11institutetext: INAF – Astronomical Observatory of Rome, Via Frascati 33,
00040 Monte Porzio Catone, Italy 22institutetext: Instituto de Estudios
Astrofísicos, Facultad de Ingeniería y Ciencias, Universidad Diego Portales,
Av. Ejército Libertador 441, Santiago, Chile. 33institutetext: Kavli Institute
for Astronomy and Astrophysics, Peking University, Beijing 100871, China
44institutetext: George Mason University, Department of Physics & Astronomy,
MS 3F3, 4400 University Drive, Fairfax, VA 22030, USA 55institutetext: Physics
Department, Tor Vergata University of Rome, Via della Ricerca Scientifica 1,
00133 Rome, Italy 66institutetext: INFN - Rome Tor Vergata, Via della Ricerca
Scientifica 1, 00133 Rome, Italy 77institutetext: Department of Astronomy,
University of Maryland, College Park, MD 20742, USA 88institutetext: NASA
Goddard Space Flight Center, Code 662, Greenbelt, MD 20771, USA
99institutetext: RIKEN Cluster for Pioneering Research, 2-1 Hirosawa, Wako,
Saitama 351-0198, Japan 1010institutetext: MIT Kavli Institute for
Astrophysics and Space Research, Massachusetts Institute of Technology,
Cambridge, MA 02139, USA 1111institutetext: INAF - Istituto di Astrofisica e
Planetologia Spaziali, Via del Fosso del Cavaliere 100, 00133, Roma, Italy
1212institutetext: Yale Center for Astronomy & Astrophysics, 52 Hillhouse
Avenue, New Haven, CT 06511, USA 1313institutetext: Joint ALMA Observatory,
Avenida Alonso de Cordova 3107, Vitacura 7630355, Santiago, Chile
1414institutetext: Eureka Scientific, 2452 Delmer Street Suite 100, Oakland,
CA 94602-3017, USA 1515institutetext: Space Science Institute, 4750 Walnut
Street, Suite 205, Boulder, CO 80301, USA 1616institutetext: Department of
Physics and Astronomy, West Virginia University, P.O. Box 6315, Morgantown, WV
26506, USA 1717institutetext: Department of Astronomy, University of Maryland,
College Park, MD 20742, USA 1818institutetext: Department of Astronomy,
University of Geneva, 1205 Versoix, Switzerland 1919institutetext: National
Radio Astronomy Observatory, 520 Edgemont Road, Charlottesville, VA 22903, USA
2020institutetext: Department of Astronomy, University of Florida, P.O. Box
112055, Gainesville, FL 32611, USA 2121institutetext: Department of Astronomy,
University of Virginia, 530 McCormick Road, Charlottesville, VA 22904, USA
2222institutetext: School of Physics and Astronomy, Tel Aviv University, Tel
Aviv 69978, Israel

#  XMM-Newton \- NuSTAR monitoring campaign of the Seyfert 1 galaxy IC 4329A

A. Tortosa<EMAIL_ADDRESS>1122 C. Ricci 223344 E. Shablovinskaia 22
F. Tombesi 5511667788 T. Kawamuro 99 E. Kara 1010 G. Mantovani 1111 M.
Balokovic 1212 C-S. Chang 1313 K. Gendreau 88 M. J. Koss 14141515 T. Liu 1616
M. Loewenstein 881717 S. Paltani 1818 G. C. Privon 191920202121 B.
Trakhtenbrot 2222

We present the results of a joint XMM-Newton and NuSTAR campaign on the active
galactic nucleus (AGN) IC 4329A, consisting of 9 $\times$ 20 ks XMM-Newton
observations, and 5 $\times$ 20 ks NuSTAR observations within nine days,
performed in August 2021. Within each observation, the AGN is not very
variable, with the fractional variability never exceeding 5%. Flux variations
are observed between the different observations, on timescales of days, with a
30% ratio between the minimum and the maximum 2–10 keV flux. These variations
follow the softer-when-brighter behavior typically observed in AGN. In all
observations, a soft excess is clearly present. Consistently with previous
observations, the X-ray spectra of the source exhibit a cut-off energy between
140 and 250 keV, constant within the error in the different observations. We
detected a narrow iron K$\alpha$ line consistent with being constant during
the monitoring, and likely originating in a distant neutral medium. We do not
find evidence of a broad component of the iron line, suggesting that the inner
disk does not produce strong reflection. We find that the reflection component
is weak ($R_{\rm max}=0.009\pm 0.002$). We also found the presence of a warm
absorber component together with an ultra-fast outflow. Looking at their
energetic, these outflows have enough mechanical power to exercise a
significant feedback impact on the AGN surrounding environment.

###### Key Words.:

galaxies:Seyfert – galaxies:active – galaxies:individual:IC4329A – black hole
physics

##  1 Introduction

Active Galactic Nuclei (AGN) are extremely luminous, compact objects located
at the center of massive galaxies. They are powered by accretion of gas onto
the central supermassive black hole (SMBH, $M_{\rm BH}>10^{5}$ $M_{\odot}$;
Salpeter 1964; Ho 2008). The X-ray emission of AGN is due to high-energy
processes occurring in the so-called hot corona, which contains high-energy
electrons and is located near the SMBH (Haardt & Maraschi 1991; Haardt et al.
1994, 1997; Merloni 2003; Fabian et al. 2009; Zoghbi et al. 2012; De Marco et
al. 2013). The hot electrons interact with thermal UV/optical photons emitted
by the accretion disc through inverse-Compton scattering, which boosts the
energy of the photons into the X-ray band. This process produces a broad power
law continuum in the X-ray spectrum with a cutoff at high energy (e.g.,
Sunyaev & Titarchuk 1980; Haardt & Maraschi 1993). Cold circumnuclear
materials, such as the optically thick accretion disc and/or the molecular
torus, absorb and reprocess the X-rays, which results in a feature in the
X-ray spectrum known as Compton reflection, which peaks at around $\sim 30$
keV (Pounds et al. 1990; George & Fabian 1991). Additionally, the reflection
process produces absorption and emission lines, such as the prominent
FeK$\alpha$ emission line at 6.4 keV (Nandra & Pounds 1994), which are the
result of photoelectric absorption and fluorescence (Matt et al. 1997).

IC 4329A is a bright nearby (z=0.01598, Koss et al. 2022) AGN and it is
classified in the optical as a Seyfert 1.5 galaxy (Oh et al. 2022). It is
located at the center of an edge-on host galaxy, with a dust lane passing
through the nucleus. It shows also a dusty lukewarm absorber in the UV band
(i.e. $T\sim 3\times 10^{4}$ K, Crenshaw & Kraemer 2001), characterized by
saturated UV absorption lines (C IV, N V) near the systemic velocity of the
host galaxy, likely responsible for reddening both the continuum and the
emission lines. The black hole mass of IC 4329A is $M_{\rm
BH}=6.8^{+1.2}_{-1.1}\times 10^{7}M_{\odot}$ (Bentz et al. 2023) and its
Eddington ratio is $\lambda_{\rm Edd}=L_{\rm bol}/L_{\rm Edd}=0.17$. This was
computed using the bolometric luminosities calculated from the intrinsic
luminosities in the 14–150 keV range, as shown in Ricci et al. 2017, with a
bolometric correction of 8 (Koss et al. 2022).

IC 4329A has been extensively studied in the X-ray band in the last decades by
all major X-rays satellites including Einstein (Petre et al. 1984; Holt et al.
1989), EXOSAT (Singh et al. 1991), Ginga (Piro et al. 1990), ASCA/RXTE (Done
et al. 2000), BeppoSAX (Perola et al. 2002; Dadina 2007), Chandra (McKernan &
Yaqoob 2004; Shu et al. 2010a), XMM-Newton (Steenbrugge et al. 2005; Nandra et
al. 2007; Mehdipour & Costantini 2018), INTEGRAL (Beckmann et al. 2006),
Swift/BAT (Winter et al. 2009; Ricci et al. 2017), Suzaku (Mantovani et al.
2014) NuSTAR (Brenneman et al. 2014a, b), AstroSat (Tripathi et al. 2021;
Dewangan et al. 2021). It has a 2–10 keV flux that ranges between $8.1\times
10^{-11}$ $\rm erg\,\rm cm^{-2}\,\rm s^{-1}$ (Winter et al. 2009) and
$1.8\times 10^{-9}$ $\rm erg\,\rm cm^{-2}\,\rm s^{-1}$ (Shu et al. 2010a). The
past BeppoSax, ASCA+RXTE INTEGRAL and Swift/BAT observations of IC 4329A
placed a rough constraints on the high-energy cut-off respectively at $E_{\rm
c}\geq 180$ keV (Perola et al. 2002), $E_{\rm c}=150-390$ keV (Done et al.
2000), $E_{\rm c}=60-300$ keV (Molina et al. 2013) and $E_{\rm c}>190$ keV
(Ricci et al. 2017). The NuSTAR hard X-ray spectrum of this AGN is
characterized by a power law with a photon-index of $\Gamma=1.73\pm 0.01$ and
a high-energy cut-off at $E_{\rm c}=186\pm 14$ keV (Brenneman et al. 2014b).
Past analysis showed also the presence of a neutral reflection component with
a reflection fraction of $R_{\rm refl}=0.42\pm 0.02$ (Mantovani et al. 2014).

This source is well known to exhibit the presence of a strong Fe K$\alpha$
line (Piro et al. 1990). The Fe K$\alpha$ line in IC 4329A shows a narrow core
at $6.4$ keV consistent with being produced by low-ionization material either
in the outer accretion disc, in the broad line region, or in the torus. The
narrow core component of the Fe K$\alpha$ shows also variability on timescales
of weeks to months (Fukazawa et al. 2016; Andonie et al. 2022). From the
analysis of the ASCA, RXTE, BeppoSAX, Suzaku and NuSTAR observations (Done et
al. 2000; Dadina 2007; Brenneman et al. 2014b; Mantovani et al. 2014) the
presence of a broad component was reported, most likely produced in the inner
part of the accretion disc, and blurred by general relativistic effects. In
other works (e.g. McKernan & Yaqoob 2004; Markowitz et al. 2006; Nandra et al.
2007; Tripathi et al. 2021) the observations suggest a modest or weak broad
iron line indicating that X–ray reflection from the inner disk is weak in this
source. Analysing the Suzaku and NuSTAR observations of IC 4329A with a
relativistic reflection model, Ogawa et al. (2019) inferred a very low
reflection fraction ($R\sim 3.2\times 10^{-3}$).

IC 4329A also shows the presence of an ionized warm absorber component with
$\log(\xi/\rm erg\,\rm s^{-1}\,\rm cm)$ in the range of $-1.37$ and $2.7$
(Steenbrugge et al. 2005) and also a highly ionised, i.e. $\log(\xi/\rm
erg\,\rm s^{-1}\,\rm cm)$ $=5.34\pm 0.94$, ultra fast outflowing component
(UFO, Tombesi et al. 2015), with outflow velocity $v/c=0.098\pm 0.004$. The
presence of this component was first detected by Markowitz et al. (2006) and
then confirmed by Tombesi et al. (2011).

Recently IC 4329A has been observed by the Imaging X-ray Polarimetry Explorer
(IXPE) for $\sim 500$ ks. From these observations it appears that the source
shows a $1\sigma$ confidence limit on polarization degree of $3.3\%\pm 1.1\%$
and a polarisation angle of $78^{\circ}\pm 10^{\circ}$, consistent with being
aligned with the radio jet (Ingram et al. 2023; Pal et al. 2023). These
finding favour coronal geometries which are more asymmetric and possibly
outflowing. However the coronal geometry is unconstrained within $3\sigma$
confidence level.

Here we present the results of the analysis of the X-ray broad-band spectra
from the simultaneous XMM-Newton and NuSTAR observations of IC 4329A. These
observations are part of a joint millimeter/X-ray study of this nearby source,
which also includes observations from Swift (10 observations), NICER (20
observations) and from the Atacama Large Millimeter/submillimeter Array (ALMA,
10 observations). The aim of the campaign is to study the relation between the
100 GHz and X-ray continuum in AGN, and to test the idea that the mm continuum
is associated to self-absorbed synchrotron emission from the X-ray corona
(e.g., Laor & Behar 2008; Inoue & Doi 2014), as suggested by the recent
discovery of a tight correlation between the X-ray luminosity and the 100 GHz
(Ricci et al. 2023) and 200 GHz (Kawamuro et al. 2022) luminosities of nearby
AGN. Since this paper is focused on the high-resolution X-ray spectroscopy of
IC 4329A, we will not include Swift and NICER data, which will be presented,
together with the results of ALMA campaign in a dedicated forthcoming paper
(Shablovinskaya et al., in prep).

The paper is organized as follows: in Sect. 2, we present the dataset analysed
in this work; in Sect. 3 and Sect. 4 we describe the timing and spectral data
analysis processes respectively; in Sect. 5 we discuss the results of our
analysis which are summarized in Sect. 6. Standard cosmological parameters
(H=70 km s${}^{-1}\rm Mpc^{-1}$, $\Omega_{\Lambda}$=0.73 and
$\Omega_{m}$=0.27) are adopted.

##  2 Data and Data Reduction

Table 1: Summary of the 2021 XMM-Newton and NuSTAR observations of IC 4329A. Telescope | Obs. ID | Start Date | Exp.
---|---|---|---
|  | yyyy-mm-dd | ks
XMM-Newton | 0862090101 | 2021-08-10 | 22.2
NuSTAR | 60702050002 | 2021-08-10 | 20.6
XMM-Newton | 0862090201 | 2021-08-11 | 17
XMM-Newton | 0862090301 | 2021-08-12 | 24
NuSTAR | 60702050004 | 2021-08-12 | 21.1
XMM-Newton | 0862090401 | 2021-08-13 | 20
XMM-Newton | 0862090501 | 2021-08-14 | 22.5
NuSTAR | 60702050006 | 2021-08-14 | 19.8
XMM-Newton | 0862090601 | 2021-08-15 | 17.9
XMM-Newton | 0862090701 | 2021-08-16 | 23
NuSTAR | 60702050008 | 2021-08-16 | 18.7
XMM-Newton | 0862090801 | 2021-08-17 | 19.5
XMM-Newton | 0862090901 | 2021-08-18 | 21.3
NuSTAR | 60702050010 | 2021-08-18 | 18
XMM-Newton | 0862091001 | 2021-08-19 | 21.3

The dataset analysed in this work consists of nine X-ray Multi-Mirror Mission
(XMM-Newton, Jansen et al. 2001) observations five of which have been
performed simultaneously with the Nuclear Spectroscopic Telescope Array
(NuSTAR, Harrison et al. 2013). Details on the duration and exposure of the
observations are reported in Table 1.

###  2.1 XMM-Newton data reduction

IC 4329A has been observed once per day for ten consecutive days (from
2021-08-10 to 2021-08-19) by XMM-Newton (P.I. C. Ricci) during the XMM-Newton
AO 19. The XMM-Newton observations have been performed with the European
Photon Imaging Camera (EPIC hereafter) detectors, and with the Reflection
Grating Spectrometer (RGS hereafter; den Herder et al. 2001). The EPIC cameras
were operated in small window and thin filter mode. The observation n.
0862090401 is not included in the analysis since, due to a problem in the
ground segment, the EPIC-pn exposure was lost.

The EPIC-pn camera (Strüder et al. 2001) event lists are extracted with the
epproc and emproc tools of the standard System Analysis Software (SAS
v.18.0.0; Gabriel et al. 2004). The MOS detectors (Turner et al. 2001) and the
RGS were not considered because, due to the lower statistics of their spectra,
they do not add information to the analysis.

The choice of the optimal time cuts for flaring particle background and of the
source and background extraction radii were performed via an iterative process
which maximizes the signal-to-noise ratio (SNR) as in Piconcelli et al.
(2004). The resulting optimal extraction radius was $30\arcsec$ and the
background spectra were extracted from source-free circular regions with radii
of $\sim$ $50\arcsec$ for the EPIC-pn for each observations. Response matrices
and auxiliary response files were generated using the SAS tools rmfgen and
arfgen, respectively. EPIC-pn spectra were binned in order to over-sample the
instrumental energy resolution by a factor larger than three and to have no
less than 20 counts in each background-subtracted spectral channel. No
significant pile-up affected the EPIC data, as indicated by the SAS task
epatplot. EPIC-pn light curves were extracted by using the same circular
regions for the source and the background as the spectra.

###  2.2 NuSTAR data reduction

The five NuSTAR observations have been performed, simultaneously with five
XMM-Newton observations, during NuSTAR AO 7 (P.I. C. Ricci). NuSTAR telescope
observed the source with its two coaligned X-ray telescopes Focal Plane
Modules A and B (FPMA and FPMB, respectively).

The NuSTAR high-level products were obtained using the NuSTAR Data Analysis
Software (NuSTARDAS) package (v2.1.1) within the heasoft package (version
6.29). Cleaned event files (level 2 data products) were produced and
calibrated using standard filtering criteria with the nupipeline task, and the
latest calibration files available in the NuSTAR calibration database (CALDB
20220802). For both FPMA and FPMB the radii of the circular region used to
extract source and background spectra were both $75\arcsec$; no other bright
X-ray source is present within $75\arcsec$ from IC 4329A, and no source was
present in the background region. The low-energy (0.2–5 keV) effective area
issue for FPMA (Madsen et al. 2020) does not affect our observations, since no
low-energy excess is found in the spectrum of this detector. The spectra were
binned in order to over-sample the instrumental resolution by at least a
factor of 2.5 and to have a SNR greater than 3 in each spectral channel. Light
curves are extracted using the nuproducts task, adopting the same circular
regions as the spectra.

##  3 Timing analysis

Figure 1: XMM-Newton EPIC-pn light curves (background subtracted) in the
0.2–10 keV, 0.2–2.5 keV and in the 2.5–10 keV are shown in the top, top-middle
and bottom-middle panels, respectively. In the bottom panels, the ratios
between the XMM-Newton EPIC-pn background subtracted light curves in the
2.5–10 keV and 0.2–2.5 keV energy bands are reported. All the light curves are
extracted using a binning time of 500 s. Black solid and dashed lines indicate
the average and the standard error of the mean, respectively. Figure 2: NuSTAR
FPMA+B background subtracted light curve in the 3–70 keV. 3–10 keV and 10–70
keV energy ranges are shown in the top, top-middle and bottom-middle panels,
respectively. In the bottom panels the ratios between the NuSTAR NuSTAR
background subtracted light curves in the 10–70 keV and 3–10 keV energy bands
are reported. All the light curves are extracted using a binning time of 500
s. Black solid and dashed lines indicate the average and the standard error of
the mean respectively.

During the monitoring, IC 4329A varied within some observations. However,
these variations do not affect the ratio between the flux in the 2.5–10 keV
and the 0.2–2.5 keV band, (¡ 15%, see bottom panels of Fig. 1).

We looked at the variability spectra of the XMM-Newton EPIC-pn observations of
IC 4329A using one of the common estimators for X-ray variability: the
fractional variability Fvar (Vaughan et al. 2004; Papadakis 2004; McHardy et
al. 2006; Ponti et al. 2012; Matzeu et al. 2016, 2017; De Marco et al. 2020).
The Fvar, i.e. the square root of the normalized excess variance
($\sigma^{2}_{\rm NXS}$ Nandra et al. 1997a; Edelson et al. 2002; Vaughan et
al. 2003), is the difference between the total variance of the light curve and
the mean squared error that is normalised for the average of the number of
flux measurements squared. Being N the number of good time intervals in a
light curve, and $x_{i}$ and $\sigma_{i}$ respectively the flux and the error
in each interval, Fvar is defined (Vaughan et al. 2003) as:

$F_{\rm var}=\sqrt{\sigma^{2}_{\rm
NXS}}=\sqrt{\frac{S^{2}-\overline{\sigma^{2}}}{\overline{x_{i}}^{2}}}$ (1)

Where: $S^{2}=\frac{1}{N-1}\sum_{i=1}^{N}[(x_{i}-\overline{x_{i}})^{2}]$ is
the sample variance, i.e. the integral of the power spectral density (PDS)
between two frequencies, and
$\overline{\sigma^{2}}=\frac{1}{N}\sum_{i=1}^{N}[\sigma_{i}^{2}]$ is the mean
square error.

Figure 3: Fractional variability spectra for the XMM-Newton EPIC-pn
observations of IC 4329A as in Table 1. The spectra of this campaign are
characterised by a mostly flat shape, with the fractional variability which
never exceeds the 5%. The dotted dashed black lines are the linear regressions
while the shaded regions represent the combined 3$\sigma$ error on the slope
and normalisation. Figure 4: Top panel: soft (0.2–1 keV) vs hard (3–10 keV)
$\sigma_{\rm NXS}^{2}$. Bottom panel: medium(1–3 keV) vs hard (3–10 keV)
$\sigma_{\rm NXS}^{2}$. The values are calculated for the XMM-Newton EPIC-pn
observations of IC 4329A as in Table 1. The black dashed lines represent the
one-to-one relation. The color code for different observations is the same as
in Fig. 3.

We computed Fvar for each observation using the background-subtracted EPIC-pn
light curves in different energy bands using a temporal bin of 1000 s on a
timescale of 20 ks. The errors are computed using Eq. B2 of Vaughan et al.
(2003). The resulting Fvar spectra of each observation are shown in Fig. 3. To
study the behaviour of the variability spectra, we used the linmix code, a
hierarchical Bayesian model for fitting a straight line to data with errors in
both the x and y directions (Kelly 2007). To perform the fitting we used a
linear model of the data using the following fitting relation:

$F_{\rm var}=\alpha\times E+\beta$ (2)

We found the behaviour of the variability spectra of this campaign to be
nearly flat, with a correlation coefficient $\rho_{\rm Pearson}$ spanning
between -0.05 and 0.05.

The different spectral component observed in AGN can lead to spectral
variability in different energy bands as each of them can dominate in a
certain energy band. For example the primary power law component or the
reflection component are dominant in the hard energy band (3–10 keV, Haardt &
Maraschi 1991, 1993; Haardt & Matt 1993) while soft-excess and warm-absorbers
(WA) can dominate the soft (0.2–1 keV, Bianchi et al. 2009) and medium (1–3
keV, Blustin et al. 2005; Tombesi et al. 2013) energy bands. We therefore
calculated the $\sigma_{\rm NXS}^{2}$ from the 0.2–1 keV (soft), 1–3 keV
(medium) and 3–10 keV (hard) light curves, to get a complete picture of the
X-ray variability of IC 4329A, and we compared these values in Fig. 4. The
$\sigma_{\rm NXS}^{2}$ in the soft, medium and hard energy band for the XMM-
Newton observations of IC 4329A in this campaign, in general, are very low.
They are below 0.1% in all the observations except for the $\sigma^{2}_{\rm
NXS}$ in the soft energy band of the Obs.ID 0862090101 (0.15%). Looking at the
$\sigma^{2}_{\rm NXS,1-3\,keV}$ vs $\sigma^{2}_{\rm NXS,3-10\,keV}$ values,
they are located in the vicinity or slightly below the one-to-one relation.
The same is for the $\sigma^{2}_{\rm NXS,0.2-1\,keV}$ vs $\sigma^{2}_{\rm
NXS,3-10\,keV}$ values except again for Obs.ID 0862090101 which is located
above the one-to-one relation. This could also be related to the small peak of
the variability spectrum of this observation at around 0.5 keV. Thus, overall
IC 4329A shows very weak variations. Our analysis shows that the soft-excess
and/or warm absorber components variations are weaker than those of the
primary continuum and/or reflection component on the timescale of 20 ks except
for the Obs.ID 0862090101 in which most likely the soft-excess or the
absorbing components are more variable on timescales less than $\sim 20$ ks
with respect to the continuum, in agreement with what expected for type 1 AGN
(Ponti et al. 2012; Tortosa et al. 2023).

Looking at the NuSTAR light curves, no spectral variations($<10$ %) are found
in the ratio between the 3–10 and 10–70 keV count rates (see bottom panels of
Fig. 2).

Since the hardness ratios of both XMM-Newton and NuSTAR observations do not
show strong changes within each pointing (see Sect. 3), we decided to use the
time-averaged spectrum for the forthcoming spectral analysis, to improve
spectral statistics.

##  4 Spectral analysis

Figure 5: Background-subtracted XMM-Newton EPIC-pn and NuSTAR FPMA-B spectra
and the corresponding background levels (shaded regions). The simultaneous
XMM-Newton and NuSTAR observations are reported with the same color. The color
code for different observations is the same as in Fig. 3. This color code is
applied throughout the paper.

We performed the spectral analysis using the xspec v.12.12.1 software package
(Arnaud 1996). The spectra obtained by the two NuSTAR modules (FPMA and FPMB)
are fitted simultaneously, with a cross-normalization constant typically less
than 5% (Madsen et al. 2015).

Errors and upper/lower limits are calculated using $\Delta\chi^{2}$ = 2.71
criterion (corresponding to the 90% confidence level for one parameter),
unless stated otherwise.

The Galactic column density at the position of the source ($N_{\rm
H}=4.1\times 10^{20}\rm cm^{-2}$, HI4PI Collaboration et al. 2016) is always
included in the fits, and it is modeled with the tbabs component (Wilms et al.
2000) in xspec, with $N_{\rm H}$ kept frozen to the quoted value. We also
assumed Solar abundances, unless stated otherwise.

Fig. 5 shows the background subtracted EPIC pn (0.3–10 keV), and FPMA-B
spectra (3–75 keV), plotted with the corresponding X-ray background. In all
figures, the spectra have been corrected for the effective area of each
detector.

###  4.1 XMM-Newton data analysis

We started our spectral analysis with the 2022 EPIC-pn XMM-Newton spectra (in
the 0.3–10 keV range). We fitted the spectra separately, adopting a fitting
model which was based on past literature results. The model is composed of a
neutral absorption component at the redshift of the source, a WA component
(Steenbrugge et al. 2005) and an UFO component (Markowitz et al. 2006; Tombesi
et al. 2011), xspec model: tbabs * ztbabs * mtable(WAcomp) * mtable(UFOcomp) *
(bbody + powerlaw + zgauss + zgauss). The two latter components are modelled
with detailed grids computed using the photoionization code xstar (Kallman &
Bautista 2001), with an input spectral energy distribution described by a
power law with a photon index of $\Gamma=2$. These tables consider standard
solar abundances from Asplund et al. (2009), and take into account absorption
lines and edges for all the metals characterized by an atomic number $Z\leq
30$. Considering the typical values of turbulent velocity for WAs (Laha et al.
2014), we used a table computed considering a turbulent velocity of 100 $\rm
km\,s^{-1}$ for the WA and a table computed considering a turbulent velocity
of 1000 $\rm km\,s^{-1}$ for the UFO. We added a black-body component to take
into account any weak underlying soft X-ray spectral feature, as the soft
excess. The primary continuum is modelled by a simple power law. Since we are
focusing on the XMM-Newton energy band range, we did not include reflection
features at this stage. Two Gaussian lines are included, to represent the iron
K$\alpha$ and K$\beta$ lines. The line widths of both lines are free to vary,
as well as their centroid energies. The fits obtained using this model are
very good (see left panels of Fig. 6). The values of the Fe K$\alpha$ line
parameters are shown in Table 2. The line width of the Fe K$\beta$ line is
$\sim 80-90$ eV because it is fitting a blend of lines which could include
also Fe xxvi . The line width of the Fe K$\alpha$ line in some observations is
unconstrained, and we found only upper limits. In all observations, the line
width of the K$\alpha$ line is always below $100$ eV. This suggests that if
there is a broad component of the Fe K$\alpha$ line it was not detectable
during our X-ray monitoring campaign.

Table 2: Details of the iron K$\alpha$ line parameters from the analysis of the XMM-Newton observations. The flux of the Fe K$\alpha$ line is given in unit of $\rm erg\,\rm cm^{-2}\,\rm s^{-1}$ . Obs.ID | E[keV] |  $\sigma$[eV] | EW [eV] | $\log(F)$
---|---|---|---|---
0862090101 | $6.39\pm 0.03$ | $<92$ | $64\pm 12$ | $-12.12\pm 0.08$
0862090201 | $6.36\pm 0.02$ | $65_{-32}^{+99}$ | $96\pm 15$ | $-12.03\pm 0.08$
0862090301 | $6.36\pm 0.01$ | $93_{-17}^{+132}$ | $112\pm 18$ | $-12.02\pm 0.07$
0862090501 | $6.37\pm 0.02$ | $<82$ | $46\pm 9$ | $-12.30\pm 0.08$
0862090601 | $6.36\pm 0.02$ | $81_{-42}^{+122}$ | $115\pm 23$ | $-11.97\pm 0.10$
0862090701 | $6.38\pm 0.02$ | $<48$ | $81\pm 19$ | $-12.09\pm 0.12$
0862090801 | $6.38\pm 0.02$ | $87_{-54}^{+114}$ | $104\pm 26$ | $-12.01\pm 0.07$
0862090901 | $6.35\pm 0.03$ | $69_{-35}^{+132}$ | $78\pm 29$ | $-12.14\pm 0.11$
0862091001 | $6.39\pm 0.02$ | $99_{-66}^{+131}$ | $108\pm 31$ | $-11.99\pm 0.07$

In the soft energy band (0.3–2 keV) we found the presence of a moderate soft
excess component with temperatures ranging from kT${}_{\rm BB}=47\pm 15$ eV to
kT${}_{\rm BB}=75\pm 6$ eV and a neutral absorption component at the redshift
of the source with $N_{\rm H}=(2.67\pm 0.07)\times 10^{21}\,$cm-2 over the 9
XMM-Newton observations. We also confirmed the presence of two X-ray ionized
absorbing components, consistent with previous studies (Steenbrugge et al.
2005; Markowitz et al. 2006; Tombesi et al. 2011). These components are
constant within their uncertainties during the monitoring. We found for the
X-ray wind with turbulent velocity 100 $\rm km\,s^{-1}$, identified with a
warm absorber, an average column density $N_{\rm H}=(1.49\pm 0.16)\times
10^{21}\,$cm-2 , an average ionization fraction $\log(\xi/\rm erg\,\rm
s^{-1}\rm cm)=1.30\pm 0.07$ and an average observed redshift of the absorption
components z=$-0.011\pm 0.005$. The component with turbulent velocity 1000
$\rm km\,s^{-1}$, identified with a UFO, shows an average column density
$N_{\rm H}=(1.97\pm 0.82)\times 10^{21}\,$cm-2, an average ionization fraction
$\log(\xi/\rm erg\,\rm s^{-1}\rm cm)=2.69\pm 0.03$ and an average z=$-0.177\pm
0.016$.

###  4.2 NuSTAR data analysis

Before analysing the broadband spectra, we first focused on the primary
emission and on its reflected component using NuSTAR FPMA+B spectra (in the
3–75 keV range). During the fitting process, we left the NuSTAR FPMB cross-
calibration constant free to vary. Again we composed a model following those
presented in the literature (e.g. Brenneman et al. 2014a, b). We modeled the
primary X-ray continuum with the cutoffpl model in xspec, which includes a
power law with high energy exponential cut-off, and the reprocessed emission
using the standard neutral reflection model xillver version [1.4.3] with
$\log(\xi)$ fixed to zero (García et al. 2013) (xspec model: tbabs*(cutoffpl+
xillver)). The cut-off energy, $E_{\rm c}$, and the photon index, $\Gamma$, of
the xillver component are linked those of the cutoffpl and the reflection
fraction, $R_{\rm refl}$, is forced to be negative. With these settings,
xillver reproduces only the reflection component. The disk inclination angle
was fixed to a value of $60^{\circ}$ (Brenneman et al. 2014b). The data are
very well reproduced by the model, the ratio between data and model are shown
in the central panels of Fig. 6, in which we reported also the $\chi^{2}_{r}$
= $\chi^{2}$/degrees of freedom (dof). In this analysis, we found a photon
index of the primary power law ranging from $\Gamma=1.67\pm 0.04$ to
$\Gamma=1.74\pm 0.03$, and a cut-off value $E_{\rm c}=145-180$ keV in
agreement with previous studies (e.g., Brenneman et al. 2014b). However,
respect to the results obtained by Brenneman et al. (2014b), we found a
significantly lower value for the reflection fraction, R${}_{\rm
refl}=0.008-0.01$, similar to that found by the analysis of Suzaku and NuSTAR
observations of IC 4329A by Ogawa et al. (2019). Together with the result of
our analysis of the Fe K$\alpha$ line with the XMM-Newton observations, which
shows faint broad features (see Sect. 4.1), this suggests that the reflection
from the inner disk is weak.

We also tested the scenario in which the Fe K$\alpha$ emission is slightly
broadened by relativistic effects and it is coming from the accretion disk by
substituting the xillver component with the relxill component (García et al.
2014; Dauser et al. 2014), which takes into account ionised relativistic
reflection from an accretion disk illuminated by a hot corona. This model
resulted in a significantly worse fit ($\chi^{2}_{r}\sim 1.4-1.7$). Moreover,
we were not able to constrain the black hole spin in any of the NuSTAR
observations of this campaign, finding just a lower limit of $a\gtrsim 0.1$.

Figure 6: Data/model ratio for the EPIC-pn XMM-Newton observations (left
panels), FPMA+B NuSTAR observations (central panels) and the X-ray broad-band
observations (right panels) of IC 4329A when the best fitting model is applied
in each case, see Sect. 4.

###  4.3 Broad-band data analysis

Table 3: Best fitting parameters for the X-ray broad-band XMM-Newton plus NuSTAR spectra of IC 4329A. Errors are at 90% confidence levels. Model | Parameter | Obs.ID
---|---|---
|  | 0862090101 | 0862090301 | 0862090501 | 0862090701 | 0862090901
|  | 60702050002 | 60702050004 | 60702050006 | 60702050008 | 60702050010
Model A | $\frac{\chi^{2}}{\rm dof}=\chi^{2}_{r}$ | 1.03 | 0.98 | 1.01 | 1.04 | 0.97
cutoffpl | $\Gamma$ | $1.71\pm 0.01$ | $1.65\pm 0.02$ | $1.71\pm 0.02$ | $1.73\pm 0.02$ | $1.77\pm 0.02$
cutoffpl |  $E_{\mathrm{cut}}$(keV) | $163^{+55}_{-35}$ | $140^{+56}_{-75}$ | $220^{+89}_{-45}$ | $249^{+70}_{-92}$ | $228^{+57}_{-41}$
xillver | R refl | $0.006\pm 0.001$ | $0.008\pm 0.002$ | $0.006\pm 0.001$ | $0.009\pm 0.002$ | $0.009\pm 0.001$
xillver | A Fe | $3.41^{+1.09}_{-0.89}$ | $3.72^{+2.46}_{-1.31}$ | $3.77^{+1.30}_{-1.17}$ | $3.1^{+1.28}_{-0.53}$ | $1.91^{+1.24}_{-0.81}$
Model B1 | $\frac{\chi^{2}}{\rm dof}=\chi^{2}_{r}$ | 1.06 | 1.002 | 1.03 | 1.06 | 1.005
cutoffpl | $\Gamma$ | $1.75\pm 0.02$ | $1.71\pm 0.01$ | $1.73\pm 0.01$ | $1.74\pm 0.02$ | $1.81\pm 0.02$
cutoffpl |  $E_{\mathrm{cut}}$(keV) | $>214$ | $>171$ | $>210$ | $>245$ | $>305$
mytorus | $N_{\rm H,tor}[10^{23}\rm cm^{-2}]$ | $2.84\pm 0.63$ | $2.78\pm 0.95$ | $2.67\pm 0.48$ | $2.32\pm 0.84$ | $2.48\pm 0.65$
Model B2 | $\frac{\chi^{2}}{\rm dof}=\chi^{2}_{r}$ | 1.07 | 1.01 | 1.06 | 1.11 | 1.006
RXTorusD | $\Gamma$ | $1.73\pm 0.01$ | $1.70\pm 0.02$ | $1.73\pm 0.04$ | $1.76\pm 0.01$ | $1.76\pm 0.01$
RXTorusD | $N_{\rm H,tor}[10^{23}\rm cm^{-2}]$ | $2.13^{+0.83}_{-0.38}$ | $3.70\pm 0.51$ | $2.34\pm 0.36$ | $5.10\pm 0.56$ | $3.57\pm 0.76$
RXTorusD | $i({\rm deg})$ | $44.91\pm 0.98$ | $44.66\pm 1.25$ | $44.04\pm 1.66$ | $44.83\pm 0.77$ | $45.00\pm 0.34$
RXTorusD | $r/R$ | $0.69\pm 0.02$ | $0.69\pm 0.01$ | $0.69\pm 0.02$ | $0.69\pm 0.01$ | $0.69\pm 0.03$
Model C | $\frac{\chi^{2}}{\rm dof}=\chi^{2}_{r}$ | 1.05 | 0.98 | 0.99 | 1.04 | 0.98
xillverCp | $\Gamma$ | $1.77\pm 0.01$ | $1.76\pm 0.03$ | $1.80\pm 0.02$ | $1.80\pm 0.01$ | $1.82\pm 0.02$
xillverCp | kTe(keV) | $78^{+32}_{-26}$ | $75^{+54}_{-61}$ | $>135$ | $>175$ | $>175$
xillverCp | R refl | $<3.37$ | $1.45\pm 0.62$ | $<0.82$ | $0.16\pm 0.04$ | $<0.96$
xillverCp | A Fe | $>3.74$ | $>1.59$ | $1.01^{+0.85}_{-0.25}$ | $0.87^{+0.63}_{-0.16}$ | $<0.50$
xillverCp | $\log(N_{\rm H,disk}/\rm cm^{-2})$ | $18.01^{+1.57}_{-0.64}$ | $19.18^{+0.69}_{-1.09}$ | $<16.20$ | $<17.21$ | $>15.03$
F${}_{2-10}(10^{-11}$ $\rm erg\,\rm cm^{-2}\,\rm s^{-1}$ ) | $11.04\pm 0.24$ | $7.79\pm 0.53$ | $9.09\pm 0.29$ | $8.64\pm 0.36$ | $8.49\pm 0.27$
Figure 7: Model A (see Sect. 4.3). The black dotted line represents the total
model, which is composed of the primary X-ray emission (light blue line), the
reprocessed component (dark blue line) and the soft excess component (green
line). Absorption and emission features are included in the model.

Figure 8: Unfolded data plus fitting model (upper panels) and ratio residuals
(lower panels) for the X-ray broad-band observations of IC 4329A when the best
fitting model is applied. We show all the model components as in Fig. 7, see
Sect. 4 for more details.

Finally, we extended the analysis to include the whole 0.3–10 keV XMM-Newton
EPIC-pn spectrum and the 3–75 keV NuSTAR FPMA/FPMB spectra for the
simultaneous observations (see Table 1). For the spectral fitting we used the
following xspec model: tbabs * ztbabs * mtable(WAcomp) * mtable(UFOcomp) *
(bbody + cutoffpl + xillver) (hereafter Model A, see Fig. 7) which is the
combination of the previous models. The data/model ratio of this analysis and
the resulting $\chi^{2}_{r}$ are shown in the right panels of Fig. 6. We show
data, fitting model and residuals ratio for all the simultaneous XMM-Newton
and NuSTAR observations in Fig. 12.

Even if from the analysis of the XMM-Newton data the Fe K$\alpha$ line does
not seems to be broad, we carried out additional tests to confirm this result.
The xillver model accounts for the narrow Fe K$\alpha$ emission line so we
included Gaussian line, for a broad component. Indeed this component was not
necessary as its normalisation was unconstrained, and we only found an upper
limit of $\sim 10^{-7}$ ph cm-2 s-1 in all the observations, with an upper
limit on the equivalent width in the range of 5–10 eV. We reported the best
fitting parameters for the X-ray primary continuum ($\Gamma$ and Ec) and the
reflection component (AFe and Rrefl) in Table 3.

The very low value of the reflection fraction and the faint broad component of
the Fe K$\alpha$ line in the fit of the broad-band spectra give strength to
the hypothesis in which the disk reflection is weak and the narrow iron
K$\alpha$ signature is coming from distant material (see Sect. 5.2). To
further test this scenario, we replaced the xillver with the mytorus model
(Murphy & Yaqoob 2009), to reproduce the reflection from distant neutral
material (Model B1). The inclination angle of the torus was fixed at 30∘. To
reproduce the Compton-scattered continuum we used a table111mytorus-
scatteredH160-v00.fits in which the termination energy is fixed by default in
the Model and is 160 keV, the closest value to the cut-off energy found with
Model A. To reproduce the fluorescent line spectra we used another
table222mytl-V000010nEp000H160-v00.fits with line centroid offset $E_{\rm
off}=0$ keV and with the same termination energy as the previous one. The
column densities of the scattered and line components were linked and free to
vary. The normalisation of the scattered and line components were tied to the
normalization of the primary continuum, i.e. the standard mytorus
configuration (the so-called ’coupled’ reprocessor; see e.g. Yaqoob 2012). The
assumed geometry corresponds to a covering fraction of 0.5.

The statistical significance of the fits with Model B1 is similar to that of
Model A. We found a value for the average column density of the torus of
$2.6\times 10^{23}$ cm-2 consistent within the errors with the results
obtained for all the observations of this campaign (see Table 3). The photon
index values obtained for the primary continuum are consistent within the
errors with those obtained with Model A but we were not able to constrain the
cut-off energies, most likely because of the degeneracies among the
parameters. Moreover, mytorus does not include a cut-off energy but a
termination energy which is fixed to be 160 keV. We reported the best-fitting
parameters of Model B1 in Table 3.

Additionally, we also tested the reflected emission from a dusty torus with
variable covering fraction (Model B2) using the
RXTorusD333https://www.astro.unige.ch/reflex/xspec-models model (Paltani &
Ricci 2017; Ricci & Paltani 2023), the first torus model to consider dusty
gas. In these tables the cutoff energy is fixed to 200 keV, which is a value
consistent with the value of $E_{\rm c}$ we found in IC 4329A (See Table 3).
The level of statistical significance observed in the fits using Model B2 is
comparable to that observed using Model B1. Also the values of the photon
index and the equatorial column density of the torus are comparable within the
errors with the ones of Model B (see Table 3). With Model B2 we were able to
measure the inner-to-outer radius ratio of the torus, i.e. $r/R=0.69\pm 0.02$,
and the viewing angle $i\sim 45^{\circ}$. These values were constant among the
observations. All the parameters values are reported in Table 3.

The results obtained with Model A, B1 and B2 are consistent with the iron
K$\alpha$ line originating from neutral Compton-thin ($N_{H}=10^{22}-10^{23}$
cm-2) material which does not produce a prominent Compton reflection.

We searched for possible degeneracies between the fitting parameters in Model
A, performing Monte Carlo Markov Chain (MCMC) using the xspec-emcee
tool444https://github.com/jeremysanders/xspec_emcee. This is an implementation
of the emcee code (Foreman-Mackey et al. 2013), to analyze X-ray spectra
within xspec. We used 50 walkers with 10000 iterations each, burning the first
1000. The walkers started at the best fit values found in xspec, following a
Gaussian distribution in each parameter, with the standard deviation set to
the delta value of that parameter. In the right panels of Fig. 12 the contour
plots are shown, resulting from the MCMC analysis of the Model A applied to
the broad-band 0.3–75 keV spectra of IC 4329A.

To directly measure the coronal temperature parameter we used the xillverCp
versions of the xillver tables (Model C). This model assumes that the primary
emission is due to the Comptonization of thermal disc photons in a hot corona
and the reflection spectrum is calculated by using a more physically motivated
primary continuum, implemented with the analytical Comptonization model
nthcomp (Zdziarski et al. 1996; Życki et al. 1999), instead of a simple cut-
off power law. The seed photon temperature is fixed at 50 eV by default in
this model. This value is consistent with the maximum disk temperature
expected for a source with $M_{\rm BH}=6.8^{+1.2}_{-1.1}\times
10^{7}M_{\odot}$ (Bentz et al. 2023) (i.e. $T_{\rm Wien}=32\pm 10$ eV). As
reported in Table 3, we were able to constrain the coronal temperature just
for two observations. The electron temperature of the corona is thought to be
related to the spectral cut-off energy, $E_{c}\sim 2-3\times kT_{\rm e}$
(Petrucci et al. 2000, 2001). Assuming this relation, the values we obtained
with Model C for the electron temperature of the corona are consistent with
the cut-off values that we obtained with Model A.

###  4.4 Re-analysis of the past observations

Table 4: Summary of the previous XMM-Newton , NuSTAR and Suzaku observations of IC 4329A. Telescope | Obs. ID | Start Date | Exp
---|---|---|---
|  | yyy-mm-dd | ks
XMM-Newton | 0101040401 | 2001-01-31 | 13.9
XMM-Newton | 0147440101 | 2003-08-06 | 136.1
Suzaku | 0702113010 | 2007-08-01 | 25.4
Suzaku | 0702113020 | 2007-08-06 | 30.6
Suzaku | 0702113030 | 2007-08-11 | 26.9
Suzaku | 0702113040 | 2007-08-16 | 24.2
Suzaku | 0702113050 | 2007-08-20 | 24.0
NuSTAR | 60001045002 | 2012-08-12 | 162.4
Suzaku | 0707025010 | 2012-08-13 | 117.4
XMM-Newton | 08800760801 | 2018-01-08 | 16.4
Figure 9: Top panel: The equivalent width (EWKα) of the neutral iron line
plotted against 2–10 keV X-ray luminosity ($L_{\rm 2-10}$). The black lines is
the linear regressions while the shaded black regions represent the combined
3$\sigma$ error on the slope and normalisation. Table 5: Best fitting
parameters of the primary continuum and details of the iron K$\alpha$ line
parameters for the analysis of the spectra of IC 4329A from the past XMM-
Newton and Suzaku observations as in Table 4. Errors are at 90% confidence
levels. The FeK$\alpha$ flux and the 2–10 keV flux are given in unit of $\rm
erg\,\rm cm^{-2}\,\rm s^{-1}$ .

Obs.ID | kTBB[keV] | $\Gamma$ |  $E_{\rm FeK\alpha}$[keV] |  $\sigma_{\rm FeK\alpha}$[eV] | EWFeKα [eV] | $\log(F_{K\alpha})$ | $\log(F_{2-10})$ | $\chi^{2}_{r}$
---|---|---|---|---|---|---|---|---
0101040401 | $0.08\pm 0.01$ | $1.68\pm 0.08$ | - | - | - | - | $-9.46\pm 0.55$ | 0.97
0147440101 | $0.42^{*}\pm 0.01$ | $1.71\pm 0.01$ | $6.35\pm 0.04$ | $96\pm 17$ | $87\pm 5$ | $-12.12\pm 0.03$ | $-10.03\pm 0.01$ | 1.74
0702113010 | $0.08\pm 0.05$ | $1.84\pm 0.01$ | $6.38\pm 0.03$ | $115\pm 32$ | $84\pm 20$ | $-11.98\pm 0.07$ | $-9.97\pm 0.03$ | 1.04
0702113020 | $0.03\pm 0.02$ | $1.88\pm 0.02$ | $6.40\pm 0.02$ | $82\pm 12$ | $48\pm 12$ | $-12.32\pm 0.07$ | $-9.89\pm 0.01$ | 1.10
0702113030 | $0.07\pm 0.01$ | $1.77\pm 0.01$ | $6.37\pm 0.04$ | $71\pm 6$ | $42\pm 12$ | $-12.08\pm 0.08$ | $-9.92\pm 0.03$ | 1.08
0702113040 | $0.07\pm 0.01$ | $1.85\pm 0.01$ | $6.39\pm 0.04$ | $78\pm 4$ | $47\pm 15$ | $-12.15\pm 0.10$ | $-9.96\pm 0.04$ | 1.06
0702113050 | $0.06\pm 0.01$ | $1.74\pm 0.01$ | $6.36\pm 0.05$ | $67\pm 10$ | $96\pm 21$ | $-12.41\pm 0.12$ | $-10.16\pm 0.01$ | 1.04
0707025010 | $0.06\pm 0.02$ | $1.75\pm 0.01$ | $6.36\pm 0.06$ | $77\pm 14$ | $51\pm 27$ | $-12.04\pm 0.03$ | $-9.96\pm 0.01$ | 1.10
08800760801 | $0.11\pm 0.05$ | $1.63\pm 0.03$ | $6.39\pm 0.03$ | $<140$ | $81\pm 13$ | $-11.91\pm 0.03$ | $-10.17\pm 0.24$ | 0.93

∗ The higher value of the kTBB of this Obs.ID is not related to the
variability of the soft excess but it is most likely due to the higher $N_{H}$
value of the UFO component (see Table 6 that could drive the black body shape
of the soft excess).

To better understand the origin of the Fe K$\alpha$ line, the evolution of the
spectral parameters of the primary continuum and of the soft energy components
(e.g., soft-excess or outflowing components) we compared the results of our
campaign with those obtained by previous observations of IC 4329A (see Table
4). This was done by re-analysing the archival XMM-Newton (2001, 2003, 2018),
NuSTAR (2012) and Suzaku (2007, 2012) data. In Fig. 10 we show the 2–10 keV
flux of IC 4329A during the years including the literature values from
observations carried out in the past four decades. During this period, the
2–10 keV flux of IC 4329A has changed by a factor $\sim 2.1$. The highest
state was in 2001 when the flux was F${}_{\rm 2-10}=1.75\times 10^{-10}$ $\rm
erg\,\rm cm^{-2}\,\rm s^{-1}$ while in 2001 the source reached the lowest
state with F${}_{\rm 2-10}=8.15\times 10^{-11}$ $\rm erg\,\rm cm^{-2}\,\rm
s^{-1}$ , $\sim 70$% lower than the value of 2007.

We applied the same fitting approach outlined in Sect. 4.1 to the archival
XMM-Newton observations and to the 2007 Suzaku observations, while we used
Model A (see Sect. 4.3) to fit the simultaneous NuSTAR and Suzaku observations
of 2012.

The continuum emission of all the past XMM-Newton observations and of the
Suzaku observations of 2007 is well fitted by a power law and a black-body
component, the latter being included to take into account the soft excess. The
simultaneous NuSTAR and Suzaku observations of 2012 are well fitted by Model A
(see Sect. 4.3). From the analysis of these simultaneous observations we found
a primary power law characterized by a photon index $\Gamma=1.74\pm 0.01$ and
an exponential cut-off at E${}_{c}=173\pm 13$ keV in agreement with what was
found by Brenneman et al. (2014b). We also found an iron abundance A${}_{\rm
Fe}=1.08\pm 0.22$ and a low value of the reflection fraction: R${}_{\rm
refl}=0.009\pm 0.004$, in agreement with the results obtained from our
monitoring campaign (see Sect. 4.3 and Table 3). We report the values of the
best fitting parameters of the primary continuum for the XMM-Newton and Suzaku
observation in Table 5. Data, fitting models and residuals ratio are shown in
Fig. 12. Moreover, all these observations, as well as the Suzaku observation
from 2012, show the presence of a WA and an UFO components, except the XMM-
Newton observation from 2001, for which the presence of ionized absorbers
could not be assessed due to the low SNR. The values of the parameters of
these components are reported in Table 6. Since during the 2007 Suzaku
monitoring (5 observations, each $\sim 5$ ks long) the spectral parameters
values are constant within the errors, we report the average values. For a
consistent analysis we also applied to the XMM-Newton observation of 2003 the
same model as in Tombesi et al. 2011 (i.e., a power law component absorbed by
Galactic column density and a xstar table with the turbulent velocity of 1000
$\rm km\,s^{-1}$ to the 4–10 keV band). We found values similar to those
reported in Tombesi et al. 2011: $N_{H}=(1.69\pm 0.23)\times 10^{22}$ cm-2,
$\log(\xi/\rm erg\,\rm s^{-1}\rm cm)=4.23\pm 0.82$, $v_{\rm w}/c=0.095\pm
0.004$.

Figure 10: The historical 2–10 keV light curve of IC 4329A as observed with
various X-ray satellites. In chronological order: Einstein, 1979 (Holt et al.
1989); EXOSAT, 1985 (Singh et al. 1991); Ginga, 1989 (Piro et al. 1990);
BeppoSAX, 1998 (Dadina 2007); Chandra, 2001 (Shu et al. 2010a); XMM-Newton
2001, 2003 (Nandra et al. 2007); Suzaku , 2007 (Mantovani et al. 2014);
Swift/XRT, 2007 (Winter et al. 2009); AstroSat, 2017 (Tripathi et al. 2021);
XMM-Newton 2022, medium flux of the observations analysed in this work.

Additionally, we found the values of the equivalent width (EWKα) of the
FeK$\alpha$ line obtained from the old observation and from the observations
of the 2021 campaign, to be inversely correlated with the X-ray luminosity in
the 2–10 keV band. Fitting these two parameters with a linear regression we
found a marginally significant ($2.8\sigma$) anti-correlation with Pearson
correlation coefficient $\rho=-0.664$ (see Fig. 9). This trend is known as the
‘X-ray Baldwin’ or ‘Iwasawa-Taniguchi’ effect (hereafter ‘IT effect’) (Iwasawa
& Taniguchi 1993). This inverse correlation is similar to the classical
’Baldwin effect’, in which the equivalent width of the [C IV] $\lambda$1550
emission line is inversely correlated with the UV continuum luminosity
(Baldwin 1977). The IT effect has been found in a large sample of different
objects studied with different instruments (Page et al. 2004; Jiang et al.
2006; Bianchi et al. 2007; Shu et al. 2010b; Ricci et al. 2014). Several
explanations have been proposed for the IT effect: the decrease in covering
factor of the material forming the FeK$\alpha$ line (Page et al. 2004; Ricci
et al. 2013), the dependency from the luminosity of the ionisation state of
the material which produce the line (Nandra et al. 1997b; Nayakshin 2000), the
variability related to the non-simultaneous reaction to flux changes of the
continuum of the reprocessing material (Jiang et al. 2006). One notable
distinction we found in the re-analysis of the past observations of IC 4329A
is the lack of a distinct broad component of the iron line. Excluding the XMM-
Newton observation of 2001 in which it was not possible to detect the
FeK$\alpha$ line for the low SNR ratio of the observation, the other
observations show a low value for both the line width and equivalent width of
the broad component of the FeK$\alpha$ (see Table 5).

##  5 Discussion

###  5.1 Spectral Variability

Figure 11: Top panels: $\Gamma$ vs F${}_{\rm 2-10\,keV}$ (left) and $\Gamma$
vs $\lambda_{\rm Edd}$ (right) relation for IC 4329A. The black lines is the
linear regressions while the shaded black regions represent the combined
1$\sigma$ error on the slope and normalisation. On the right panel the dotted,
dotted dashed and dashed lines represent respectively the relations from
Shemmer et al. (2008), Brightman et al. (2013) and Trakhtenbrot et al. (2017).
Bottom left panel: Ec vs $\lambda_{\rm Edd}$. Bottom right panel: Rrefl vs
$\lambda_{\rm Edd}$. The plots include the data from our campaign and data
obtained from archival XMM-Newton , NuSTAR and Suzaku observations. Ec and
Rrefl can not be measured if there are just XMM-Newton data available. The
Eddington ratio are computed using the bolometric luminosity estimated using
the bolometric correction to the 2–10 keV X-ray luminosity ($\kappa_{2-10}$)
from Lusso et al. (2010)

The continuum 2–10 keV flux of IC 4329A showed fluctuations over the past 43
years (Fig. 10) and we also observed significant variability during our
campaign. In fact the maximum value of the observed 2–10 keV flux in our
monitoring is $\sim$30% larger than the minimum observed value (see Table 3).
Consequently, we aimed to investigate whether the spectral properties of the
sources are connected to the flux level of the source. The photon index of the
power law shows some evidence for variability between the different
observations. Including the values from archival observations, we found that
$\Gamma$ shows a moderate correlation with the 2–10 keV flux with a Pearson
correlation coefficient of 0.63 corresponding to a $1-P_{\rm value}$ of 98%,
see top left panel of Fig. 11. This is in agreement with the softer-when-
brighter behavior typically observed in AGN (Sobolewska & Papadakis 2009;
Trakhtenbrot et al. 2017). Specifically, it has been found that as the 2–10
keV flux increases, the photon index tends to become steeper (Shemmer et al.
2006). This could be related to changes in the physical conditions of the
accretion disk as the flux increases (Haardt & Maraschi 1991), or to the
effect of pair production in the X-ray corona (Ricci et al. 2018).

Several studies during the past decades have shown that $\Gamma$ is strongly
correlated with $\lambda_{\rm Edd}$ (Lu & Yu 1999; Wang et al. 2004; Shemmer
et al. 2006, 2008; Brightman et al. 2013; Trakhtenbrot et al. 2017). We
checked the presence of this correlation between the values of the photon
index of IC 4392A found in our monitoring and from archival observations, with
the Eddington ratio computed using the bolometric luminosity estimated using
the bolometric correction to the 2–10 keV X-ray luminosity ($\kappa_{2-10}$)
from Lusso et al. (2010). We found a moderate $\Gamma$ vs $\lambda_{\rm Edd}$
correlation (Pearson correlation coefficient $\rho_{\rm corr}=0.62$,
corresponding to a $1-P_{\rm value}$ of 97%, see top right panel of Fig. 11).
The linear regression has a slope value of $1.69\pm 0.80$ which is consistent
within the errors with the values expected from the literature (e.g. Shemmer
et al. 2008; Brightman et al. 2013; Trakhtenbrot et al. 2017) but with a lower
normalisation value: $1.52\pm 0.11$.

This result implies that IC 4329A fits well the picture in which the mass
accretion rate is responsible for driving the physical conditions of the AGN
corona, responsible for the shape of the X-ray spectrum. Therefore the cut-off
energy is expected to show a correlation with the 2–10 keV flux/luminosity and
with the Eddington ratio. In Tortosa et al. (2018) we did not find a clear
evidence of a significant correlation between $E_{\rm c}$ and $\lambda_{\rm
Edd}$ in the sample of bright AGN observed at the time with NuSTAR . However,
the sample was still small, and did not allow the exclusion of the existence
of this kind of relation. Ricci et al. (2018), using the values from the BASS
X-ray catalogue (Ricci et al. 2017), found a negative correlation in the
$E_{\rm c}$ vs $\lambda_{\rm Edd}$ relation suggesting that the Eddington
ration is the main parameter driving the cut-off energy in AGN. The high
energy cut-off values of IC 4329A are marginally different among the
observations, but still consistent within the errors with previous literature
values (Brenneman et al. 2014b) and with the median values of bright nearby
AGN (Ricci et al. 2018). We do not observe a clear correlation in this
monitoring between this parameter and $\lambda_{\rm Edd}$ (see lower panels of
Fig. 11).

Among the different observations of our campaign, the reflection fraction is
consistent with being constant within the errors (see lower right panel of
Fig. 11). This parameter shows a very low value, consistent with what has been
observed in the analysis of Suzaku and NuSTAR observations of IC 4329A by
Ogawa et al. (2019). This result, in addition to our analysis of the Fe
K$\alpha$ line based on XMM-Newton observations, which shows faint broad
features (see Sect. 4.1), suggests that the reflection from the inner disk is
weak. This interpretation is further supported by the fact that the narrow
features of the iron line seen in the XMM-Newton observations are consistent
with being associated to emission from distant neutral material.

IC 4329A shows also a clear presence of a moderate soft excess component.
However this component does not vary among the observations. It has a
temperature ranging from kT${}_{\rm BB}=68\pm 5$ eV to kT${}_{\rm BB}=75\pm 6$
eV, consistent within the error among the observations of this campaign and
with the temperature found analysing the archival observations (see Table 4).
The flux of the soft excess component is constant among the observations of
our campaign, being $(2.31-2.89)\pm 0.75\times 10^{-4}$ ph cm-2 s-1.

###  5.2 The Fe K$\boldsymbol{\alpha}$ line

The central energy of the Gaussian line in our analysis is always consistent,
within the errors, with low-ionization or neutral material among all the
observations of this campaign (see Table2). In fact, IC 4329A is well known to
exhibit the presence of a strong Fe K$\alpha$ line (Piro et al. 1990) with a
narrow core at $6.4$ keV. A broad component, most likely produced in the inner
part of the accretion disc and blurred by general relativistic effects, has
been found in some X-ray observations (Done et al. 2000; Dadina 2007;
Brenneman et al. 2014b; Mantovani et al. 2014). Our modelling of the
observations of this campaign and of the past archival observations does not
need a relativistic broad component of the Fe K$\alpha$ line.

Using the mean value of the measured $\sigma_{\rm FeK\alpha}$, without
considering the upper limits (see Table 2) we computed a full width at half-
maximum (FWHM) of the iron K$\alpha$ line: FWH$M_{\rm FeK\alpha}=907\pm 132$
$\rm km\,/s$. With this value, we estimated the radius of the iron K$\alpha$
line emission, assuming that the line width represents the Keplerian velocity
emitting material, using the following equation:

$R_{\rm FeK\alpha}=\frac{G\,M_{\rm BH}}{v^{2}}$ (3)

where $G$ is the gravitational constant and $M_{\rm BH}$ is the central
supermassive black hole mass (Netzer 1990; Peterson et al. 2004) and assuming
a spherical geometry and an isotropic velocity distribution which results in
the assumption of $v=\sqrt{3}/2v_{\rm FWHM}$. We obtained $R_{\rm
FeK\alpha}=0.48\pm 0.14$ pc. As in Nandra (2006); Gandhi et al. (2015), we can
compare the radial position of the iron K$\alpha$ line emitting region with
the optical BLR radius RHβ, inferred from H$\beta$ reverberation studies, and
with the value of the dust sublimation radius Rsub (i.e., inner radius of the
dusty torus, Nenkova et al. 2008). Rsub is defined as:

$R_{\rm sub}=0.4\left(\frac{L_{\rm
bol}}{10^{45}\rm{ergs^{-1}}}\right)^{0.5}\left(\frac{1500\,K}{T_{\rm
sub}}\right)^{2.6}[pc]$ (4)

where $L_{\rm bol}$ is the bolometric luminosity and $T_{\rm sub}$ is the dust
sublimation temperature, assumed to be the sublimation temperature of graphite
grains ($T\sim 1500$ K, Kishimoto et al. 2007). Considering a bolometric
luminosity of $\log(L_{\rm bol}/\rm erg\,s^{-1})=45.04$ (Koss et al. 2017), we
obtain R${}_{\rm sub}=0.42$ pc while R${}_{\rm H\beta}=0.012$ pc from both
Bentz et al. 2009, 2023. Looking at these values, the bulk of the Fe K$\alpha$
emission line of IC 4329A appears to likely originate in the dusty torus. This
finding is also consistent with the value of the dust-continuum size ($0.14\pm
0.01$ pc) found by Gravity Collaboration et al. (2023).

###  5.3 X-ray outflows

Table 6: Spectral parameters, velocity, upper and lower limits on the radial location and upper limits on the energetic parameters for each X-ray absorber during the years. $v_{\rm turb}=100$ $\rm km\,s^{-1}$ | 2003 | 2007 | 2012 | 2018 | 2021
---|---|---|---|---|---
$N_{\rm H}(10^{21}\rm cm^{-2})$ | $1.35\pm 0.01$ | $1.42\pm 0.16$ | $1.03\pm 0.37$ | $1.75\pm 0.72$ | $1.49\pm 0.16$
$\log(\xi/\rm erg\,\rm s^{-1}\rm cm)$ | $1.28\pm 0.08$ | $1.54\pm 0.08$ | $1.78\pm 0.03$ | $1.70\pm 0.67$ | $1.30\pm 0.07$
$z_{\rm obs}$ | $-0.003\pm 0.001$ | $-0.016\pm 0.011$ | $-0.025\pm 0.009$ | $-0.024\pm 0.056$ | $-0.011\pm 0.005$
$v_{\rm w}/c$ | $0.02\pm 0.01$ | $0.03\pm 0.01$ | $0.04\pm 0.01$ | $0.04\pm 0.05$ | $0.03\pm 0.05$
$r_{\mathrm{min}}[10^{-2}\mathrm{pc}]$ | $1.18\pm 0.01$ | $0.63\pm 0.04$ | $0.38\pm 0.01$ | $0.40\pm 0.01$ | $0.88\pm 0.03$
$r_{\mathrm{max}}[\mathrm{pc}]$ | $1586\pm 13$ | $828\pm 93$ | $657\pm 236$ | $465\pm 164$ | $1372\pm 147$
$\dot{M}_{\rm max}/\dot{M}_{\rm acc}$ | 0.38 | 0.24 | 0.13 | 0.24 | 0.30
$\dot{p}_{\rm max}/\dot{p}_{\rm rad}$ | 0.80 | 1.27 | 1.21 | 1.38 | 1.56
$\dot{K}_{\rm max}/L_{\rm b,out}$ | 0.76 | 2.04 | 2.49 | 2.78 | 2.12
Pnull | 2.15E-6 | 3.25E-7 | 4.63E-5 | 2.74E-5 | 4.39E-6
$v_{\rm turb}=1000$ $\rm km\,s^{-1}$ | 2003 | 2007 | 2012 | 2018 | 2021
$N_{\rm H}(10^{21}\rm cm^{-2})$ | $6.68\pm 0.19$ | $2.58\pm 0.04$ | $2.29\pm 0.22$ | $1.49\pm 0.02$ | $1.97\pm 0.82$
$\log(\xi/\rm erg\,\rm s^{-1}\rm cm)$ | $2.44\pm 0.05$ | $2.15\pm 0.09$ | $2.15\pm 0.03$ | $2.02\pm 0.11$ | $2.69\pm 0.03$
$z_{\rm obs}$ | $-0.083\pm 0.001$ | $-0.076\pm 0.011$ | $-0.086\pm 0.010$ | $-0.094\pm 0.020$ | $-0.177\pm 0.016$
$v_{\rm w}/c$ | $0.102\pm 0.001$ | $0.094\pm 0.011$ | $0.105\pm 0.011$ | $0.114\pm 0.022$ | $0.207\pm 0.019$
$r_{\mathrm{min}}[10^{-4}\mathrm{pc}]$ | $6.19\pm 0.13$ | $7.22\pm 0.16$ | $5.82\pm 0.12$ | $4.97\pm 0.19$ | $1.50\pm 0.28$
$r_{\mathrm{max}}[\mathrm{pc}]$ | $55.27\pm 3.91$ | $11.94\pm 1.73$ | $126.12\pm 12.16$ | $261.48\pm 3.52$ | $42.28\pm 17.60$
$\dot{M}_{\rm max}/\dot{M}_{\rm acc}$ | 0.77 | 0.44 | 0.30 | 0.20 | 0.39
$\dot{p}_{\rm max}/\dot{p}_{\rm rad}$ | 1.63 | 2.72 | 3.38 | 5.35 | 3.78
$\dot{K}_{\rm max}/L_{\rm b,out}$ | 8.34 | 12.92 | 17.85 | 30.54 | 39.32
Pnull | 9.56E-8 | 5.34E-9 | 2.13E-9 | 7.26E-8 | 1.42E-8

During the fitting process we took into account ionized absorption by
including two xstar tables, as outlined in Sect. 4.1. The resulting parameters
for the absorption components obtained by our analysis of the 2021 campaign,
together with the results obtained for the archival observations, are listed
in Table 6.

Following the same approach outlined in different studies (e.g. Tombesi et al.
2013; Serafinelli et al. 2019; Tortosa et al. 2022), we computed upper and
lower bounds on the radial position of the absorbers. The lower limit on the
radial location of the outflows can be placed by estimating the radius at
which the observed velocity equals the escape velocity:
$r_{\mathrm{min}}=2GM_{\rm BH}v^{-2}$. The upper limit is placed using the
definition of the ionization parameter $\xi=L_{\rm ion}/nr^{2}$ (Tarter et al.
1969) where $L_{\rm ion}$ is the unabsorbed ionizing luminosity emitted by the
source between 1 Ryd and 1000 Ryd (1 Ryd = 13.6 eV), $n$ is the number density
of the absorbing material and $r$ is the distance from the central source. To
apply this definition we need to assume that the thickness of the absorber
does not exceed the distance from the SMBH, thus that the absorbers are
somewhat compact (Crenshaw & Kraemer 2012; Tombesi et al. 2013):
$r_{\mathrm{max}}=L_{\mathrm{ion}}N_{\rm H}^{-1}\xi^{-1}$ with $N_{\rm H}$ and
$\xi$ respectively the column density and the ionization fraction of the
absorber. We computed the ionizing luminosity using the luminosity task in
xspec on the unabsorbed best fit spectral model (see Sect. 4 and Sect. 4)
between 13.6 eV and 13.6 keV for all the observations.

The values of the estimated radial location of the absorber components in IC
4329A are reported in Table 6. Comparing these values with the literature it
is possible to see that the outflow component with turbulent velocity 100 $\rm
km\,s^{-1}$ (hereafter Wind 1) shows the typical upper and lower limits of the
distance of the WAs for the type 1 Seyfert galaxies (Tombesi et al. 2013)
while the outflow component with turbulent velocity 1000 $\rm km\,s^{-1}$
(hereafter Wind 2) is within the range of the average locations of ultra-fast
outflows ($\sim 3\times 10^{-4}-3\times 10^{-2}$ pc, see Tombesi et al. 2012).

Regarding the outflow energetics, we computed the mass outflow rate:
$\dot{M}_{\rm out}=4\pi rN_{\rm H}\mu m_{p}C_{g}v_{\rm w}$ (Crenshaw & Kraemer
2012), where $r$ is the radial location of the absorber, $N_{\rm H}$ is the
equivalent hydrogen column density, $\mu$ is the mean atomic mass per proton
(= 1.4 for solar abundances), $m_{p}$ is the mass of the proton, $C_{\rm g}$
is the global covering factor ($\simeq 0.5$, Tombesi et al. 2010) and $v_{\rm
w}$ is the radial-velocity centroid. It is interesting to compare this
quantity with the mass accretion rate of the source (see Table 6). The typical
value of the mass outflow rate for sources accreting below or close to the
Eddington limit is $\dot{M}_{\rm out}\gtrsim 5-10\%\dot{M}_{\rm acc}$, for
both UFOs and slower outflows (Tombesi et al. 2012). Both X-ray absorbers are
in agreement with this scenario in which the mechanical power is enough to
exercise a significant feedback impact on the surrounding environment.

We computed also the value of the momentum rate of the outflow, that is the
rate at which the outflow transports momentum into the environment of the host
galaxy: $\dot{p}=\dot{M}_{\rm out}v_{w}$. This quantity is compared in Table 6
with the momentum of the radiation of the source, which is defined as the
ratio between the observed luminosity and the speed of light, thus for IC
4329A, it is: $\log(\dot{p}_{\rm rad}/\rm erg\,\rm cm^{-1})=35.58$. Outflows
accelerated through the continuum radiation pressure are expected to have a
$\dot{p}_{\rm out}/\dot{p}_{\rm rad}\gtrsim 1$ (King & Pounds 2015). The
median value of this ratio for UFOs is $\sim 0.96$ after the relativistic
correction and $\sim 0.64$ without the relativistic corrections (Luminari et
al. 2020). The values of $\dot{p}_{\rm out}/\dot{p}_{\rm rad}$ we found for IC
4329A are in the range between $0.80$ and $5.35$. This suggests that radiation
pressure is the mechanism accelerating the material to the escape velocity.

The last parameter we computed was the instantaneous kinetic power of the
outflow: $\dot{K}=\frac{1}{2}\dot{M}_{\rm out}v_{w}^{2}$, and we compared it
with the outflowing observed bolometric luminosity of the source. Tombesi et
al. (2012) showed that in order to have a significant feedback impact in the
environment surrounding an AGN, a minimum ratio between the mechanical power
of the outflow and the bolometric luminosity of $\sim 0.3\%$ for UFOs and
$\sim 0.8\%$ for WAs is required. Looking at the values found for IC 4329A,
the source fits well in this scenario in which the outflowing winds can
impress a strong feedback.

##  6 Conclusions

Here we have presented the detailed broad-band analysis of the simultaneous
XMM-Newton \- NuSTAR monitoring of the Seyfert 1 galaxy IC 4329A carried out
in 2021. In summary the results of our analysis are the following:

  * •

Our spectral analysis shows that the X-ray broad-band spectra of IC 4329A from
the 2021 X-ray monitoring campaign are well fitted by a model which included a
soft excess component, a warm absorber (WA) and an ultrafast outflow (UFO),
primary emission modelled by a power law with a cut-off at high energy and a
distant neutral reflection component accounting for the narrow Fe fluorescence
(see Fig. 7 and Sect. 4).

  * •

The soft-excess or warm absorber components variations are weaker than those
of the primary continuum on timescales less than 20 ks.

  * •

We do not find significant variations in the XMM-Newton and NuSTAR hardness
ratios of the source (less than 10%, see Fig. 1,2 and Sect. 3) within the
single observations ($\sim 20$ ks). The variability spectra of IC 4329A during
the campaign is nearly flat (see Fig. 3). The analysis of the excess variance
values of the observations of this monitoring also suggests that, overall, IC
4329A shows very weak spectral variations (see Fig. 4).

  * •

We found a photon index of the primary power law ranging from $\Gamma=1.65\pm
0.02$ to $\Gamma=1.77\pm 0.02$, and a cut-off value $E_{\rm c}=140-250$ keV in
agreement with the results of Brenneman et al. (2014b). However, respect to
the results from Brenneman et al. (2014b), we found a lower value for the
reflection fraction, R${}_{\rm refl}=0.006-0.009$, as found by the analysis of
Suzaku and NuSTAR observations of IC 4329A by Ogawa et al. (2019). Together
with the very faint broad component of the Fe K$\alpha$ (see Sect. 4.1), this
suggests that the reflection from the inner disk is weak.

  * •

We do not find the presence of a broad component of the iron K$\alpha$ line in
IC 4329A. The presence of this component, attributed to relativistic effects,
has been a matter of debate in the past decades. It was reported in the
analysis of past ASCA, RXTE, BeppoSAX, Suzaku and NuSTAR observations (Done et
al. 2000; Dadina 2007; Brenneman et al. 2014b; Mantovani et al. 2014) while
other other works (e.g. McKernan & Yaqoob 2004; Markowitz et al. 2006; Nandra
et al. 2007; Tripathi et al. 2021) reported a modest or weak broad iron line
indicating that X–ray reflection from the inner disk is weak in this source.
Our findings are in agreements with the latter.

  * •

Small spectral changes are observed, following the softer-when-brighter
behaviour typically observed in unobscured AGN. This is most likely related to
changes in the physical conditions of the accretion flow (see Sect. 5.1). In
fact, we found a statistically significant $\Gamma$ vs $\lambda_{\rm Edd}$
correlation (see upper right panel of Fig. 11), indicating that the Eddington
ratio may drive the physical conditions of the X-ray emitting corona.

  * •

We estimated the distance from the SMBH of the iron K$\alpha$ line emitting
region from the mean FWHM of the line. We find that the Fe K$\alpha$ emission
line of IC 4329A is consistent with originating in the dusty torus (see Sect.
5.2).

  * •

From the analysis of the XMM-Newton and NuSTAR observations of IC 4329A from
this campaign and from archival data, we found the presence of X-ray outflows
composed by two phases, one ultrafast outflow and one warm absorber with lower
velocity (see Sect. 5.3). From the energetic of these X-ray winds we conclude
that both components are powered by radiation-pressure and that they could
exert a significant feedback impact on the surrounding environment.

###### Acknowledgements.

C.R. acknowledges support from the Fondecyt Regular grant 1230345 and ANID
BASAL project FB210003. E.S. acknowledges support from ANID BASAL project
FB210003 and Gemini ANID ASTRO21-0003. T.K. is grateful for support from RIKEN
Special Postdoctoral Researcher Program and is supported by JSPS KAKENHI grant
number JP23K13153. This work is based on observations obtained with the ESA
science mission XMM-Newton , with instruments and contributions directly
funded by ESA Member States and the USA (NASA), the NuSTAR mission, a project
led by the California Institute of Technology, managed by the Jet Propulsion
Laboratory and funded by NASA. This research has made use of the NuSTAR Data
Analysis Software (NuSTARDAS) jointly developed by the ASI Space Science Data
Center (SSDC, Italy) and the California Institute of Technology (Caltech,
USA).

## References

  * Andonie et al. (2022) Andonie, C., Bauer, F. E., Carraro, R., et al. 2022, A&A, 664, A46
  * Arnaud (1996) Arnaud, K. A. 1996, in Astronomical Society of the Pacific Conference Series, Vol. 101, Astronomical Data Analysis Software and Systems V, ed. G. H. Jacoby & J. Barnes, 17
  * Asplund et al. (2009) Asplund, M., Grevesse, N., Sauval, A. J., & Scott, P. 2009, ARA&A, 47, 481
  * Baldwin (1977) Baldwin, J. A. 1977, ApJ, 214, 679
  * Beckmann et al. (2006) Beckmann, V., Gehrels, N., Shrader, C. R., & Soldi, S. 2006, ApJ, 638, 642
  * Bentz et al. (2023) Bentz, M. C., Onken, C. A., Street, R., & Valluri, M. 2023, ApJ, 944, 29
  * Bentz et al. (2009) Bentz, M. C., Peterson, B. M., Netzer, H., Pogge, R. W., & Vestergaard, M. 2009, ApJ, 697, 160
  * Bianchi et al. (2007) Bianchi, S., Guainazzi, M., Matt, G., & Fonseca Bonilla, N. 2007, A&A, 467, L19
  * Bianchi et al. (2009) Bianchi, S., Guainazzi, M., Matt, G., Fonseca Bonilla, N., & Ponti, G. 2009, A&A, 495, 421
  * Blustin et al. (2005) Blustin, A. J., Page, M. J., Fuerst, S. V., Branduardi-Raymont, G., & Ashton, C. E. 2005, A&A, 431, 111
  * Brenneman et al. (2014a) Brenneman, L. W., Madejski, G., Fuerst, F., et al. 2014a, ApJ, 781, 83
  * Brenneman et al. (2014b) Brenneman, L. W., Madejski, G., Fuerst, F., et al. 2014b, ApJ, 788, 61
  * Brightman et al. (2013) Brightman, M., Silverman, J. D., Mainieri, V., et al. 2013, MNRAS, 433, 2485
  * Crenshaw & Kraemer (2001) Crenshaw, D. M. & Kraemer, S. B. 2001, ApJ, 562, L29
  * Crenshaw & Kraemer (2012) Crenshaw, D. M. & Kraemer, S. B. 2012, The Astrophysical Journal, 753, 75
  * Dadina (2007) Dadina, M. 2007, A&A, 461, 1209
  * Dauser et al. (2014) Dauser, T., Garcia, J., Parker, M. L., Fabian, A. C., & Wilms, J. 2014, MNRAS, 444, L100
  * De Marco et al. (2020) De Marco, B., Adhikari, T. P., Ponti, G., et al. 2020, A&A, 634, A65
  * De Marco et al. (2013) De Marco, B., Ponti, G., Cappi, M., et al. 2013, MNRAS, 431, 2441
  * den Herder et al. (2001) den Herder, J. W., Brinkman, A. C., Kahn, S. M., et al. 2001, A&A, 365, L7
  * Dewangan et al. (2021) Dewangan, G. C., Tripathi, P., Papadakis, I. E., & Singh, K. P. 2021, MNRAS, 504, 4015
  * Done et al. (2000) Done, C., Madejski, G. M., & Życki, P. T. 2000, ApJ, 536, 213
  * Edelson et al. (2002) Edelson, R., Turner, T. J., Pounds, K., et al. 2002, ApJ, 568, 610
  * Fabian et al. (2009) Fabian, A. C., Zoghbi, A., Ross, R. R., et al. 2009, Nature, 459, 540
  * Foreman-Mackey et al. (2013) Foreman-Mackey, D., Hogg, D. W., Lang, D., & Goodman, J. 2013, Publications of the Astronomical Society of the Pacific, 125, 306
  * Fukazawa et al. (2016) Fukazawa, Y., Furui, S., Hayashi, K., et al. 2016, ApJ, 821, 15
  * Gabriel et al. (2004) Gabriel, C., Denby, M., Fyfe, D. J., et al. 2004, in Astronomical Society of the Pacific Conference Series, Vol. 314, Astronomical Data Analysis Software and Systems (ADASS) XIII, ed. F. Ochsenbein, M. G. Allen, & D. Egret, 759
  * Gandhi et al. (2015) Gandhi, P., Hönig, S. F., & Kishimoto, M. 2015, ApJ, 812, 113
  * García et al. (2014) García, J., Dauser, T., Lohfink, A., et al. 2014, ApJ, 782, 76
  * García et al. (2013) García, J., Dauser, T., Reynolds, C. S., et al. 2013, The Astrophysical Journal, 768, 146
  * George & Fabian (1991) George, I. M. & Fabian, A. C. 1991, MNRAS, 249, 352
  * Gravity Collaboration et al. (2023) Gravity Collaboration, Amorim, A., Bourdarot, G., et al. 2023, A&A, 669, A14
  * Haardt & Maraschi (1991) Haardt, F. & Maraschi, L. 1991, ApJ, 380, L51
  * Haardt & Maraschi (1993) Haardt, F. & Maraschi, L. 1993, ApJ, 413, 507
  * Haardt et al. (1994) Haardt, F., Maraschi, L., & Ghisellini, G. 1994, ApJ, 432, L95
  * Haardt et al. (1997) Haardt, F., Maraschi, L., & Ghisellini, G. 1997, ApJ, 476, 620
  * Haardt & Matt (1993) Haardt, F. & Matt, G. 1993, MNRAS, 261, 346
  * Harrison et al. (2013) Harrison, F. A., Craig, W. W., Christensen, F. E., et al. 2013, ApJ, 770, 103
  * HI4PI Collaboration et al. (2016) HI4PI Collaboration, Ben Bekhti, N., Flöer, L., et al. 2016, A&A, 594, A116
  * Ho (2008) Ho, L. C. 2008, ARA&A, 46, 475
  * Holt et al. (1989) Holt, S. S., Turner, T. J., Mushotzky, R. F., & Weaver, K. 1989, in ESA Special Publication, Vol. 296, Two Topics in X-Ray Astronomy, Volume 1: X Ray Binaries. Volume 2: AGN and the X Ray Background, ed. J. Hunt & B. Battrick, 1105–1110
  * Ingram et al. (2023) Ingram, A., Ewing, M., Marinucci, A., et al. 2023, MNRAS, 525, 5437
  * Inoue & Doi (2014) Inoue, Y. & Doi, A. 2014, PASJ, 66, L8
  * Iwasawa & Taniguchi (1993) Iwasawa, K. & Taniguchi, Y. 1993, ApJ, 413, L15
  * Jansen et al. (2001) Jansen, F., Lumb, D., Altieri, B., et al. 2001, A&A, 365, L1
  * Jiang et al. (2006) Jiang, P., Wang, J. X., & Wang, T. G. 2006, ApJ, 644, 725
  * Kallman & Bautista (2001) Kallman, T. & Bautista, M. 2001, ApJS, 133, 221
  * Kawamuro et al. (2022) Kawamuro, T., Ricci, C., Imanishi, M., et al. 2022, ApJ, 938, 87
  * Kelly (2007) Kelly, B. C. 2007, The Astrophysical Journal, 665, 1489
  * King & Pounds (2015) King, A. & Pounds, K. 2015, ARA&A, 53, 115
  * Kishimoto et al. (2007) Kishimoto, M., Hönig, S. F., Beckert, T., & Weigelt, G. 2007, A&A, 476, 713
  * Koss et al. (2017) Koss, M., Trakhtenbrot, B., Ricci, C., et al. 2017, ApJ, 850, 74
  * Koss et al. (2022) Koss, M. J., Ricci, C., Trakhtenbrot, B., et al. 2022, The Astrophysical Journal Supplement Series, 261, 2
  * Koss et al. (2022) Koss, M. J., Trakhtenbrot, B., Ricci, C., et al. 2022, ApJS, 261, 1
  * Laha et al. (2014) Laha, S., Guainazzi, M., Dewangan, G. C., Chakravorty, S., & Kembhavi, A. K. 2014, MNRAS, 441, 2613
  * Laor & Behar (2008) Laor, A. & Behar, E. 2008, MNRAS, 390, 847
  * Lu & Yu (1999) Lu, Y. & Yu, Q. 1999, ApJ, 526, L5
  * Luminari et al. (2020) Luminari, A., Tombesi, F., Piconcelli, E., et al. 2020, A&A, 633, A55
  * Lusso et al. (2010) Lusso, E., Comastri, A., Vignali, C., et al. 2010, A&A, 512, A34
  * Madsen et al. (2020) Madsen, K. K., Grefenstette, B. W., Pike, S., et al. 2020, arXiv e-prints, arXiv:2005.00569
  * Madsen et al. (2015) Madsen, K. K., Harrison, F. A., Markwardt, C. B., et al. 2015, ApJS, 220, 8
  * Mantovani et al. (2014) Mantovani, G., Nandra, K., & Ponti, G. 2014, MNRAS, 442, L95
  * Markowitz et al. (2006) Markowitz, A., Reeves, J. N., & Braito, V. 2006, ApJ, 646, 783
  * Matt et al. (1997) Matt, G., Fabian, A. C., & Reynolds, C. S. 1997, MNRAS, 289, 175
  * Matzeu et al. (2016) Matzeu, G. A., Reeves, J. N., Nardini, E., et al. 2016, MNRAS, 458, 1311
  * Matzeu et al. (2017) Matzeu, G. A., Reeves, J. N., Nardini, E., et al. 2017, MNRAS, 465, 2804
  * McHardy et al. (2006) McHardy, I. M., Koerding, E., Knigge, C., Uttley, P., & Fender, R. P. 2006, Nature, 444, 730
  * McKernan & Yaqoob (2004) McKernan, B. & Yaqoob, T. 2004, ApJ, 608, 157
  * Mehdipour & Costantini (2018) Mehdipour, M. & Costantini, E. 2018, A&A, 619, A20
  * Merloni (2003) Merloni, A. 2003, MNRAS, 341, 1051
  * Molina et al. (2013) Molina, M., Bassani, L., Malizia, A., et al. 2013, MNRAS, 433, 1687
  * Murphy & Yaqoob (2009) Murphy, K. D. & Yaqoob, T. 2009, MNRAS, 397, 1549
  * Nandra (2006) Nandra, K. 2006, MNRAS, 368, L62
  * Nandra et al. (1997a) Nandra, K., George, I. M., Mushotzky, R. F., Turner, T. J., & Yaqoob, T. 1997a, ApJ, 476, 70
  * Nandra et al. (1997b) Nandra, K., George, I. M., Mushotzky, R. F., Turner, T. J., & Yaqoob, T. 1997b, ApJ, 488, L91
  * Nandra et al. (2007) Nandra, K., O’Neill, P. M., George, I. M., & Reeves, J. N. 2007, MNRAS, 382, 194
  * Nandra & Pounds (1994) Nandra, K. & Pounds, K. A. 1994, MNRAS, 268, 405
  * Nayakshin (2000) Nayakshin, S. 2000, ApJ, 534, 718
  * Nenkova et al. (2008) Nenkova, M., Sirocky, M. M., Ivezić, Ž., & Elitzur, M. 2008, ApJ, 685, 147
  * Netzer (1990) Netzer, H. 1990, in Active Galactic Nuclei, ed. R. D. Blandford, H. Netzer, L. Woltjer, T. J. L. Courvoisier, & M. Mayor, 57–160
  * Ogawa et al. (2019) Ogawa, S., Ueda, Y., Yamada, S., Tanimoto, A., & Kawaguchi, T. 2019, ApJ, 875, 115
  * Oh et al. (2022) Oh, K., Koss, M. J., Ueda, Y., et al. 2022, ApJS, 261, 4
  * Page et al. (2004) Page, K. L., O’Brien, P. T., Reeves, J. N., & Turner, M. J. L. 2004, MNRAS, 347, 316
  * Pal et al. (2023) Pal, I., Stalin, C. S., Chatterjee, R., & Agrawal, V. K. 2023, Journal of Astrophysics and Astronomy, 44, 87
  * Paltani & Ricci (2017) Paltani, S. & Ricci, C. 2017, A&A, 607, A31
  * Papadakis (2004) Papadakis, I. E. 2004, MNRAS, 348, 207
  * Perola et al. (2002) Perola, G. C., Matt, G., Cappi, M., et al. 2002, A&A, 389, 802
  * Peterson et al. (2004) Peterson, B. M., Ferrarese, L., Gilbert, K. M., et al. 2004, ApJ, 613, 682
  * Petre et al. (1984) Petre, R., Mushotzky, R. F., Krolik, J. H., & Holt, S. S. 1984, ApJ, 280, 499
  * Petrucci et al. (2001) Petrucci, P. O., Haardt, F., Maraschi, L., et al. 2001, ApJ, 556, 716
  * Petrucci et al. (2000) Petrucci, P. O., Haardt, F., Maraschi, L., et al. 2000, ApJ, 540, 131
  * Piconcelli et al. (2004) Piconcelli, E., Jimenez-Bailón, E., Guainazzi, M., et al. 2004, MNRAS, 351, 161
  * Piro et al. (1990) Piro, L., Yamauchi, M., & Matsuoka, M. 1990, ApJ, 360, L35
  * Ponti et al. (2012) Ponti, G., Papadakis, I., Bianchi, S., et al. 2012, A&A, 542, A83
  * Pounds et al. (1990) Pounds, K. A., Nandra, K., Stewart, G. C., George, I. M., & Fabian, A. C. 1990, Nature, 344, 132
  * Ricci et al. (2023) Ricci, C., Chang, C.-S., Kawamuro, T., et al. 2023, ApJ, 952, L28
  * Ricci et al. (2018) Ricci, C., Ho, L. C., Fabian, A. C., et al. 2018, MNRAS, 480, 1819
  * Ricci & Paltani (2023) Ricci, C. & Paltani, S. 2023, ApJ, 945, 55
  * Ricci et al. (2013) Ricci, C., Paltani, S., Awaki, H., et al. 2013, A&A, 553, A29
  * Ricci et al. (2017) Ricci, C., Trakhtenbrot, B., Koss, M. J., et al. 2017, ApJS, 233, 17
  * Ricci et al. (2014) Ricci, C., Ueda, Y., Paltani, S., et al. 2014, MNRAS, 441, 3622
  * Salpeter (1964) Salpeter, E. E. 1964, ApJ, 140, 796
  * Serafinelli et al. (2019) Serafinelli, R., Tombesi, F., Vagnetti, F., et al. 2019, A&A, 627, A121
  * Shemmer et al. (2006) Shemmer, O., Brandt, W. N., Netzer, H., Maiolino, R., & Kaspi, S. 2006, ApJ, 646, L29
  * Shemmer et al. (2008) Shemmer, O., Brandt, W. N., Netzer, H., Maiolino, R., & Kaspi, S. 2008, ApJ, 682, 81
  * Shu et al. (2010a) Shu, X. W., Yaqoob, T., & Wang, J. X. 2010a, ApJS, 187, 581
  * Shu et al. (2010b) Shu, X. W., Yaqoob, T., & Wang, J. X. 2010b, ApJS, 187, 581
  * Singh et al. (1991) Singh, K. P., Rao, A. R., & Vahia, M. N. 1991, ApJ, 377, 417
  * Sobolewska & Papadakis (2009) Sobolewska, M. A. & Papadakis, I. E. 2009, MNRAS, 399, 1597
  * Steenbrugge et al. (2005) Steenbrugge, K. C., Kaastra, J. S., Sako, M., et al. 2005, A&A, 432, 453
  * Strüder et al. (2001) Strüder, L., Briel, U., Dennerl, K., et al. 2001, A&A, 365, L18
  * Sunyaev & Titarchuk (1980) Sunyaev, R. A. & Titarchuk, L. G. 1980, A&A, 500, 167
  * Tarter et al. (1969) Tarter, C. B., Tucker, W. H., & Salpeter, E. E. 1969, ApJ, 156, 943
  * Tombesi et al. (2012) Tombesi, F., Cappi, M., Reeves, J. N., & Braito, V. 2012, MNRAS, 422, L1
  * Tombesi et al. (2013) Tombesi, F., Cappi, M., Reeves, J. N., et al. 2013, MNRAS, 430, 1102
  * Tombesi et al. (2011) Tombesi, F., Cappi, M., Reeves, J. N., et al. 2011, ApJ, 742, 44
  * Tombesi et al. (2010) Tombesi, F., Cappi, M., Reeves, J. N., et al. 2010, A&A, 521, A57
  * Tombesi et al. (2015) Tombesi, F., Meléndez, M., Veilleux, S., et al. 2015, Nature, 519, 436
  * Tortosa et al. (2018) Tortosa, A., Bianchi, S., Marinucci, A., Matt, G., & Petrucci, P. O. 2018, A&A, 614, A37
  * Tortosa et al. (2023) Tortosa, A., Ricci, C., Arévalo, P., et al. 2023, MNRAS, 526, 1687
  * Tortosa et al. (2022) Tortosa, A., Ricci, C., Tombesi, F., et al. 2022, MNRAS, 509, 3599
  * Trakhtenbrot et al. (2017) Trakhtenbrot, B., Ricci, C., Koss, M. J., et al. 2017, MNRAS, 470, 800
  * Tripathi et al. (2021) Tripathi, P., Dewangan, G. C., Papadakis, I. E., & Singh, K. P. 2021, ApJ, 915, 25
  * Turner et al. (2001) Turner, M. J. L., Abbey, A., Arnaud, M., et al. 2001, A&A, 365, L27
  * Vaughan et al. (2003) Vaughan, S., Edelson, R., Warwick, R. S., & Uttley, P. 2003, MNRAS, 345, 1271
  * Vaughan et al. (2004) Vaughan, S., Fabian, A. C., Ballantyne, D. R., et al. 2004, MNRAS, 351, 193
  * Wang et al. (2004) Wang, J.-M., Watarai, K.-Y., & Mineshige, S. 2004, ApJ, 607, L107
  * Wilms et al. (2000) Wilms, J., Allen, A., & McCray, R. 2000, ApJ, 542, 914
  * Winter et al. (2009) Winter, L. M., Mushotzky, R. F., Reynolds, C. S., & Tueller, J. 2009, ApJ, 690, 1322
  * Yaqoob (2012) Yaqoob, T. 2012, MNRAS, 423, 3360
  * Zdziarski et al. (1996) Zdziarski, A. A., Johnson, W. N., & Magdziarz, P. 1996, MNRAS, 283, 193
  * Zoghbi et al. (2012) Zoghbi, A., Fabian, A. C., Reynolds, C. S., & Cackett, E. M. 2012, MNRAS, 422, 129
  * Życki et al. (1999) Życki, P. T., Done, C., & Smith, D. A. 1999, Monthly Notices of the Royal Astronomical Society, 309, 561

##  Appendix A Spectra

Figure 12: Left panels: data, fitting model and residuals ratio for XMM-Newton
EPIC-pn (black) and NuSTAR FPMA (blue) and FPMB (magenta) spectra of the
simultaneous observations of IC 4329A when Model A (see Sect. 4.3) is applied
to the data. Right panels: 68%, 90% and 99% contour plots resulting from the
MCMC analysis of the Model A applied to the broad-band XMM-Newton and NuSTAR
0.3–75 keV spectra of IC 4329A. We show the outputs for photon index
($\Gamma$), cut-off energy ($E_{\mathrm{cut}}$[keV]), NuSTAR normalization
(nor$M_{\rm nu}$ [ph/keV/cm2/s]), iron abundance ($A_{\rm Fe}$), and
reflection fraction ($R_{\rm refl}$). We reported from top to bottom the
Obs.ID:086209101+60702050002,086209103+60702050004,086209105+60702050006,086209107+60702050008,086209109+60702050010.