Accretion Flow Properties of GRS 1716-249 during its $2016-17$ ‘failed’ Outburst

Stellar-mass black holes (SBHs) are one of the most fascinating astronomical objects to study especially in X-ray, $\gamma$-ray. Black holes (BHs) are considered to be identified when the accretion flows around them emit radiations with spectra consistent with the black hole boundary conditions. Accretion disk are formed around the BHs when they accrete matter from the binary companion(s) through Roche-lobe overflow or wind. High energy radiation bears more information about the compact objects as they are emitted from closer radius of the accretion disk. Some of the stellar mass BHs are transient in nature, and show outbursts in a gap of months to years. An outburst is generally considered to be triggered, when rise in viscosity at the outer-edge of the disk occurs (Chakrabarti 1996; Ebisawa et al. 1996). According to Chakrabarti and his collaborators (Chakrabarti et al. 2019; Bhowmick et al. 2020), accreted matter from the companion first piles up at the far location (within Lagrange point L1), known as the pile-up radius ($X_{p}$) inside the Roche lobe. Due to the accumulation of large amount of matter at $X_{p}$, viscosity rises and crosses its critical value to trigger an outburst. During an outburst, a BH generally passes through four different spectral states. They are hard state (HS), hard intermediate state (HIMS), soft intermediate state (SIMS) and soft state (SS) (see, Remillard & McClintock 2006; Nandi et al 2012). Normally, when an outburst starts, it begins in the HS. Then as the luminosity increases, it goes to the HIMS and then to SIMS. Finally it moves to the SS when luminosity of the source is the maximum. This is known as the rising phase of the outburst. Then, the luminosity starts to decrease slowly to the minimum and the source transits back to the HS in the reverse cycle. This is known as the declining phase of the outburst. In short, the spectral state transition of a BH takes place in the following sequence to form a hysteresis loop: HS (rising) $\rightarrow$ HIMS (rising) $\rightarrow$ SIMS (rising) $\rightarrow$ SS $\rightarrow$ SIMS (declining) $\rightarrow$ HIMS (declining) $\rightarrow$ HS (declining).

Depending upon the nature of outbursts, transient black hole candidates (BHCs) show two types of outburst, such as type-I and type-II (Debnath et al. 2017). Those outbursts, where all four canonical spectral states are observed in the above mentioned sequence, are known as type-I or classical outburst. If any outburst does not go to softer states (SIMS, SS), it is termed as type-II or harder type of outburst. The later outbursts are also known as ‘failed’ outbursts. Each of these states are defined based on the spectral and temporal properties. Hardness intensity diagram (HID; Belloni et al. 2005) and accretion rate ratio intensity diagram (ARRID; Jana et al. 2016) correlate the spectral and temporal properties during various spectral states. The radiation spectrum of a BH is defined by a multicolor thermal black body or disk black body (DBB) and a non-thermal power-law (PL) component. While harder states (HS, HIMS) are dominated by the non-thermal high energy radiations, the softer (SIMS, SS) states are dominated by thermal black body radiation. Quasi periodic oscillations (QPOs) are one of the most common features of a stellar mass black hole (SBH). Low frequency type C QPOs (van der Klis 2005) are most common which are generally observed in HS and HIMS (McClintock & Remillard 2005). There are also the existence of sporadic type (A and B; Casella et al. 2005; van der Klis 2005) low frequency QPOs (LFQPOs). They generally show monotonic increase and decrease in frequency in rising and declining phases respectively. In SS, generally no QPOs are seen (Nandi et al. 2012; Debnath et al. 2010, 2013).

Over the years, various models were introduced to understand physics of accretion around the BHs. Bondi flow (Bondi 1952), Standard disk model (Shakura & Sunyaev 1973), Thick disk model (Paczynski & Witta 1980) are a few of them. These models could explain the radiation spectra of a BHC successfully to some extent. However, they all were more of a special solution of a general picture. The soft component in the BH spectra is, as mentioned before, the black body radiation that comes from the Keplerian disk (Shakura & Sunyaev 1973) and the PL component comes from the hot Compton cloud (Sunyaev & Titarchuk 1980). In 1995, Chakrabarti and his collaborators came up with a general picture of accretion disk, namely the two component advective flow (TCAF) solution (Chakrabarti & Titarchuk 1995, hereafter CT95; Chakrabarti 1997), which is a solution of radiation transfer equation considering both heating and cooling effects. In this TCAF model, there are two types of accretion flows: one is the high viscous, high angular momentum Keplerian matter, another one is the low viscous, lower angular momentum sub-Keplerian component. According to this model, the sub-Keplerian component, being less viscous, moves faster than the Keplerian one and forms an axisymmetric shock close to black hole when the centrifugal force roughly becomes comparable to the gravitational force acting on it. In the post-shock location as matter slows down and moves in a sub-sonic speed, matter starts to pile up and becomes hot. This post-shock puffed-up matter acts a ‘hot’ Compton cloud as of Sunyaev & Titarchuk (1980), known as the CENtrifugal pressure supported BOundary Layer (CENBOL). The soft photons from the Keplerian disk contribute to the soft multicolor thermal DBB part of the spectra. A fraction of these soft photons are intercepted by the CENBOL and are inverse-Comptonized multiple times by the ‘hot’ electrons to become higher energy photons before they emit from the CENBOL. These Comptonized photons contribute to the non-thermal hard PL tail of the spectra.

In general, at the beginning of an outburst, we see high dominance of the PL component. When an outburst is triggered due to a sudden rise in viscosity at the pile up radius i.e., at the outer disk, both types of matter start to move towards the BH after forming accretion disk (Ebisawa et al. 1996; Chakrabarti et al. 2019). As discussed before, being low viscous and low angular momentum, the sub-Keplerian matter moves closer to the BH much faster (roughly in a free fall time) than Keplerian component (moves in viscous time), and dominates at the beginning of an outburst. We see presence of a larger CENBOL. The source stays in HS during this phase of the outburst as sub-Keplerian matter dominates. The Keplerian component is referred to as the disk rate (${\dot{m}}_{d}$) while the sub-Keplerian component is referred to as the halo rate (${\dot{m}}_{h}$) in the TCAF model. As outburst progresses, more and more Keplerian matter comes into the picture to increase cooling rate, resulting shrinking of the CENBOL size. In this phase, we see the presence of the intermediate (HIMS and SIMS) spectral states. We observe soft spectral state, when CENBOL cools down completely due to higher dominance of the Keplerian matter. The declining phase starts, when viscous effect at the outer edge is turned off. This reduces supply of matter from the pile-up radius. In this phase, we see spectral states in the reverse sequence (SS$\rightarrow$SIMS$\rightarrow$HIMS$\rightarrow$HS) as supply from the ${\dot{m}}_{d}$ reduces in a faster rate compared to the ${\dot{m}}_{h}$. Debnath and his collaborators have successfully implemented this TCAF solution as a local additive table model in HEASARC’s spectral analysis software package xspec to fit BH spectra (Debnath et al. 2014; 2015a). The model requires supply of initial guess of the four basic flow parameters: (i) Keplerian disk rate (${\dot{m}}_{d}$ in ${\dot{m}}_{Edd}$), (ii) sub-Keplerian halo rate (${\dot{m}}_{h}$ in ${\dot{m}}_{Edd}$), (iii) shock location ($X_{s}$ in Schwarzschild radius $r_{s}$), which is the boundary of the CENBOL, (iv) compression ratio ($R$), which is the ratio of the matter densities of the post to pre shock flows ($R={\rho}_{+}/{\rho}_{-}$), other than (v) mass of the BH ($M_{BH}$ in $M_{\odot}$) and (vi) normalization ($N$).

In the HS, we generally see presence of a strong shock (higher $R$) at a larger $X_{s}$. Since the QPO frequency is inversely proportional to the location of the shock ($\nu_{qpo}\propto X_{s}^{-3/2}$; Chakrabarti & Manickam 2000), we observe lower frequency QPOs in the HS. In the rising phase, as the outburst progresses, shock moves inward with reducing strength as post-shock cooling rate increases with increasing ${\dot{m}}_{d}$. We see monotonically increasing frequency of the generally observed type-C QPOs. On the HIMS to SIMS transition day, shock becomes weaker ($R\sim 1$) and we see highest frequency of the evolving type-C QPOs. These evolving QPOs are generally formed due to the resonance shock oscillation when cooling time scale roughly matches with the infall time scale (Molteni, Sponholtz & Chakrabarti 1996; Chakrabarti et al. 2015). In the SIMS, generally type-B or A QPOs are observed sporadically on and off. The origin of the SIMS QPOs are (i) due to the weak resonance of a weakly resonating CENBOL (for type-B) or a shockless centrifugal barrier (for type-A); or (ii) due to non-satisfaction of the Rankine-Hugoniot (R-H) conditions to form a stable shock (all types). Recently, Chatterjee et al. (2020) observed presence of type-C QPOs in the SIMS of MAXI J1535-571 (using AstroSat LAXPC data), originating due to non-satisfaction of the R-H conditions. No QPOs are observed in SS. In declining intermediate and hard spectral states, we also see presence of the low frequency QPOs. Similar to the rising HS and HIMS, in the declining HS and HIMS, monotonically evolving type-C QPOS (with reducing frequencies as shock moves outward) are also generally seen. Jets are associated with the hard and intermediate spectral states. CENBOL acts as the base of the jets. Jets are more active in the intermediate spectral states. Compact jets are generally observed in the harder spectral states (HS and HIMS), while discrete or blobby jets could be seen in the SIMS. Due to the quenching of the CENBOL in the SS, similar to no QPOs, no jet is observed (Chakrabarti & Nandi 2000 and references therein).

Analysis with TCAF model provides us a great understanding about the source’s intrinsic property (i.e., mass) along with its spectral and temporal properties. Spectral fit with the TCAF model, directly provide us information about the flow dynamics via two accretion rates and two shock parameters. If mass of the BH is not well known, one can also estimate its probable value from spectral fit by keeping it as a free parameter. The normalization depends on the source’s distance and inclination, which does not change significantly during an outburst. That is why $N$ generally stays almost constant, or, varies in a narrow range during an outburst. Because of the non-inclusion of the jet phenomena in the current version of the TCAF model fits file, $N$ varies in a broad range in presence of jets. It can vary in a broad range if other phenomena like disk precision, bulk motion Comptonization, etc. are present. Recently, after its (TCAF) inclusion as an additive table model in XSPEC (Debnath et al. 2014, 2015a), spectra of many BHC have been successfully fitted and described by our group (Debnath et al. 2015a,b,2017,2020; Mondal et al. 2014, 2016; Chatterjee et al. 2016,2019; Jana et al. 2016,2020a,b; Chatterjee et al. 2020, Shang et al. 2019). Estimation of mass of these sources has also been done using the TCAF model. In our group’s recent papers, X-ray contribution only from jets/outflows are also estimated using the method as described by Jana, Chakrabarti & Debnath (2017, hereafter JCD17).

The well known Galactic transient BHC GRS 1716-249 (also known as GRO J1719-24; Bharali et al. 2019) is an X-ray transient, which was discovered in 1993 by CGRO/BATSE and Granat/SIGMA telescopes (Ballet et al. 1993; Harmon et al. 1993). This source is located at a distance of $2.4\pm 0.4$ kpc (della Valle, Mirabel & Rodriguez 1994). Masetti et al. (1996) estimated the lower limit of the mass of the source, which is $\sim 4.9M_{\odot}$. According to them, it has a K-type companion star of mass $\sim 1.6M_{\odot}$ with an orbital period of $\sim 14.7$ hr. This BHC has a hydrogen column density of $\sim 0.4\times 10^{22}~{}cm^{-2}$ (Tanaka 1993). Since its last outburst in 1993, it was in a quiescence phase. In 2016, it became active and showed an outburst. On 2016 December 18, it was detected by MAXI (Masumitsu et al. 2016; Negoro et al. 2016). According to Negoro et al. (2016), during December 2016, the source had a photon index ($\Gamma$) $\sim 1.62\pm 0.06$, which suggests that the source was in harder states during that time. Using Chandra satellite data, Miller et al. (2017) reported that during February 2017, $\Gamma$ still had a low value ($\sim 1.53$), which suggests that the source was still in harder state. Although Armas Padilla & Munoz-Darias (2017) reported that the source was making a transition towards softer states during 2017 March 27 & April 4, it was reported during the 2017 May 5 & 11 epoch that the source again returned back to the harder state (Bassi et al. 2017). According to them, the source did not go through the high soft state and hence they termed the $2016-17$ outburst of GRS 1716-249 as ‘failed’ outburst.

In this paper, we intend to study the accretion flow properties of the BHC GRS 1716-249 during its recent $2016-17$ outburst by careful analysis of the spectral and temporal properties and their evolution using the phenomenological disk black body plus power-law (DBB+PL) and physical TCAF models. The paper is organized in the following way. In §2, we discuss about the data reduction and data analysis method. In §3, we show the results from the spectral and temporal analysis. Finally, in §4, we make discussions about the results and try to draw some conclusions from our work in §5.

Swift monitored the $2016-17$ outburst of the BHC GRS 1716-279 roughly on a daily basis, starting from 2016 January 28, ($\sim$ MJD 57781.00) till the end of the outburst. However, the outburst actually started in December. On 2016 December 18, MAXI detected this outburst (Masumitsu et al. 2016; Negoro et al. 2016). For our work, we make use of the archival Swift/XRT ($1-10$, $0.6-8$, $0.6-4$ and $4-10$ keV), MAXI/GSC ($2-10$ and $6-20$ keV) and Swift/BAT ($15-50$ and $15-30$ keV) data. Since we did not find the archival (both in the HEASARC and ASDC) Swift data before 2017 January 28, we make use of the on demand MAXI/GSC data from 2016 January 28. For broadband analysis, we also make use of the AstroSat/LAXPC data from the ISSDC archive and NuSTAR data from HEASARC archive. The data reduction for different satellites/instruments are done using the methods mentioned in the following sub-section.

We have studied the accretion flow properties of the source based on the temporal and spectral analysis. Because every spectrum is fundamentally dependent on the mass of the black hole which controls the thermodynamic quantities of a disk, it is possible to estimate the mass of the black hole from each spectrum if the fit is proper. TCAF fits the data in absolute sense up to a normalization (a function of distance). Thus we have also obtained the probable mass of the source from TCAF model while fitting of the spectra. Our results are presented below.

The Galactic transient BHC GRS 1716-249 was monitored in multi-wave band by Swift roughly on a daily basis starting from MJD 57781.00 (2017 January 28). The data of the initial rising phase of the 2016-17 outburst is missing as Swift started pointing the source few days after its discovery. MAXI also monitored the outburst. We have performed spectral analysis of GRS 1716-249 during its $2016-17$ outburst using archival data of Swift/XRT ($0.6-8$ keV), MAXI/GSC ($6-20$ keV) and Swift/BAT ($15-30$ keV). To study in the higher energy bands, we use two observations of NuSTAR (on MJD 57850.60; 2017 April 07 and 57853.69; 2017 April 10) and also two observations of AstroSat/LAXPC (on MJD 57799.68; 2017 February 15 and 57849.69; 2017 April 06). Combining NuSTAR or LAXPC data with simultaneous XRT observations, a broad energy study is done in the range of $0.6-77$ keV and $0.6-80$ keV bands respectively. The spectral analysis is done using both phenomenological (DBB+PL) and physical (TCAF) models. The timing analysis is done mainly using XRT ($1-10$ keV) and LAXPC ($3-80$ keV) data.

Low frequency QPOs are one of the most common phenomena in hard and intermediate spectral states of stellar mass BHCs. For QPO analysis, we have made use of $0.01$ sec time binned light curves from both XRT and LAXPC data in $1-10$ keV and $3-80$ keV energy bands respectively. With Lorentzian model fitting, we have extracted the centroid frequencies of the QPOs. Using XRT data, we have found QPOs in 4 observations on MJDs 57781.00 (2017 January 28; obs. Id.=00034924001), 57806.57 (2017 February 22; obs. Id.=00034924019), 57808.37 (2017 February 24; obs. Id.=00034924020) and 57821.53 (2017 March 09; obs. Id.=00034924022). The power density spectra, made from XRT data, has low signal to noise ratio. We think that there might be existence of LFQPOs on other dates where no prominent QPOs are observed due to the noise. We also checked PDS from LAXPC data on MJDs 57799.68 (2017 February 15) and 57849.69 (2017 April 06). On both these observations, we found existence of QPOs. The PDS from LAXPC data are much less noisy than that of XRT. In Fig. 2, we show model fitted PDS of the first LAXPC observation in the frequency band $0.2-10$ Hz. This Figure shows a QPO with a harmonic, which are fitted using the combination of Lorentzian and power-law models. Using all the QPO information from both set of data, we show that $\nu_{qpo}$ increased monotonically in the rising phase (see, Fig. 3). We have not found the existence of QPOs in the declining phase. To find the origin of these QPOs i.e., they are due to satisfaction of the resonance oscillation of the shock or not (see, Molteni et al. 1996; Chakrabarti et al. 2015), we calculated cooling time ($t_{c}$)and infall time ($t_{i}$) scales of the post shock matter and found that their ratio ($t_{c}/t_{i}$) deviates largely from the unity. So, we conclude that during the both rising and declining phases of the outburst, resonance shock condition was not satisfied. Possibly due to this, in the declining phase of the outburst, we have not observed any QPO.

We have also shown the energy dependence of power density spectrum using AstroSat/LAXPC data. We have observed the presence of both the fundamental QPO and the harmonic appear in energies $<40$ keV, while there is no sign of harmonic above 40 keV. The QPO signature becomes weaker in the PDS above 60 keV. It is evident from the Figure 4(a-d) that as the energy is increasing the QPO shape is changing along with reducing rms. With increasing energy, the QPO rms is decreasing while the broadband noise is increasing. The disappearance of QPO shape in high energy (above 60 keV here) in the LAXPC data could be due to the reason that there is presence of low number of photons in high energy. It also could be due to the fact that the detector effective area decreases to $4100-2200~{}cm^{2}$ (Antia et al. 2017) in the energy range of $60-80$ keV. This was also reported by Sreehari et al. (2019) for the BHC MAXI J1535-571. Interestingly, we see one more low frequency QPO $\sim 0.62$ Hz in the energy range above $40$ keV. This QPO becomes more significant in the $60-80$ keV band as compared to $40-60$ keV band. This has an rms of $\sim 3.6\%$ in $60-80$ keV band.This could be one more local oscillation in the higher radius.

From the spectral analysis with the TCAF model, we see that at the beginning of the outburst, the supply of both the high viscous Keplerian and low viscous sub-Keplerian matter stays low. According to the TCAF model, due to higher radial velocity of the sub-Keplerian component, it moves much faster than the Keplerian component (moves in viscous timescale). The Keplerian component, having high angular momentum forms the disk, with sub-Keplerian matter residing over it, forming the halo. This halo component forms the hot Compton cloud or CENBOL at the post-shock region. Due to inverse Comptonization, the soft photons from the disk gains energy and produced the power-law tail in the observed spectrum. At the beginning of the outburst, the shock was located far away ($\sim 240r_{s}$) and the cooling was inefficient. As a result the shock was strong. The shock then gradually moved closer and reached $\sim 150r_{s}$ after MJD $\sim$ 57800 (2017 February 16). This means that the Keplerian matter moved closer and the cooling was increased. During this time, the total accretion (disk + halo ) rate was increased. Although the halo rate was decreasing from the start, it did not decrease that much during this period of observations (from $\sim 0.49$ to $\sim 0.47$). While the disk rate steadily increased during this period, the halo rate was always dominating. Also, during this time, the Photon index ($\Gamma$) was also on the lower side ($1.25-1.66$) of its values, with $T_{in}$ decreasing from $0.74$ keV to $0.58$ keV gradually. The ARR (Fig. 6d) was also higher ($\sim 1.4-2.1$) during this period of the outburst. Since the total rate continued to increase in this period, we term it as the rising phase. The shock location was stalling at $\sim 150r_{s}$ after 2017 April 07 (MJD 57850.37) for some days until 2017 April 21 (MJD 57864.07). From MJD=57864.07, we observe outward movement of the shock, when the total rate was maximum. The disk rate also reached to its peak just one day prior to this on 2017 April 14 (MJD 57857.02). It was still near its peak on MJD=57864.07. This is when we believe that the declining phase started. After this day, the disk rate started to decrease with already decreasing halo rate. Due to this, the efficiency of cooling of the CENBOL decreased and the boundary i.e., $X_{s}$ started to move outward. On 2017 May 05 (MJD 57878.74), $X_{s}$ became $\sim 194r_{s}$ and then moved away rapidly from the source. On MJD=57857.02, the photon index ($\Gamma$) of power-law model was on its highest value ($\sim 1.76$). ARR became $\sim 1.3$ during this period. After MJD=57878.74, as the supply of Keplerian matter decreased further, the $X_{s}$ continued to move outside and reached to a similar value ($\sim 240r_{s}$) as in the beginning of the outburst. After MJD=57878.74, $\Gamma$ again became $<1.6$ and both the ${\dot{m}}_{d}$ and ${\dot{m}}_{h}$ further decreased, as the source progressed towards the quiescence. ARR did not increase that much except only at the back end of the outburst when it became $\sim 1.7$ on MJD=57975.27 (2017 August 11). Since the supply of the Keplerian rate did not become very high during the outburst, the CENBOL was not cooled down completely during the middle phase of the outburst. As a result it started to move outside when supply resides and ultimately source did not reach intermediate and softer states. Due to the absence of softer states, the $2016-17$ outburst of GRS 1716-249 can be referred to as a harder type (type-II) or ‘failed’. According to Bharali et al. (2019), the orbital period of the source is $14.7$ hrs. So, this nature of the non-observation of the softer spectral states is consistent with other short orbital period objects, such as MAXI J1836-194 (Jana et al. 2016), Swift J1753.5-0127 (Debnath et al. 2017), XTE J1118+480 (Chatterjee et al. 2019). A similar presence of only HS during 2000 & 2005 outbursts of the short orbital period ($\sim 4.1$ hrs) transient BHC XTE J1118+480 was noticed (see, Chatterjee et al. 2019; Debnath et al. 2020).

We also estimated the probable mass of the BH ($M_{BH}$), from our spectral analysis with the TCAF model. In TCAF, $M_{BH}$ is a model input parameter. Since, mass of GRS 1716-249 is not well known, while fitting spectra we kept $M_{BH}$ as free. We observed $M_{BH}$ in the range of the source in between $4.50-5.93M_{\odot}$. The observed variation of our estimated mass is mainly due to instrumental and fitting errors. The flux from black body emission depends on the fourth power of the temperature ($T$) of the Keplerian disk. So any small change in the value of $T$ while fitting the black body component (which is possible due to poorly understood absorption due to intervening medium) can lead to a significant change in the estimated mass. Taking the average, we conclude that the most probable mass of the BHC GRS 1716-249 is $5.02^{+0.91}_{-0.52}M_{\odot}$, which is in good agreement with the earlier findings (Masetti et al. 1996).

We have done spectral and temporal analysis of the BHC GRS 1716-249 during its recent $2016-17$ outburst. To study outburst profile in different energy bands and hardness ratios, we used MAXI/GSC, Swift/BAT and Swift/XRT data. The temporal analysis of the QPOs is done using the Swift/XRT and AstroSat/LAXPC data using powspec of XRONOS package. We have also studied energy dependent QPOs using the LAXPC data. The spectral analysis is done using Swift/XRT, MAXI/GSC, Swift/BAT, AstroSat/LAXPC and NuSTAR/FPMA data, using both the phenomenological DBB+PL and physical TCAF models separately in XSPEC. Depending on our detailed analysis of GRS 1716-249 during 2016-17 outburst, we may summarize following results -

(i) The source showed low frequency QPOs during the rising phase of the outburst. These QPO frequencies were found to rise monotonically. QPOs were not originated due to the resonance shock oscillation of the CENBOL.

(ii) The nature of the QPOs are found to be energy dependent. We found that the fundamental QPO is absent above 60 keV, while the harmonic is non-detectable above 40 keV in the case of AstroSat/LAXPC data. As the energy increases the rms decreases along with the increment of the broadband noise.

(iii) The $2016-17$ is a ‘failed’ or harder type of outburst which did not go through the usual soft and intermediate spectral states. The source was observed only in the hard state (HS) during entire period of the outburst. This non-observation of the softer spectral states is similar to what was observed in other short orbital period sources.

(iv) The probable mass of the black hole was estimated to lie in the range of $4.50-5.93M_{\odot}$ or $5.02^{+0.91}_{-0.52}M_{\odot}$, from our spectral analysis with the TCAF model.

This work has made use of the archival XRT and BAT data provided by UK Swift Science Data Centre at the University of Leicester, MAXI GSC data provided by RIKEN, JAXA, and the MAXI team. We have also made use of LAXPC data from Indian Space Research Organization’s (ISRO’s) successful operation of AstroSat mission. We have used FPMA data from NuSTAR mission, a project led by Caltech, funded by NASA and managed by NASA/JPL. We have utilised the NuSTARDAS software package, jointly developed by the ASDC, Italy and Caltech, USA.

K.C. acknowledges support from DST/INSPIRE Fellowship (IF170233). Research of D.D. and S.K.C. is supported in part by the Higher Education Dept. of the Govt. of West Bengal, India. D.C. and D.D. acknowledge support from DST/SERB sponsored Extra Mural Research project (EMR/2016/003918) fund. A.J. acknowledges post-doctoral fellowship of Physical Research Laboratory, Ahmedabad, funded by the Department of Space, Government of India. S.N., D.D. and S.K.C. acknowledge partial support from ISRO sponsored RESPOND project (ISRO/RES/2/418/17-18) fund. R.B. acknowledges support from CSIR-UGC NET qualified UGC fellowship (June-2018, 527223).