.68}_{-0.01}^{+0.01}$	${1.62}_{-0.08}^{+0.06}$	${1.66}_{-0.02}^{+0.03}$	${1.65}_{-0.02}^{+0.02}$	${1.61}_{-0.01}^{+0.01}$
MYTorusL	i (deg)	45 (f)	45(f)	45(f)	45(f)	45(f)	45(f)	45(f)	45(f)
	norm (10−3)	${5.62}_{-2.29}^{+2.81}$	${4.59}_{-1.16}^{+1.21}$	${4.38}_{-0.72}^{+0.73}$	${3.31}_{-0.76}^{+0.79}$	${2.09}_{-0.99}^{+1.05}$	${5.56}_{-2.05}^{+2.10}$	${7.67}_{-1.24}^{+1.26}$	${4.34}_{-0.04}^{+0.88}$
MYTorusS	NH (1024 cm−2)	10.0(f)	10.0(f)	10.0(t)	10.0(t)	10.0(t)	10.0(t)	10.0(t)	10.0(t)
	norm (10−3)	5.14 (f)	3.00(f)	<0.48	<1.09	<0.76	<0.52	<0.05	<0.06
	χ2/dof	160/124	218/163	898/831	546/550	189/209	340/316	972/840	632/615

Note. In XSPEC, the models read as (constant × tbabs × (diskbb + optxagnf+ MYTorus)). The quoted best-fit values for Suzaku observations (obs3 and obs4) are from the spectral fitting of X-ray band only. (f) indicates a frozen parameter. (t) indicates a tied parameter between different observation. (*) indicates that parameters are not constrained. (a): in units of 10⁷ M_⊙.

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The optxagnf model produced comparable fit statistics of the broadband source spectra for all the observations except for obs7 and obs8. For obs7 and obs8, we find that this set of models provided a poor description of the observed high-energy cutoff above 20 keV. Figure 2 shows the residuals and the theoretical model for these two observations. The optxagnf model can describe the UV bump, and hence a separate diskbb is not required in principle. However, given that we get a poor fit using the optxagnf model alone for obs1, obs2, and obs5, we added a separate diskbb component, which improved the fit statistics by Δχ² ∼120–130. Clearly the optical–UV data require this additional component. We notice that the optical–UV flux measured during obs1, obs2, and obs5 is relatively lower compared to obs6, obs7, and obs8. However, we were unable to constrain the optical depth (τ) and the coronal radius (r_corona) for most of the observations. We found a sub-Eddington accretion rate (1%–3%) in all observations. The variability in power-law photon index (Γ) between observations was not statistically significant. We found very high values of f_pl in all observations except for obs6. Interestingly, for obs6 we found the electron temperature (kT_e ) to be very low compared to other observations. The MyTorusL flux normalization values were consistent between observations. We could not constrain the MyTorusS flux normalization and only got an upper limit on the flux.

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3.2.3. Summary

Our spectral analysis revealed a significant variability in soft excess flux between observations of Mrk 590. The soft excess was present (${3.7}_{-0.5}^{+1.0}\times {10}^{-13}\,\mathrm{erg}\,{\mathrm{cm}}^{-2}\,{{\rm{s}}}^{-1}$) in 2004 (obs2) but vanished/was undetected ( <1.6 × 10⁻¹³ erg cm⁻² s⁻¹) in 2011 (obs3 and obs4), within a period of 7 yr. This excess emission never reappeared in any of the later observations until 2021. We calculated the 0.3 − 2.0 keV soft excess flux from the phenomenological best fit and quoted these values in Table 5. The soft excess fluxes (F_SE) in obs1 and obs2 are ${F}_{\mathrm{SE}}={4.3}_{-0.6}^{+0.6}\,\times {10}^{-13}\,\mathrm{erg}\,{\mathrm{cm}}^{-2}\,{{\rm{s}}}^{-1}$ and ${3.7}_{-0.5}^{+1.0}\times {10}^{-13}\,\mathrm{erg}\,{\mathrm{cm}}^{-2}\,{{\rm{s}}}^{-1}$, respectively. We do not detect the soft excess in obs3 and obs4, and the corresponding upper limits are F_SE <1.6 × 10⁻¹³ erg cm⁻² s⁻¹ and <1.4 × 10⁻¹³ erg cm⁻² s⁻¹, respectively. For obs5, obs6, obs7, and obs8 (NuSTAR plus Swift observations) the upper limits on the soft excess flux are F_SE <0.7 × 10⁻¹³ erg cm⁻² s⁻¹, <0.6 × 10⁻¹³ erg cm⁻² s⁻¹, <1.9 × 10⁻¹³ erg cm⁻² s⁻¹, and <0.8 × 10⁻¹³ erg cm⁻² s⁻¹, respectively. From Figure 3 (top panel) we see that the soft excess flux drops by a factor of four within 9 yr, from 2002 to 2011.

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Table 5. Fluxes of the Different Spectral Components of Mrk 590 Obtained from the Observations Used in Our Work

Spectral	Flux	Flux	Flux	Flux	Flux	Flux	Flux	Flux
Component	obs1	obs2	obs3	obs4	obs5	obs6	obs7	obs8
	(Jan. 2002)	(July 2004)	(Jan. 2011)	(Jan. 2011)	(Feb. 2016)	(Oct. 2018)	(Jan. 2020)	(Jan. 2021)
Soft excess (×10−13)	${4.27}_{-0.55}^{+0.63}$	${3.72}_{-0.48}^{+0.96}$	<1.60	<1.40	<0.65	<0.55	<1.97	<0.79
Power law a (×10−12)	${4.17}_{-0.19}^{+0.51}$	${5.89}_{-0.27}^{+0.28}$	${7.24}_{-0.16}^{+0.17}$	${6.76}_{-0.15}^{+0.16}$	${2.75}_{-0.18}^{+0.20}$	${9.12}_{-0.41}^{+0.43}$	${12.10}_{-0.21}^{+0.20}$	${6.13}_{-0.13}^{+0.12}$
Fe Kα emission line (×10−13)	${1.02}_{-0.44}^{+0.80}$	${0.79}_{-0.27}^{+0.21}$	${1.15}_{-0.24}^{+0.26}$	${1.00}_{-0.28}^{+0.32}$	${0.71}_{-0.39}^{+0.46}$	${2.24}_{-1.09}^{+1.07}$	${2.78}_{-0.98}^{+1.04}$	${1.61}_{-0.41}^{+0.47}$
Neutral reflection b (×10−13)	⋯	⋯	<4.47	<8.32	${3.31}_{-3.30}^{+3.30}$	<1.23	<0.04	<0.03
UV monochromatic c
UVW2 (×10−15)	2.80 ± 0.10	2.03 ± 0.08	⋯	⋯	1.95 ± 0.08	6.51 ± 0.14	5.08 ± 0.14	3.19 ± 0.11
UVM2 (×10−15)	2.60 ± 0.08	3.19 ± 0.05	⋯	⋯	1.68 ± 0.09	5.73 ± 0.20	4.29 ± 0.18	⋯
UVW1 (×10−15)	3.50 ± 0.05	3.19 ± 0.04	⋯	⋯	2.17 ± 0.06	6.07 ± 0.13	4.83 ± 0.16	3.17 ± 0.13
U (×10−15)	2.67 ± 0.03	3.13 ± 0.02	⋯	⋯	2.39 ± 0.06	4.88 ± 0.12	3.67 ± 0.13	2.81 ± 0.10
F2 keV (×10−12)	${1.12}_{-0.05}^{+0.05}$	${1.51}_{-0.03}^{+0.04}$	${1.78}_{-0.04}^{+0.04}$	${1.59}_{-0.04}^{+0.04}$	${0.63}_{-0.04}^{+0.03}$	${1.91}_{-0.09}^{+0.09}$	${2.49}_{-0.08}^{+0.09}$	${1.29}_{-0.03}^{+0.04}$
αOX	1.228	1.163	⋯	⋯	1.244	1.231	1.148	1.188
logL2−10 keV	${42.82}_{-0.02}^{+0.05}$	${42.96}_{-0.02}^{+0.02}$	${43.06}_{-0.01}^{+0.01}$	${43.02}_{-0.01}^{+0.01}$	${42.64}_{-0.03}^{+0.03}$	${43.15}_{-0.01}^{+0.01}$	${43.27}_{-0.03}^{+0.04}$	${42.98}_{-0.01}^{+0.01}$
Lbol(0.001–100 keV)	${43.89}_{-0.02}^{+0.02}$	${43.99}_{-0.01}^{+0.01}$	⋯	⋯	${43.65}_{-0.02}^{+0.02}$	${44.05}_{-0.01}^{+0.01}$	${44.04}_{-0.01}^{+0.01}$	${43.82}_{-0.01}^{+0.01}$
λEdd	0.0130	0.0163	⋯	⋯	0.0075	0.0187	0.0183	0.0110

Notes.

^a Unabsorbed power-law flux estimated in the energy range 2–10 keV. ^b The reflected emission due to Compton down-scattering of the hard X-ray photons by a neutral medium, as estimated using the model pexrav. We did not quote this for obs1 and obs2, which are XMM-Newton observations and do not cover the range above 10 keV. ^c The UV monochromatic fluxes are measured from XMM-Newton OM and Swift UVOT. See Table 2 for the model fit. The fluxes are in units of erg cm⁻² s⁻¹ Å⁻¹.

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We detected the presence of a weak Fe Kα emission line in the source spectra for all observations with a flux of ${1.0}_{-0.4}^{+0.8}\times {10}^{-13}\,\mathrm{erg}\,{\mathrm{cm}}^{-2}\,{{\rm{s}}}^{-1}$ for obs1, which remains consistent within 3σ uncertainties for all the observations (see Table 5). The iron line is narrow in nature (σ <0.1 keV) for obs1, obs2, obs3, and obs4. We found an upper limit on the Fe line width of σ <0.45 and <0.44 keV for obs5 and obs6, respectively. For obs7 and obs8, due to longer exposure of NuSTAR, we were able to constrain the Fe line width and found $\sigma ={0.24}_{-0.16}^{+0.33}\,\mathrm{keV}$ and ${0.19}_{-0.10}^{+0.14}\,\mathrm{keV}$ for obs7 and obs8, respectively. In all observations, the addition of pexrav to the phenomenological set of models did not improve the fit statistic, and we only got an upper limit on the reflection fraction. This result shows no Compton hump present above 10 keV in the source spectra.

We found the power-law photon index Γ for obs1 and obs2 to be slightly steeper than in the rest of the observations. Although the best-fit value of Γ remains within errors for the optxagnf model, it showed a significant variation when we used the model relxill (e.g., ${\rm{\Gamma }}={1.88}_{-0.08}^{+0.02}$ in 2002 and ${1.58}_{-0.03}^{+0.02}$ in 2021). These results are in agreement with previous studies (Laha et al. 2018a; Ezhikode et al. 2020). We have calculated the 2–10 keV unabsorbed power-law flux and quoted these values in Table 5. We also estimated the UV monochromatic flux at 2500 Å using the UVW1 band from the OM and UVOT filter of XMM-Newton and Swift, respectively (see Figure 3). We corrected the source count rates for the Galactic extinction using the CCM extinction law (Cardelli et al. 1989) with a color excess of E(B − V) = 0.0134 and a ratio of total to selective extinction of R_V = A_V /E(B − V) = 3.1, where A_V is the extinction in the V band. From Table 5, it is evident that both 2–10 keV power-law flux and UV monochromatic flux varied significantly between observations. The power-law flux rises from ${4.2}_{-0.2}^{+0.5}\,\times {10}^{-12}\,\mathrm{erg}\,{\mathrm{cm}}^{-2}\,{{\rm{s}}}^{-1}$ in 2002 to ${12.1}_{-0.2}^{+0.2}\,\times {10}^{-12}\,\mathrm{erg}\,{\mathrm{cm}}^{-2}\,{{\rm{s}}}^{-1}$ in 2020 and then declines to ${6.1}_{-0.1}^{+0.1}\times {10}^{-12}\,\mathrm{erg}\,{\mathrm{cm}}^{-2}\,{{\rm{s}}}^{-1}$ in 2021. The UV monochromatic flux (UVW1) also rises from (3.50 ± 0.05)×10⁻¹⁵ erg cm⁻² s⁻¹Å⁻¹ in 2002 to (6.07 ±0.13)×10⁻¹⁵ erg cm⁻² s⁻¹ Å⁻¹ in 2018 and then declines to (3.17 ± 0.13)×10⁻¹⁵ erg cm⁻² s⁻¹ Å⁻¹ in 2021. Figure 3 shows the soft excess flux, the power-law flux, the UV monochromatic flux, and the iron emission-line flux variation for the past two decades. We notice that the power-law flux and the UV flux follow the same temporal pattern. However, the soft excess flux does not follow the trend and shows a unique spectral or flux evolution.

We summarize the spectral evolution of Mrk 590 in Figure 4 using the relxill model as described in Table 3. The figure shows the best-fitting UV to X-ray broadband models derived from all the observations used in our work, obs1 in black, obs2 in red, obs5 in green, obs6 in blue, obs7 in cyan, and obs8 in magenta. The ionized reflection model describing the soft excess and power-law emission and the optical–UV emission described by the diskbb are shown for each epoch with different markers and colors. The MyTorus model components, MyTorusL and MyTorusS, are shown with dotted and dashed lines, respectively. Clearly Mrk 590 has shown some unique disk and corona properties over the past few decades.

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We have calculated the bolometric luminosities (L_bol) of the source at each epoch from our physical broadband spectral modeling in the energy range 0.001–100 keV. We preferred the ionized reflection model for this purpose, which provided a relatively better description of the source spectra at all epochs. Next, we estimated the λ_Edd = L_bol/L_Edd during each epoch assuming a black hole mass of 4.5×10⁷ M_BH. We find a sub-Eddington accretion rate for all the observations and the values are listed in Table 5. However, the values are not consistent between observations, and we plotted them (λ_Edd) in Figure 3. As expected, the value of accretion rate follows a similar trend of variation to that of power-law and UV monochromatic flux. We also plotted the logarithm of the Eddington ratio versus the power-law slope Γ at each epoch (Figure 5).

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We used absorption-corrected UV monochromatic and 2 keV fluxes to calculate ${\alpha }_{\mathrm{OX}}=-0.384\mathrm{log}[{L}_{2500{\rm{\mathring{\rm A} }}}/{L}_{2\,\mathrm{keV}}]$ (Tananbaum et al. 1979), defined as the power-law slope joining the 2 keV and the 2500 Å flux for a given source. The α_OX values show variation between observations and are quoted in Table 5. Figure 5 shows the correlation between the α_OX versus L_2500Å and the L_{2 keV} versus L_2500Å. To compare our results, we overplot the best-fit α_OX versus L_2500Å relation found by Lusso & Risaliti (2016), and we see that they do not follow this relation.

The soft excess emission below 2 keV is very common in type 1 AGNs, and the origin is still under debate. Our phenomenological set of models revealed the presence of soft excess emission only in obs1 and obs2. Our result is consistent with Rivers et al. (2012).

We first investigate whether variable obscuration is responsible for this flux variation. An obscuring cloud crossing the line of sight and causing changes in observed light curves (Goodrich 1989; Guo et al. 2016) may cause the soft excess to vanish. However, we did not find the presence of intrinsic absorption in any of the source spectra. This clearly rules out variable obscuration to be the reason behind this vanishing soft excess. Next, we discuss the results of two physical models used in our work to find the origin of the soft excess emission.

In the thermal Comptonization model optxagnf, soft excess emission is produced owing to thermal Comptonization of disk optical–UV photons by a warm (kT ∼ 0.1–0.2 keV), optically thick (τ ∼ 10−20) corona surrounding the inner regions of the disk. Hence, the vanishing soft excess requires a vanishing warm corona or simply a significant change in the size of the warm Comptonizing medium (r_cor). Here, the disk makes transitions between a cold+warm disk and a cold disk. This excess energy generation in the innermost disk must be provided by the increase in accretion rate, which could not all be released in the form of radiation. This excess energy raises the temperature and pressure and transforms the innermost accretion disk to the warm Comptonization medium. Hence, the absence of soft excess should follow with a decrease in the accretion rate (Done et al. 2012; Tripathi & Dewangan 2022). If this scenario is true, there must be a correlation between the soft excess flux, optical–UV flux, and mass accretion rate. In addition, the spectral state transition due to disk evaporation and change in the accretion rate leads to spectral hardening of the photons arising from the hot corona.

However, in Mrk 590, optxagnf needed an additional diskbb component to model the UV bump in obs1, obs2, and obs5, as the model overestimated the optical–UV flux when extrapolated from the spectral fitting of the X-ray energy band. Now, if the soft excess is very weak or absent, the warm corona responsible for the soft excess may instead contribute to the optical–UV band. However, we found that we did not require this diskbb component for obs6, obs7, and obs8, where soft excess is absent. The improvement in fit statistics after the addition of the diskbb component was significant. Table 5 shows that UV monochromatic flux measured for obs1, obs2, and obs5 is significantly lower than UV flux measured during obs6, obs7, and obs8. In Mrk 590, we see the soft excess flux drop four times within 7 yr. We expect an increase in the value of r_cor and a decrease in the accretion rate as found in other CLAGNs, e.g., Mrk 1018 (Noda & Done 2018) and NGC 1566 (Tripathi & Dewangan 2022). In NGC 1566, the soft excess flux component decreases by a factor >45, and a significant change in the size of the warm Comptonizing medium (r_cor) is found, where r_cor increased from ∼26r _g during the high flux state in 2015 to 50r _g during the low flux state in 2018. However, for Mrk 590, we could not determine the exact size (r_corona), electron temperature, or optical depth (kT_e and τ) of the Comptonizing corona even when the soft excess is present. We did not find any significant decrease in luminosity and accretion rate. In addition, the optxagnf best fit (Table 4) shows no significant change in power-law Γ or f_pl, indicating no spectral hardening between observations. The results indicate no correlation between the optical–UV and soft excess flux. Figures 6 and 7 show this lack of correlation, where we plot the soft excess versus the UV flux and the accretion rate, respectively. We also found that optxagnf could not model the observed exponential cutoff in power-law continuum in obs7 and obs8. The source spectra did not require black hole spin to model the soft excess emission. This result is consistent with Mrk 1018 and NGC 1566 but contradicts the recent findings in other type 1 AGNs where soft excess is present (García et al. 2019; Ghosh & Laha 2020, 2021; Xu et al. 2021). These results show that the thermal Comptonization of disk photons, which successfully explained the soft excess flux variation in other CLAGNs such as NGC 1566 and Mrk 1018, is unable to explain the vanishing soft excess and high-energy cutoff observed in Mrk 590.

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The ionized reflection model relxill provided a good fit for all the observations. In this model, the untruncated accretion disk approaches the innermost stable circular orbit owing to high black hole spin. The hard X-ray photons from the corona illuminate the disk, ionize it, and emit fluorescent emission lines. These lines are blurred owing to extreme gravity near SMBH. In this scenario, the power-law flux and the soft excess flux should have a strong correlation between them. Hence, the decrease in soft excess flux may occur owing to changes in disk and corona properties. Some possibilities are (a) the disk becoming a truncated one (r_in >50r_g ) owing to disk evaporation or (b) the disk becoming highly ionized (ξ ∼ 10⁴ erg cm s⁻¹; Ross & Fabian 2005). The disk may become highly ionized owing to spectral hardening, as harder illuminating spectra have greater ionizing power. But spectral hardening will also give rise to a strong and broad Fe emission line and a Compton hump. The change in soft excess strength should also affect the reflected flux or the reflection fraction that determines the ratio of intensity emitted toward the disk compared to escaping to infinity.

We note that Mrk 590 does not fit this description. The soft excesses in obs1 and obs2 were described by a nonrotating black hole (r_in fixed at 6r_g ) and a slightly ionized accretion disk. This result contradicts recent studies of other type 1 AGNs, where the reflection model favors a rotating black hole and the inner part of the disk approaches the innermost stable circular orbit (1.25r_g ; García et al. 2019; Ghosh & Laha2020, 2021). We note a significant increase in disk ionization between observations (from $\mathrm{log}\xi ={0.52}_{-0.30}^{+0.77}$ in obs2 to $\mathrm{log}\xi \sim 3$ in obs3). We also found a significant decrease in Γ value from obs1 to obs8, indicating a spectral hardening. For obs7 and obs8, we found a well-constrained, high-energy cutoff of ${92}_{-25}^{+55}\,\mathrm{keV}$ and ${60}_{-8}^{+10}\,\mathrm{keV}$, respectively, better described by the relxill model. Although these values are relatively lower compared to other Seyfert 1 galaxies (200–300 keV; Ghosh et al. 2016; Ricci et al. 2017; Fabian et al. 2017; Akylas & Georgantopoulos 2021), similar low values of E_cut have been found in recent sample studies of Swift/BAT-selected AGNs (Kamraj et al. 2022). This low-energy cutoff may indicate a decrease in the plasma temperature of the corona only if it was higher during obs1 and obs2. But we cannot test this scenario owing to the nonavailability of data beyond 10 keV. Now hard X-ray photons illuminate the disk, so less energetic photons mean low illumination. However, we note that the disk is still moderately ionized and is capable of Ross & Fabian (2005) producing fluorescent emission lines. More importantly, we did not find any broad Fe emission line or Compton hump in any of the source spectra. The reflected flux and reflection fraction do not show statistically significant variation and remain within the 3σ value. When we plotted the soft excess flux versus the power-law flux (2–10 keV), we did not find any significant correlation (see Figure 8). This result is consistent with Boissay et al. (2016), where the shape of reflection at hard X-rays stays constant when the soft excess varies, showing an absence of a link between reflection and soft excess. In Mrk 590, the power law, the Fe line emission, and UV monochromatic flux follow the same temporal pattern (Figure 3). This result suggests that the disk and corona are most likely evolving together. However, the soft excess is not responding to this change in disk–corona properties. Hence, the soft excess emission observed in Mrk 590 is not due to ionized reflection from the disk.

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Hence, we have two possibilities. Either we cannot distinguish the differences in the ionized reflection and warm Comptonization models owing to low-quality data, or these models cannot describe the vanishing soft excess feature observed in Mrk 590. Next, we discuss in detail the possibility of change in the accretion disk profile behind the spectral variability in Mrk 590 as found in other CLAGNs.

Previously, state change due to disk evaporation or condensation associated with a factor 2−4 decrease or increase in luminosity, significant mass accretion rate change, or the combination of both has been suggested as the reason behind the changing-look nature of Mrk 590 (Mathur et al. 2018; Noda & Done 2018; Yang et al. 2021). The soft excess emission in Mrk 590 vanished within 7 yr. This timescale puts an upper limit on the position of the reprocessing material within 7 lt-yr, or roughly 2 pc. This distance is significantly large compared to the distance (10–100 lt-days) of the BLR from SMBH in a typical AGN but comparable to the distance of the torus (∼few parsecs). The BLR in Mrk 590 also has gone through some dramatic changes, as after ∼10 yr of absence, the optical broad emission lines of Mrk 590 have reappeared (Raimundo et al. 2019). The absence of soft excess even when the source changed its type suggests a lack of correlation between the two phenomena. For further investigation, we study the disk instability in Mrk 590 that may cause the observed flux variation.

In Mrk 590, we see a drop in total accretion luminosity from 9.8×10⁴³ erg s⁻¹ (in obs2) to 4.4×10⁴³ erg s⁻¹ (in obs5), rising again to 1.2×10⁴⁴ erg s⁻¹ (in obs6) over a timescale of ∼14 yr. The amplitude of bolometric drop requires a change in either mass accretion rate or efficiency (or both) of the accretion flow. But the observed timescale (∼14 yr) is too fast for the mass accretion rate to change through a standard disk (viscous timescale) by many orders of magnitude (Noda & Done 2018; Wang et al. 2019; Ricci et al. 2020) and poses a problem for any standard disk model. If we compare the three timescales for a standard thin disk, the dynamical timescale is the fastest, then the thermal timescale, and then the viscous timescale. To compare the disk variability timescale in Mrk 590, we calculate the accretion disk timescales at the inner radius of the disk obtained from the best-fit broadband spectral model. The dynamical (t_dyn), thermal (t_th), and viscous (t_vis) timescales of the accretion disk are given by (Czerny 2006)

$$\begin{eqnarray}&&{t}_{\mathrm{dyn}}={\left(\displaystyle \frac{{R}^{3}}{{{GM}}_{\mathrm{BH}}}\right)}^{1/2}\end{eqnarray} \tag{ 1 }$$

$$\begin{eqnarray}&&{t}_{\mathrm{th}}=\displaystyle \frac{1}{\alpha }{t}_{\mathrm{dyn}},\end{eqnarray} \tag{ 2 }$$

and

$$\begin{eqnarray}&&{t}_{\mathrm{vis}}\sim \displaystyle \frac{1}{\alpha }{\left(\displaystyle \frac{R}{H}\right)}^{2}{t}_{\mathrm{dyn}},\end{eqnarray} \tag{ 3 }$$

where R is the radial distance in the disk, α is the viscosity parameter, and H is the height of the disk. To estimate these values, we first calculate the accretion disk temperature of ∼1.23 eV for an inner disk radius of ∼100R_g , accretion rate of $\dot{M}\sim 0.018$, and a black hole mass of 4.75×10⁷ M_⊙. This implies H/R = c_s /V_ϕ ∼ 3.6×10⁻⁴ (where ${c}_{s}=\sqrt{{kT}/{m}_{p}}$ is the sound speed and ${v}_{\phi }=\sqrt{{{GM}}_{\mathrm{BH}}/R}$ is the Keplerian orbital velocity). Assuming α = 0.1, we finally estimated the dynamical, thermal, and viscous timescales to be t_dyn = 3.7 days, t_th = 37 days, and t_vis = 7×10⁵ yr, respectively. So the timescale (7 yr) of flux variability in Mrk 590 is much smaller compared to the viscous timescale but longer than the dynamical and thermal timescales at an inner radius of 100R_g . If we consider the sound crossing time of t_s ∼ 100R_g /c_s ∼ 20 yr, the value is still an order of magnitude higher than the changing-look time of Mrk 590. So the flux variability in Mrk 590 is not likely due to pressure instabilities in the disk. However, only if we consider an untruncated thin accretion disk up to 10r_g , then, does the variability timescale become much shorter and comparable to the observed timescale in Mrk 590. Similar procedures mentioned above estimate the timescales to be t_dyn = 2.84 hr, t_th = 28 hr, and t_vis = 4×10³ yr, respectively. The sound crossing time becomes t_s ∼ 10R_g /c_s ∼ 1 yr. Following this assumption, we have an inner disk temperature of ∼ 6 eV, which is inconsistent with the disk temperature we got from our spectral best fit using the ionized reflection model (see Table 3). In addition, the disk this close to SMBH will affect the Keplerian frequency (Kato 2001) and should give rise to stronger reflection features in the source spectra, common in other type 1 AGNs but absent in our spectral analysis. This points toward a possible disk truncation above 10r_g . In addition, the change in accretion profile should affect the accretion rate in Mrk 590. Raimundo et al. (2019) found that the broad Balmer broad emission lines in Mrk 590 reappeared in 2017 October, within a timescale of decades. A similar behavior has also been observed in Mrk 1018 (McElroy et al. 2016; Husemann et al. 2016). We note that this reappearance of Balmer lines coincides with the increase in the mass accretion rate (see Table 5 and Figure 3). However, when we plot the soft excess flux and λ_Edd (Figure 7), we do not find any correlation between these two parameters as previously suggested by Noda & Done (2018), Mathur et al. (2018), and Yang et al. (2021). Instead, we found a relatively higher Eddington ratio even when the soft excess is not present. We also note that our best-fit accretion rate (∼0.02 ± 0.01) is consistent with previous studies (∼0.03 ± 0.01; Laha et al. 2018a). So the origin of soft excess emission in Mrk 590 is likely to be not related to the change in the accretion rate.

In Mrk 590, the observed UV and power-law flux variability follows the same temporal pattern. The corona cools down, and the disk becomes more ionized between observations. Hence, a change in the nature of the accretion disk and corona is evident, but not the fundamental process through which the disk–corona evolves. To investigate further, we studied the relation between the 2–10 keV power-law slope, Γ, and the Eddington ratio (λ_Edd) at each epoch. This exercise helps us check how efficiently the disk photons are coupled with the hot corona and how efficient the central engines are. A strong coupling between Γ and λ_Edd implies that a higher accretion rate cools off the corona faster, leading to steeper power-law slopes (Pounds et al. 1995). Previously Baumgartner et al. (2013) and Trakhtenbrot et al. (2017) studied a sample of radio-quiet AGNs and a BAT-selected AGN sample, respectively, and found a strong correlation between Γ and λ_Edd. Gu & Cao (2009) investigated the relation for a sample of 57 low-luminosity AGNs in the local universe and found that they follow an anticorrelation. This contradiction suggests the possibility of two modes of accretion above/below some critical transition value of λ_Edd. From Figure 5, we find that Mrk 590 does not show any such strong correlation or anticorrelation between the spectral slope and the Eddington rate. This is consistent with Laha et al. (2018b), who did not find any such strong correlation in a sample of low-luminosity QSOs. These results clearly show that the disk/corona interaction in Mrk 590 does not follow the typical disk–corona properties of Seyfert 1 AGNs and has unique characteristics. In this context we note that there has been a recent discovery of a changing-look phenomenon in AGN 1ES 1927+654 (Trakhtenbrot et al. 2019; Ricci et al. 2020; Laha et al. 2022), the origin of which is still debated. However, the radio, optical, UV, and X-ray observations point toward an increase of accretion, probably due to magnetic flux inversion, as the primary cause of this event (Scepi et al. 2021; Laha et al. 2022).

The relativistic reflection from the ionized accretion disk cannot explain the spectral variability in Mrk 590. The soft excess flux variation does not correlate with the power-law continuum flux. The lack of a broad Fe emission line in the spectra indicates that relativistic reflection does not dominate the source spectra. We were unable to constrain the Fe abundance of the disk, which is previously observed in other Seyfert 1 AGNs (Fabian et al. 2009; Dauser et al. 2012; García et al. 2018; Ghosh & Laha 2020; Laha & Ghosh 2021) as well. A narrow Fe emission line in the X-ray spectra suggests a distant neutral reflection of the hard X-ray continuum from the outer part of the disk or torus. In Mrk 590, we found that both Fe line emission and the power-law flux follow the same temporal pattern, supporting the idea that narrow Fe line emission is most likely due to a neutral reflection of hard X-ray photons from the outer part of the disk. However, we do not see any Compton hump that arises owing to Compton down-scattering of high-energy photons by the cold disk or torus. This result is inconsistent with the typical neutral reflection observed in the X-ray spectra of type 1 AGNs. The reflection fraction value was within 3σ significance throughout observations, and we did not find any change in the disk properties except for the ionization. These results indicate a complex reflection scenario that does not follow the typical disk–corona interaction in reflection-dominated type 1 AGNs.