On the formation of solar wind & switchbacks, and quiet Sun heating

The temperature of the solar atmosphere varies from $\approx 5500$ K at the photosphere to $\geq 10^{6}$ K in the corona. The intervening chromosphere variably lies at around $\approx 10^{4}$ K, with a steep rise occurring in the transition region (TR). The solar atmosphere further extends outward from the corona, filling the heliosphere with the streaming solar wind. All these layers of the solar atmosphere interact through a continuous exchange of mass and energy, and are tightly coupled by the highly dynamic magnetic field. Various processes have been proposed, and in some cases, observed to be the drivers of this mass and energy transport. The two primary processes thought to be responsible are dissipation of MagnetoHydroDynamic (MHD) waves (see e.g. Alfvén, 1947), and that of currents produced due to the braiding of magnetic field lines (see, e.g. Parker, 1972; Chiueh & Zweibel, 1987; Parker, 1988, and also Klimchuk (2006); Parnell & De Moortel (2012) for comprehensive reviews).

In the solar corona, we usually find three morphologically different regions, viz. Coronal Holes (CHs), Active Regions (ARs) and Quiet Sun (QS). The CHs are features seen as dark structures in Extreme UltraViolet (EUV) and X-ray images of the solar corona, while ARs are seen as localized bright structures. In addition to these dark and bright structures, the omnipresent background over which the CHs and ARs are observed is called the QS. Note that while CHs are clearly distinguishable from QS in the EUV and X-ray images, these two regions appear extremely similar at lower heights viz the chromosphere and photosphere (see e.g., Stucki et al., 2000, 1999; Kayshap et al., 2018; Tripathi et al., 2021; Upendran & Tripathi, 2021). However, note that the He I 10830 Å (an absorption line) shows excess intensity in CHs (Harvey & Sheeley, 1977; Kahler et al., 1983), while the He I 584 Å (an emission line) shows lower intensity (Jordan et al., 2001), and excess blueshift, line width in the CHs, over QS (Peter, 1999). However, these differences may be attributed to the sensitivity of these lines to coronal radiation, reflecting conditions in the corona. Furthermore, at 17 GHz in microwave, CHs are found to be brighter than QS (Gopalswamy et al., 1999), while this difference is not observed in radio wavelengths at 1.2 mm (Brajša et al., 2018). Thus, a gross differentiation of a given region in QS or CH is markedly seen predominantly in the coronal observations.

Observations by Krieger et al. (1973) indicated that the high velocity streams of the solar wind may be traced back to the CHs on the Sun. The “slow” wind, on the other hand, has been variously traced to the edges of ARs (Brooks et al., 2015) and equatorial CHs (see e.g. Bale et al., 2019). Similar results have been obtained by various authors (see e.g. Wang & Sheeley Jr, 1990; Schwenn, 2006; Janardhan et al., 2008) and more recently through the application of Deep learning based localization by Upendran et al. (2020). These results demonstrate that CHs are potential source regions of the solar wind and thus provide an opportunity to investigate the physical processes involved in the formation of solar wind. Since CHs and QS hardly differ at the lower atmospheric heights, studying them in tandem may allow us to explore the possibility of a unified explanation of heating in the QS and CHs, including the formation of solar wind Tripathi et al. (2021).

Hassler et al. (1999) investigated differences between CH and QS using the Si II and Ne VIII lines, and found a relation between the network regions and blueshifts of Ne VIII, with more blueshifts in CHs. Comprehensive studies of CHs and QS were undertaken by Stucki et al. (1999, 2000) using spectral lines sensitive to a range of temperatures from $\approx 8\times 10^{3}$ K to $\approx 1.4\times 10^{6}$ K. While CHs showed a clear deficit in intensity, excess blueshift and excess line width with respect to QS for spectral lines forming at a temperature higher than $\approx 4\times 10^{5}$ K, at chromospheric temperatures the differences were negligible, and within the measurement error. Similarly, Xia et al. (2004) studied the relationship between Doppler shifts of C II, H I Ly$\beta$, and O VI in CHs, and found a direct relation between the Doppler shifts of O VI with that of C II and H I Ly$\beta$. These correlated shifts led Xia et al. (2004) to conclude that these shifts are signatures of solar wind in the chromosphere. However, note that the associated uncertainties in the velocity scatter obtained by Xia et al. (2004) were large. Moreover, while the average chromospheric velocities in bins of the O VI velocities were studied, the systematic associations between red and blue shifted pixels, separately, for these lines, was not performed by Xia et al. (2004).

The correspondence between network region and outflows in the CHs, using Ne VIII line, demonstrated by Hassler et al. (1999) was further investigated by Tu et al. (2005), by mapping the formation heights of Si II, C IV and Ne VIII in a CH. On further detailed investigation, Tu et al. (2005) showed a clear relation between the Ne VIII blueshifts and the underlying magnetic field configuration, obtained using the potential field extrapolation of photospheric magnetic field. Thus, Tu et al. (2005) suggested a modulation of the solar wind velocities due to the underlying magnetic field configuration.

More recently, Kayshap et al. (2018) investigated the intensity differences between CH and QS in the Mg II k line, observed by the Interface Region Imaging Spectrometer (IRIS, De Pontieu et al., 2014). They find a clear deficit of intensity in CHs over QS for regions with similar absolute photospheric magnetic flux density ($|$B$|$), and with larger difference for larger $|$B$|$. Similar analysis for the intensity, velocity and non thermal widths for Si IV was performed by Tripathi et al. (2021) (henceforth referred to as Paper I). Similar to the results of Kayshap et al. (2018), intensity deficit in CHs over QS for regions with similar $|$B$|$ was observed. Moreover, CH (QS) was more blueshifted (redshifted) over QS (CH) for identical $|$B$|$. However, no significant difference was observed in the non-thermal width between CH and QS. The excess CH blueshifts were interpreted to be signatures of nascent solar wind at Si IV formation heights in Paper I. Thus, while a clear signal of solar wind was reported in the hotter Ne VIII line by Tu et al. (2005), the signatures are already present in the upper TR line Si IV, if the underlying photospheric magnetic flux density distribution is taken into account. Furthermore, since the regions with identical $|$B$|$ were compared, the deficit in intensity in CHs over QS would mean energy to be either used to accelerate the solar wind, or heat up the corona. Thus, a unified picture of solar wind formation & coronal heating was presented in Paper I.

In a companion paper (Upendran & Tripathi, 2021, hereafter referred to as Paper II), we perform investigation, similar to Paper I, on the C II $1334$ Å line. We find intensity deficit and excess blueshifts in CHs over QS for regions with identical $|$B$|$, similar to the results from Paper I in Si IV line. However, we also find excess redshifts in CHs over QS for regions with identical $|$B$|$, which is the opposite of what is observed in the Si IV line. Finally, we find the total line width to be larger in CH over QS for regions with similar $|$B$|$, while the line skew & kurtosis were found to be similar in both CH and QS, suggesting that similar physical processes are at work in the two regions, as was also concluded in Paper I.

All these investigations, taken together, raise several questions: Where does the solar wind actually originate? Can a single height be even ascribed to it? At what height does the differentiation between CH and QS start? Does the solar wind emergence have any relation to the origin of a million degree Kelvin corona? These questions are critical to diagnose the formation signatures of the solar wind. Furthermore, attempting to answer these questions will also narrow down on the viability and contribution of various heating mechanisms in the solar atmosphere.

Keeping the above questions in mind, we first study the Mg II h & k line dynamics in CH and QS. We then go ahead and explore the correlations between the Mg II, C II and Si IV lines in CH and QS, with the intention to explain the observations. The remainder of the paper is structured as follows: In §2, we describe our observations, with feature extraction in §2.1. In §3, we present results of the Mg II lines, while we recapitulate results for the C II and Si IV lines from Paper II and Paper I in §4 and §5. In §6, we summarize and discuss our results. Finally, we provide an interpretation on context of origin of solar wind, switchbacks (Bale et al., 2019) and QS coronal heating §8.

In this study, we use the observations recorded by IRIS, the Atmospheric Imaging Assembly (AIA; Boerner et al., 2012) and the Helioseismic and Magnetic Imager (HMI; Scherrer et al., 2012). AIA and HMI are both on board the Solar Dynamics Observatory (SDO; Pesnell et al., 2012). IRIS observes the Sun in two modes, viz. slit-jaw imaging (SJI) and spectroscopy. In the SJI mode, IRIS provides photometric context images in NUV and FUV with a pixel size of $\approx 0.16$″, and an approximate cadence of $\approx 63$ s. We consider the SJI data centered around 1330 Å and 2796 Å for co-alignment purposes. The spectroscopy is performed in three windows, one in near ultraviolet (NUV) from 2782.7 to 2851.1 Å, and two in the far ultraviolet (FUV), from 1332 to 1358 Å (FUV 1), and from 1389 to 1407 Å (FUV 2). These have a pixel size of $\approx 0.16$″ along the slit, and sample at $\approx.33$″ across the field of view (FOV). The spectral pixel size in these rasters is $\approx 25.9$ mÅ. Time cadence between successive slit positions is $\approx 30$ s. For further details on IRIS, see De Pontieu et al. (2014).

AIA observes the Sun’s atmosphere in UV and EUV using eight different passbands sensitive to plasma at different temperatures (O’Dwyer et al., 2010; Boerner et al., 2012). Here, we use the 193 Å images to distinguish between CHs and QS, and the 1600 Å images to co-align the IRIS, AIA and HMI observations, as described in Paper II. We obtain the information on the photospheric absolute magnetic flux density (i.e. $|$B$|$) from the line-of-sight (LOS) magnetograms obtained with HMI. The AIA images are taken with a pixel size of $\approx$0.6″, with the EUV images a time cadence of $\approx 12$ s, while the Ultraviolet images are taken at $\approx 24$ s cadence. HMI obtains the $B_{LOS}$ magnetograms at $\approx 45$ s cadence with a pixel size of 0.5″.

For our study, we have analyzed five different sets of observations recorded by IRIS in spectroscopic mode. The main criteria used to select these observations are that the raster must include CH and QS within the same FOV, and that they must be taken within latitude and longitude of $\pm$60^∘. The IRIS observation details are given in Table. 1. Note that the same data-set have been analyzed in Paper II to the study the properties of C II lines in CHs and QS. Out of these, three of the observations viz. DS1, DS2 and DS5 were also studied in Paper I to characterize the similarities and differences in QS and CHs in TR using Si IV line. We use corresponding coordinated AIA data cubes, with cutouts from the full disk data used from HMI.

The Mg II h & k lines form near 2803.53 Å and 2796.35 Å, respectively, while the C II and Si IV lines form near 1334.53 Å and 1393.755 Å, respectively. The Mg II and C II lines form in an optically thick chromosphere under non local thermodynamic equilibrium conditions (see, for e.g. Leenaarts et al., 2013; Rathore et al., 2015). Thus, these lines show extremely complex features and have non-trivial associations with local plasma properties. They have been explored in detail in Paper II, and Rathore et al. (2015); Leenaarts et al. (2013). For all practical purposes, the Si IV line, forming in QS TR, can be considered to be formed in optically thin conditions (Tripathi et al., 2020; Gontikakis & Vial, 2018) and its properties in QS and CH are studied in detail in Paper I.

The properties of C II lines as a function of $|$B$|$ in CHs and QS using exactly the same five dataset have been studied in great detail in Paper II. Here we summarize the salient conclusions from Paper II, which are relevant to this paper. The relevant results are graphically summarized in Fig. 11, which shows the variation of intensity (panel a), signed average velocity (panel b), average velocity in upflowing (panel c) and downflowing pixels (panel d) as a function of $|$B$|$. These variations are by and large similar to the results obtained from the analysis of Mg II lines in §3. The intensity, average velocity, upflow and downflow pixels show variation similar to Mg II with $|$B$|$. The differences in CHs and QS in all parameters here are similar to the differences seen in Mg II features.

There are a few notable differences between the properties in C II and Mg II lines, though. The QS shows excess average redshifts for $|$B$|$$\leq 30$ Gauss, while the CHs show excess average redshifts for larger $|$B$|$ in the C II line (see Fig. 11.b). Such a relation is not observed for the Mg II line (see Fig. 10). This arises due to consistent redshifts in both CHs and QS for $|$B$|$$\leq 35$ Gauss from Fig. 11.d. Note furthermore that the upflow-$|$B$|$ relation in C II is weaker than in Mg II features.

We now present the results from the intensity and Doppler shift of the Si IV line for all the datasets. The only difference between results presented here and those from Paper I is the inclusion of two additional observations in this work, thereby increasing the statistical significance of the results. We obtain the Si IV line parameters by fitting the spectra with a single Gaussian and constant continuum. The relevant results are graphically summarized in Fig. 12. These results are in complete agreement with those obtained in Paper I. Moreover, the results obtained for Si IV bear some similarities with those obtained for C II and Mg II lines.

The intensity (Fig. 12.a) and blueshift (Fig. 12.c) differences between CHs and QS, and their relations with $|$B$|$ are consistent across all three lines, though more enhanced for Si IV line. However, the signed average velocities clearly indicate reduced average redshifts in CHs over QS (Fig. 12.b). These average redshifts increase with $|$B$|$, and saturate at $\approx 40$ Gauss. The CHs show excess blueshifts over QS (Fig. 20.c), with the variation similar to those exhibited by the C II line (Fig. 11.c). The redshifted pixels alone also show a direct relation to $|$B$|$, but the QS is more redshifted than CHs (Fig. 12.d). Note also that the upflow and downflow velocities obtained for Si IV are much larger than those inferred from Mg II (Fig. 10) or the C II lines (Paper II).

The velocities and intensities show a highly non-trivial relation as a function of the formation height of different spectral lines. Therefore, to investigate if there is any correlation between Doppler signatures observed in the three different spectral lines viz. Mg II, C II and Si IV, we consider the approximate formation height of these lines obtained from numerical simulations. It has been suggested that on average C II lines form slightly higher in the atmosphere than the Mg II lines. Within the Mg II lines, the k line forms higher than the h line. Moreover, it has also been found that the line cores of both k & h lines forms higher than their respective peaks (Leenaarts et al., 2013; Rathore et al., 2015). The Si IV line forms in optically thin conditions, so ascribing an exact formation height is not possible. However, it forms at a higher temperature in the TR. We may, therefore, ascribe a greater height to Si IV than the Mg II and C II lines. With this prior, we may assume that the formation height (ascending order) is approximately Mg II h2 $\leq$ Mg II k2 $<$ Mg II h3 $\leq$ Mg II k3 $\approx$ C II $<$ Si IV.

The obtained velocities in different line features of Mg II (Fig. 10. b, c, e & f), C II (Fig.11. c & d) and Si IV (Fig. 12. c & d) clearly show that the velocity magnitude increases with increasing formation height. Considering mass flux conserving flows (Avrett et al., 2013), and that the density decreases as a function of height in the solar atmosphere, it is plausible to hypothesize that the upflows (downflows) at lower (greater) heights are enhanced (reduced) while traveling towards greater (lower) heights.

To check this hypothesis, we investigate the correlations between Mg II, C II and Si IV velocities. Since Mg II and C II form at approximately the same height, we expect these two lines to have similar properties vis-à-vis the Si IV line. In Paper I, it is suggested that the increase in Si IV blueshift with increasing with $|$B$|$ may indicate the signatures of the solar wind emergence. This motivates us to explore if the observed flows in chromosphere detected in Mg II and C II lines are in any way related to those obtained from Si IV. For this analysis, we use the results obtained from the combined dataset. Note that we only consider the Mg II k2, Mg II k3, C II 1334 Å and Si IV lines in this analysis. We emphasize that the results from Mg II h follow the results from k feature and hence are not shown for brevity.

We split the velocities observed in Mg II, C II and Si IV into sets of pixels containing upflows and downflows. Then, we consider scatter plots between flows in intersection of these sets, e.g., relation between pixels showing upflows in Mg II k3 & upflows in Si IV, upflows in Mg II k3 & downflows in Si IV and so on. These scatter plots are obtained for Si IV velocities, in Mg II and C II velocity bins to improve statistics. Note that the bins here are selected in deciles, i.e, every 10% of the data for each Mg II or C II feature is considered to be in one bin.

In Fig. 13, we plot the correlations between downflows (top row) and upflows (bottom row) observed in Si IV with those observed in Mg II k2 (panels a & d), Mg II k3 (panels b & e) and C II (panels c and f). Panels a, b and c demonstrate that the downflows observed in Si IV are strongly correlated with those observed in Mg II k2, k3 and C II. For a given value of downflow in Si IV, the downflows are stronger in k3 and C II than those in k2. Note though, that the downflows in k3 and C II are very similar. Moreover, Si IV displays excess downflows in QS vis-à-vis CH for similar C II and Mg II downflows. These differences in QS-CH Si IV downflows are also observed to increase with increasing Mg II and C II downflows. Similarly, Panels d, e and f suggest that the upflows in Si IV have slightly better correlation with upflows in Mg II k3 and C II than those in Mg II k2. Moreover, like downflows, we find that for the similar upflows in Mg II and C II, the CH exhibit larger upflows vis-à-vis QS in Si IV. We further note that there is a slight hint of increase in the difference of upflows observed in CH and QS in Mg II and C II lines.

In Fig. 14, we study the correlations between the upflows in Si IV with downflows in Mg II and C II and vice-versa as shown in the top and bottom rows, respectively. We find that the upflows in Si IV have a monotonic relation with the downflows in Mg II and C II (top row). In addition, we note that for similar downflows observed in Mg II and C II, Si IV shows stronger upflows in CH and than in QS. On the other hand, the Si IV downflows do not show any particular correlation with Mg II and C II upflows. Furthermore, the Si IV downflows in CH and QS remain consistent for similar upflows in Mg II and C II. We do not find any relation between the downflows in Si IV with upflows in Mg II and C II.

The problems of coronal heating in QS and CHs, and the formation and acceleration of solar wind are intimately tied to the structure and dynamics of the magnetic field in the respective regions. Therefore, comparative studies between CHs and QS become important in understanding the plasma dynamics, and the underlying processes. The Mg II, C II and Si IV lines observed by IRIS probe different layers in the chromosphere & TR, have provided us with a unique opportunity to understand the dynamics of these regions as a dynamically coupled system. In this paper, we characterize, in detail, the dynamics of the Mg II line by combining the information related to the plasma dynamics with that of the magnetic field. In addition, we also probe the correlations between the Doppler shift obtained for Mg II lines, and those obtained for the Si IV line from Paper I and the C II line from Paper II, to investigate if a common origin may be ascribed to the observed dynamics in the different lines. Below we summarize and discuss our results followed by an interpretation in §8.

While the occurrence of impulsive events at the interface between chromosphere and corona may explain the observed flow variations with $|$B$|$ in different lines and their interrelations, the question remains as to what leads to the observed differences between the intensities as well as flows in CHs and QS. To explain the differences in the intensities, in §7.1 we invoked the loop statistics in CHs and QS derived by Wiegelmann & Solanki (2004). The predominant velocity differences between CH and QS are: i) reduced Si IV downflows in CHs over QS for similar downflows in the chromospheric lines, and ii) excess Si IV upflows in CHs over QS for similar upflows and downflows in the chromospheric lines. These results indicate excess acceleration of upflows in CHs and excess deceleration of downflows in QS. Furthermore, while the QS shows enhanced intensity over CHs for regions with similar $|$B$|$, the CHs show larger flow speeds (except Si IV downflows) over QS. Such an observation thus hints towards a unified scenario of heating of the corona in QS and CH as well as the emergence of the solar wind. Therefore, we then ask if it is possible to combine the loop statistics and the occurrence of flows due to impulsive events illustrated in Fig. 15, to explain the observed differences in intensities as well as the Doppler shifts in CHs and QS, similar to that is discussed in Tripathi et al. (2021) for Si IV.

A graphic depicting the scenario we propose is shown in Fig. 16. The left panel depicts the predominant topology in CHs while the right panel is for QS regions, based on the loop statistics of Wiegelmann & Solanki (2004), according to which both CHs and QS have equal number of short closed loops but QS has predominantly large closed loops and CHs have open field lines. In CH regions, the interchange reconnection (e.g., Fisk, 2005; Janardhan et al., 2008), leading to impulsive events, may occur between closed and open field lines, while in QS, the impulsive events will be due to reconnection among closed-closed loops, similar to those observed in the core of active regions during transient formation of loops (see Tripathi, 2021). The excess open, and expanding flux tubes in CHs may cause preferential acceleration of upflows in CHs over QS. In principle, the scenario proposed here is similar to those employed by Tian et al. (2008b, a); He et al. (2007) to explain the Doppler shifts observed in QS-CH in coronal and TR lines, and similar to the Fig. 5 of He et al. (2010). Note that the concept of interchange reconnection has been invoked to explain active region outflows by Del Zanna et al. (2011); Barczynski et al. (2021), solar wind disappearance events by Janardhan et al. (2008), active region jets, X-ray & cool jets in polar coronal holes (Moore et al., 2011, 2013, 2015) and type III radio burst (e.g., Mulay et al., 2016).

Under the scenario presented in Fig. 16, the impulsive events may occur across a range of heights via magnetic reconnection among open field and closed loops of various heights in the CHs. Thus, if the energy dumping events were to occur below the formation height of Mg II, the upflowing plasma may be accelerated preferentially in CHs, and reach Si IV heights, where a strong correlation is obtained. Similarly, if the event were to occur at much greater heights, or if plasma launched from the lower heights (e.g. type II spicules-like events) are returning to the low solar atmosphere, the downflowing plasma may be falling from Si IV formation height, that will show deceleration due to increasing density towards lower atmosphere mapped by Mg II. Assuming a mass flux conserving flow, we have $\rho\mathrm{V}=\mathrm{const}$, where $\rho$ is the mass density and $\mathrm{v}$ is the velocity. For a given downflow in Si IV, the downflows in chromospheric lines are smaller in QS over CH (Fig. 13.a-c). Hence, the density increase from Si IV formation heights to Mg II formation heights is larger in QS over CH, resulting in a larger velocity reduction in QS. However, note that since the mass flux is typically different for CHs and QS, a quantitative comparison is beyond the scope of this work.

For bidirectional flows due to reconnection event between Mg II and Si IV formation heights, the counterpart upflows will be preferentially accelerated into Si IV formation height in CHs over QS due to excess open expanding flux tubes in CHs over QS. Since the QS has predominantly closed loops, the closed loop reconnection only serves to fill the loop with plasma, raising its intensity. Thus, impulsive events occurring across a range of heights combined with loop statistics in CHs and QS, elegantly ties in together all our observations, and provides a unified scenario for QS heating and solar wind emergence.

Finally, the scenario we present in Fig. 16 is also appealing to explain the switchbacks observed in the near-Sun solar wind (Balogh et al., 1999; Bale et al., 2019) using Parker Solar Probe. One of the competing scenarios for the formation of these switchbacks is through interchange reconnection events occurring in the TR and lower corona (Fisk & Kasper, 2020; Mozer et al., 2021; Zank et al., 2020; Tripathi et al., 2021; Fargette et al., 2021; Bale et al., 2021; Sterling et al., 2020; Sterling & Moore, 2020, see also Liang et al. (2021) for an assessment of viability of switchbacks from the linear theory of Zank et al. (2020)). The kinked-field lines as a result of reconnection between the close loop and open field in the coronal holes, as shown by black arrow in Fig. 16, may be transported outwards into the solar wind, which are then observed as rotations in the magnetic field. In such a scenario, the flows reported in this paper serve as constraints and modeling inputs for solar wind switchback simulations.

A straightforward association between the scenario we present, and polar coronal hole jets is clearly seen. Interchange reconnection seems to play the predominant role in generation of mass flux and plasma heating in these events. However, note that the jets observed by Moore et al. (2015) have velocity almost two orders of magnitude more than the velocities we report, and show morphological differences arising due to twist and shear in the ambient magnetic field (see also Moore et al., 2013). Since we are averaging over multiple pixels for boosting signal, checking for such morphological signatures is beyond the scope of this work. However, newly-emerged bipoles may interact with the ambient vertical field similar to the scenario proposed by Moore et al. (2011), giving rise to Spicule-like events. The interaction height, and amount of magnetic flux converted into thermal energy would then determine the lines which show correlated flow.

The observational results and the scenario presented in this paper explain the solar wind formation including switchbacks and the dynamics observed in the QS. We, however, stress that disentangling the absolute effects of wave propagation v/s impulsive upflows is needed. Furthermore, disentangling the effect of return of spicule-like events v/s downflows due to impulsive events occurring higher up in the TR is also difficult. Disentangling these different effects requires further high resolution spectroscopic observations simultaneously taken at different heights, combined with numerical simulations incorporating radiative transfer & evolution of solar corona into the solar wind. Such observations may be provided with the EUV High-Throughput Spectroscopic Telescope (EUVST) on the upcoming Solar-C mission (Shimizu et al., 2020).

The authors thank Téo Bloch (University of Reading) for suggesting binning in deciles and Vilangot Nhalil Nived (Armagh Observatory and Planetarium, U.K ) for providing coalignment and Si IV fitting codes in IDL for cross check. The authors also thank Mats Carlsson (University of Oslo), Mark Cheung (Lockheed Martin Solar and Astrophysics laboratory), Giulio Del Zanna (University of Cambridge), Jim Klimchuk (NASA Goddard) and Aveek Sarkar (PRL, Ahmedabad, India) for discussions and suggestions. The authors sincerely thank Nishant Singh (IUCAA, Pune, India) for comments on the manuscript. U.V. thanks his sister Vaishnavi Upendran for suggestions regarding aesthetics of Fig. 15 and Fig. 16. We acknowledge the use of data from IRIS, AIA and HMI. IRIS is a NASA small explorer mission developed and operated by LMSAL with mission operations executed at NASA Ames Research Center and major contributions to downlink communications funded by ESA and the Norwegian Space Centre. AIA and HMI are instruments on board SDO, a mission for NASA’s Living With a Star program.