The solar magnetic field extends from the solar interior, across the photosphere, through the chromosphere and transition region, and into the corona, permeating the entire volume and out into the heliosphere. Plasma motions in the solar convective envelope produce a highly structured photospheric magnetic-field distribution, which can be considered the boundary condition for the magnetic field above. As the photospheric boundary field evolves, the field above must reconfigure, and, because the atmosphere extends over many scale heights, the energetics and dynamics of the response are diverse and multi-scaled. A central goal of the DKIST science mission is a quantitative understanding of the complex interplay between flows, radiation, heat conduction, wave propagation and dissipation, and reconnection in the solar atmosphere. This understanding is critical to the resolution of long-standing problems in solar and stellar astrophysics such as chromospheric and coronal heating, the origin and acceleration of the solar wind, and the propagation of magnetic disturbances into the heliosphere.
Understanding the complex, connected solar atmosphere, with its widely varying ionized and partially ionized plasma regimes, requires diverse, flexible, and multi-spectral-line polarimetric instrumentation capable of simultaneously probing many heights at high temporal cadence and over small spatial length scales. DKIST is at the forefront of these observational challenges. Its new capabilities in high-precision spectropolarimetry will allow the inference of magnetic-field properties with spatial and temporal resolutions never before achieved. Those inferences will depend on advanced and challenging inversions. For example, while the distribution of solar photospheric magnetic fields has been routinely measured since the invention of the Zeeman-effect-based magnetograph, measurement of the chromospheric field and its connections to the upper solar atmosphere requires interpretation of spectral lines formed under conditions of non-local thermodynamic equilibrium (NLTE) in which both Zeeman and atomic-level polarization processes are important. DKIST aims to facilitate these with chromospheric data of unprecedented precision.
DKIST is also at the forefront of coronal spectropolarimetry. Extrapolation of the magnetic field into the solar corona is inherently limited by assumptions about the distribution of currents in the magnetized volume. To advance our understanding of, for example, the available free magnetic energy driving solar flares and eruptions, remote sensing of the magnetic field in the corona itself is necessary. DKIST, as the world’s largest coronagraphic polarimeter, will allow the measurement of the full Stokes spectra of the forbidden magnetic-dipole emission lines of highly ionized metals in the corona. These can be used to determine the topology and evolution of the coronal field and, along with DKIST chromospheric diagnostics, its connectivity to the lower atmosphere.
Several critical science topics in this research area are discussed in detail below, including i) the mass and energy cycle in the low solar atmosphere, ii) the origin and acceleration of the solar wind, iii) magnetic reconnection throughout the solar atmosphere, iv) waves in the solar atmosphere, v) impact of flux emergence on the non-eruptive solar atmosphere, vi) multilayer magnetometry, and vii) large-scale magnetic topology, helicity, and structures.
Mass and Energy Cycle in Low Solar Atmosphere
How important are spicules and other jet-like phenomena to the chromosphere–corona mass cycle? What is the role of spicule heating in the coronal energy balance? Can we characterize and model the coronal-rain phenomenon well enough to understand its role as a return flow?
All coronal plasma has its origins in the lower solar atmosphere, with the coronal mass budget determined by a balance between upward flows (evaporative or eruptive), downward mass transport (e.g. coronal rain), and solar-wind losses. One of the key outstanding challenges to furthering our understanding of the solar atmosphere is determining which of many process that appear to be involved dominates the transfer of mass between the cool chromosphere and hot corona in regions of differing magnetic topology. Previous observations have provided a partial view of the mass cycle, but the small spatial and temporal scales involved have significantly hindered advancement. DKIST will provide a more complete view, especially when combined with coordinated observations using space-based observatories, such as IRIS, Hinode, and Solar Orbiter, or the ground-based radio telescope array ALMA (Atacama Large Millimeter/sub-millimeter Array), which allows complementary temperature and soon polarization diagnostics of the solar chromosphere (see Yokoyama et al., 2018).
Jets play an important role in the chromosphere–corona mass cycle. They are observed to have a wide range of spatial scales, from spicules with widths of a few hundred kilometers, to larger chromospheric jets with widths of \(\approx 1000~\mbox{km}\), to coronal jets with widths of a few thousand kilometers. Recent observations show that some jets are driven by the eruption of a small-scale flux ropes (Sterling et al., 2015) triggered by flux cancelation at the magnetic neutral line (Panesar et al., 2016, 2018; Panesar, Sterling, and Moore, 2017), but the full range of magnetic-field reconfiguration scenarios that lead to the jets observed is not known. DKIST’s spectropolarimetric capabilities will significantly contribute to our understanding of jet initiation, allowing us to better quantify how effectively each jet type injects plasma into the solar atmosphere.
Spicules appear to be the most ubiquitous jet-like feature in the solar chromosphere (Figure 9). They are highly dynamic, vary on time scales of 10 – 30 seconds, and are finely structured, with widths \(<300~\mbox{km}\) and substructure at scales below that. Spicules transport plasma upwards into the solar atmosphere at speeds of \(10\,\text{--}\,200~\text{km}\,\text{s}^{-1}\) and may, either alone (De Pontieu et al., 2011) or in combination with other slower jets (e.g. Morton et al., 2012), play a significant role in the mass and energy balance of the corona and solar wind. Yet, despite having been observed over a wide range of wavelengths from EUV to visible, we do not understand the mechanisms responsible for spicule formation, how they are heated, often to transition-region temperatures or higher (Pereira et al., 2014; Henriques et al., 2016; Samanta et al., 2019), the role that hydromagnetic waves play in their dynamic evolution, or their full impact on the mass and energy balance of the outer solar atmosphere (Pereira, 2019). That coronal signatures of spicules have been observed at the smallest scales above both active regions (De Pontieu et al., 2011) and the quiet Sun (Henriques et al., 2016) suggests that they may play a global role as heat and mass conduits between the chromosphere and corona, but our understanding of their importance is limited by current instrumentation capabilities. Spicule-induced mass transport to coronal heights is estimated by some authors to be two orders of magnitude larger than the solar-wind mass flux (Beckers, 1968; Sterling, 2000), but others suggest that the role of spicules in the mass cycle is rather limited compared to that of more uniform chromospheric evaporation due to small-scale flaring processes (e.g. Klimchuk, 2012). In fact, there is still some debate about the fundamental nature of spicules, whether only one or multiple spicule types exist (De Pontieu et al., 2007a; Zhang et al., 2012, but see Pereira, De Pontieu, and Carlsson, 2012), or whether they are indeed jets of accelerated plasma, with an alternative suggestion that the apparent upward motions are due instead to line-of-sight integration through evolving warped two-dimensional magnetic sheet-like structures (Judge, Tritschler, and Low, 2011; Judge, Reardon, and Cauzzi, 2012; Lipartito et al., 2014).
DKIST observations, employing a wide range of chromospheric spectral lines, will provide revolutionary new views of the fine-scale structure, magnetic and electric fields (Anan, Casini, and Ichimoto, 2014) within, and thermodynamic evolution of, spicules and other chromospheric jets. DKIST will be able to capture the dynamic evolution of jets at high cadence (\(\lesssim 3~\text{seconds}\)) and make simultaneous measurements (with \(\lesssim 15\)-second cadence) of the chromospheric and photospheric magnetic fields and flows that underlie jet initiation and acceleration. Together these will help clarify issues such as the relationship between jet initiation and the apparently ubiquitous Alfvén pulses that are excited by swirling motions in the solar photosphere (Goodman, 2012; Liu et al., 2019a,b) and the roles and relative importances of processes such as reconnection (Section 5.3 below), magnetic tension amplification by ion–neutral coupling (Martínez-Sykora et al., 2017), micro-filament eruption (Sterling and Moore, 2016), and other small-scale processes.
As already indicated, some have suggested that, while jet-like features are perhaps the most prominent, they may not be the most important component of the chromosphere–corona mass cycle, that more gentle processes play an important, even dominant, role. One sensitive signature of mass transport, momentum flux, and wave heating is the degree of elemental fractionation, i.e. the degree to which elements of low (lower than about 10 eV) first ionization potential (FIP) are enriched or depleted compared to other elements (the FIP or inverse-FIP effects: Meyer, 1985a,b; Feldman, 1992; Feldman and Laming, 2000; Laming, 2015). Because FIP fractionation is sensitive to the thermodynamic and electromagnetic plasma environment, it depends on the atmospheric height of the region being considered, the magnetic-field geometry (open or closed) and connectivity, the heating mechanisms at play, the bulk plasma flow speed, and the dominant magnetohydrodynamic (MHD) wave modes present (Laming, 2015). In fact, spicules themselves may be responsible for the absence of the FIP effect at temperatures below \(10^{6}~\text{K}\) on the Sun (Laming, 2015, and the references therein). These sensitivities, when understood well, allow FIP fractionation measurements to be used to constrain the solar-wind source region and thus the origin of plasma sampled during in-situ heliospheric measurements (e.g. Geiss, Gloeckler, and von Steiger, 1995; Parenti et al., 2000; Brooks and Warren, 2011; Brooks, Ugarte-Urra, and Warren, 2015; Baker et al., 2015). Coordinated observations (Section 6.5), combining chromospheric measurements by DKIST with in-situ measurements of the heliospheric plasma properties by Solar Orbiter and Parker Solar Probe, will allow unprecedented identification and characterization of the sources of the fast and slow solar wind.
Another key aspect of the chromosphere–corona mass cycle is the return flow from the corona to the lower atmosphere. One component of this is prominently visible at the solar limb as coronal rain: a finely structured and multi-thermal flow that appears to be driven by cooling instabilities (e.g. Antolin and Rouppe van der Voort, 2012; Antolin et al., 2015b; Mason, Antiochos, and Viall, 2019, and the references therein). The descending material is visible in the off-limb active corona at transition-region and chromospheric temperatures, and can be observed in what are typically chromospheric optical spectral lines (Kawaguchi, 1970; Schrijver, 2001; Antolin, Shibata, and Vissers, 2010). Current observations are unable to fully resolve coronal rain: the peak of the coronal-rain element-width distribution remains unresolved (Scullion et al., 2014). Thus we do not yet have a full accounting of the total rain drainage rate. Measurements of loop oscillations can be used to determine the thermally unstable mass fraction in a coronal-loop system and thus provide further constraint on that rate. Since coronal-rain condensations are coupled to the magnetic-field lines they track the oscillatory motion of the loops (Kohutova and Verwichte, 2016) and the evolution of the oscillations can be used to deduce the fraction of the loop plasma mass that becomes thermally unstable and drains with time. This seismic estimation provides more than a consistency check on the drainage rate because it also allows measurement of the unresolved coronal-rain mass fraction (Froment et al., 2018), which in turn constrains the fundamental spatial scales of the rain and the fine-scale structure of coronal loops. DKIST observations will push these inferences even further.
One-dimensional hydrodynamic loop models (Mendoza-Briceño, Erdélyi, and Sigalotti, 2002; Müller, Hansteen, and Peter, 2003; Mendoza-Briceño, Sigalotti, and Erdélyi, 2005) produce catastrophic cooling events that generate intermittent and repeating rain-like downflows even when steady, exponentially decaying with height, foot-point concentrated heating is employed. Similarly, 2.5-dimensional and recent three-dimensional simulations form fine-scaled localized rain elements even when a spatially smooth but localized foot-point concentrated heating function is applied (Fang, Xia, and Keppens, 2013; Moschou et al., 2015). These numerical studies, along with observational evidence for foot-point concentrated heating in active regions (Aschwanden, 2001), suggest that studies of coronal rain are important, not just in the context of the mass cycle, but also in constraining heating mechanisms. The statistical properties of the rain depend on the spatial distribution of the heating because the heating location influences the thermal stability of the plasma in a coronal-loop system (Antolin, Shibata, and Vissers, 2010).
Observational contributions to our understanding of coronal rain depend on leveraging the high-resolution and multi-wavelength capabilities of DKIST to unravel its evolving, complex, multi-thermal behavior. With a mix of hot (ionized) and cold (neutral) gas, coronal rain is an ideal context within which to study ion–neutral interaction effects, such as ambipolar diffusion, that are expected to play key roles more broadly in the dynamics of the partially ionized chromosphere. But the rain’s thermal state is complex and difficult to characterize. Hydrogen is likely out of ionization equilibrium and other elemental ionization ratios are poorly constrained (Antolin, Shibata, and Vissers, 2010; Antolin et al., 2015b). DKIST will also be able to provide spectropolarimetric measurements from which the vector magnetic field within rain-producing loop systems can be inferred. Such inferences are fundamental in characterizing the underlying flow instabilities (e.g. Martínez-Gómez et al., 2020), but they will require careful theoretical underpinning. Recent studies have been able to exploit the Zeeman effect in the Ca ii 854.2 nm chromospheric line to assess the presence of strong fields (100 – 300 G) in bright post-flare loops systems (Kuridze et al., 2019), but this approach may be difficult in more quiescent cases, which are characterized by much weaker fields and lower photon counts. Significant progress may be possible if Hanle diagnostics in lines of neutral Helium such as the He i 1083 and 587.6 nm multiplets are employed (Schad, 2018), and polarimetric observations of these two He i multiplets with DKIST may provide the first view of sub-arcsecond scale weak magnetic fields (with sensitivity of a few Gauss) in quiescent loop systems.
Coronal Heating, Solar Wind Origin and Acceleration
Waves or nanoflares? What are the relative importances of different proposed coronal-heating mechanisms? What are the solar-wind momentum sources? What role does the chromosphere play in coronal heating?
The steep transition from the cool chromosphere to the million-degree and higher corona is the direct result of the cooling catastrophe that results when hydrogen in the solar atmosphere becomes fully ionized so that radiative recombination can no longer cool the optically thin plasma (Woods, Holzer, and MacGregor, 1990a,b). The plasma temperature climbs high enough so that electron conduction back down to the chromosphere, and resulting radiative losses from there, are sufficient to balance the heating above. Solving the coronal-heating problem requires identifying the heat source, which in the statistically steady state balances thermal conduction to the chromosphere and any other smaller direct energy losses from the corona by radiation or advection. Significant progress has been made in the last decades, and it is now apparent that no single heating mechanism is likely universally dominant (e.g. Kuperus, Ionson, and Spicer, 1981; Zirker, 1993; Klimchuk, 2006; Erdélyi and Ballai, 2007; Parnell and De Moortel, 2012; De Moortel and Browning, 2015). The importance of mechanisms such as reconnection, MHD or plasma-wave dissipation, or turbulent dissipation likely vary depending on coronal conditions, particularly the magnetic-field configuration. Moreover, the mechanisms are highly intertwined and interdependent; dissipation of current sheets can produce waves and waves in magnetically structured media can induce current sheets (e.g. Velli et al., 2015), and observations indicate that even on very large scales eruptive flares can trigger oscillations in coronal loops and filament oscillations can induce eruption (e.g. Jess et al., 2015; Russell, Simões, and Fletcher, 2015, and the references therein) and associated flaring. Turbulence is similarly likely ubiquitous, and theoretical work suggests an important role for Alfvén-wave induced turbulence (e.g. van Ballegooijen et al., 2011; Asgari-Targhi et al., 2013; van der Holst et al., 2014).
Thus, while a number of general properties of coronal heating (e.g. De Moortel and Browning, 2015) have been observationally established (that coronal heating is unsteady and impulsive, that coronal magnetic fields store energy that can be dissipated via reconnection, that the corona supports a rich wave field, and that the corona can only be understood in conjunction with its coupling to the chromosphere) details are less certain (e.g. Klimchuk, 2015; Schmelz and Winebarger, 2015). Questions remaining include: What is the relative importance of different heating mechanisms? Do current sheets in the corona play an important role in heating, and if so what plasma processes are involved in their dissipation? What is the fundamental scale of the substructure in multi-thermal coronal loops, and what processes determine this? During small-scale reconnection events, how much energy is dissipated directly, how much is radiated as waves, and how much goes into acceleration of non-thermal electrons (Testa et al., 2014)? How does magnetic reconnection heat the plasma (Longcope and Tarr, 2015)? How and where are the ubiquitous MHD waves in the corona dissipated (e.g. Poedts, 2002; Gupta, 2017)? How far out in the solar wind does heating extend (Martinović, Klein, and Bourouaine, 2019, and the references therein)? What are the characteristic scales and magnitudes of the heating events? What triggers them?
DKIST will be fundamental in answering these questions because it will enable careful, repeated, and frequent measurements of the plasma properties of the inner corona, including the local magnetic-field strength and direction, at multiple heights. These will facilitate the quantitative evaluation of suggested heating processes in varying magnetic environments. Measurement of the line–continuum and line–line intensity ratios of ionic species, such as Fe ix to xv, that are also found in in-situ measurements of the fast and slow solar wind, will allow determination of the electron-density, temperature, and charge-state evolution of the solar-wind plasma, informing our understanding of the acceleration and heating processes (Figure 10, Landi et al., 2012; Landi, Habbal, and Tomczyk, 2016; Boe et al., 2018). Comparison between observations and theoretical studies of fast and slow magneto-acoustic wave-mode conversion, shock formation, and dissipation (e.g. Zhugzhda, Bromm, and Ulmschneider, 1995; Carlsson and Stein, 1997) will help constrain the role of these processes as a function of height in the chromosphere. For example, in models, high-frequency (\(>10~\text{mHz}\)) propagating acoustic waves develop into radiatively damped weak shocks within the first few hundred kilometers above the photosphere (Carlsson and Stein, 2002) while lower-frequency waves (\(\approx4\,\text{--}\,10~\text{mHz}\)) develop into strong shocks in the chromosphere (above 1 Mm) where radiative damping is less effective (Priest, 2000). Multi-height DKIST observations will allow careful measurement of this frequency dependence in regions of differing magnetic-field configurations and thus an assessment of wave-energy deposition with height. Moreover, previous high-resolution observational work has reported evidence for shock-induced turbulence (Reardon et al., 2008), and such turbulence may provide a mechanism for the dispersal of the wave energy beyond the local shock region itself. Turbulence may also play a direct role in wave-mode conversion, coupling the compressive motions to Alfvénic fluctuations which can continue to propagate outward, transmitting energy to higher layers of the solar atmosphere. Alternatively, counter-propagating Alfvén waves may nonlinearly interact to produce MHD turbulence, dissipating the waves, and heating the plasma (van Ballegooijen et al., 2011). DKIST observations can differentiate between these processes in the solar atmosphere and determine their occurrence frequency.
Further, even though direct observation of individual nanoflare heating events lies beyond the capabilities of DKIST, DKIST observations will help distinguish between specific reconnection heating mechanisms. Heating models based on flux cancelation have been previously motivated by high-resolution Imaging Magnetograph eXperiment (IMAX) data from Sunrise (Priest, Chitta, and Syntelis, 2018) and heating events in simulations of three-dimensional kink-unstable flux ropes may be diagnosable using DKIST coronal lines (Snow et al., 2018). With DKIST, high-resolution chromospheric observations of intensity fluctuations and motions at the foot points of hot coronal loops should reveal key telltale signatures of nanoflares in the overlying corona (Testa et al., 2014). When coupled with radiative-hydrodynamic modeling (e.g. Kerr et al., 2016; Polito et al., 2018), such observations may be able to constrain the properties of the non-thermal particles or waves generated at nanoflare sites. Finally, although the details are uncertain, it has been suggested that the chromosphere may play an important role in coronal heating (e.g. Withbroe and Noyes, 1977; Sturrock, 1999; De Pontieu et al., 2017) since chromospheric plasma can be heated to transition-region temperatures and higher and be carried via jets to coronal heights. If so, the jet studies outlined in Section 5.1 are very relevant to this research topic as well.
The hot, outer solar corona escapes the Sun as the solar wind. The solar wind carries plasma out into the heliosphere, where it interactions with the planetary space environments and importantly influences CME arrival times and geo-effectiveness, yet there is no consensus understanding of where the solar wind originates or how it is accelerated. The fast solar wind and slow solar wind have different physical properties (e.g. Feldman, Landi, and Schwadron, 2005; Ebert et al., 2009), likely come from different source regions, and are possibly subject to different acceleration mechanisms. The fast wind originates in coronal holes, but how the detailed properties of the field and plasma within a coronal hole lead to the observed properties of the wind is not clear. For example, additional sources of momentum, beyond Parker’s original gas-pressure gradient mechanism (Parker, 1958, 1963), are required for the plasma to reach the observed fast-wind speeds. Possible momentum sources included large-amplitude MHD waves (Alazraki and Couturier, 1971; Jacques, 1977; De Pontieu et al., 2007b; Thurgood, Morton, and McLaughlin, 2014), Type-II spicules (e.g. De Pontieu et al., 2009; Moore et al., 2011), resonant interactions with ion-cyclotron waves (Hollweg and Isenberg, 2002), and others (e.g. Cranmer and Winebarger, 2019, and the references therein), but these are to date poorly constrained by observations. Similarly, the slow wind has variously been proposed to originate from the edges of coronal holes, closed–open field boundaries within and bordering active regions, streamers (particularly streamer tops), or small coronal holes, but while these locations are all associated with magnetic reconnection between closed magnetic flux systems and open ones that connect to the wind, there is no consensus on the dominant underlying field configuration or acceleration mechanism (e.g. Feldman, Landi, and Schwadron, 2005; Cranmer, 2009). In fact, some of the observed differences between the fast and slow solar winds may have more to do with the expansion properties of the background magnetic fields along which they are streaming than with differences between their source regions (Wang and Sheeley, 2003).
The solar wind can be studied using either remote-sensing or in-situ techniques. In-situ measurements typically provide direct information on plasma properties only after the plasma has undergone much of its evolution, although Parker Solar Probe (PSP) is revolutionizing these measurements, aiming to sample, during its closest perihelia, regions of solar-wind heating and acceleration directly (e.g. Venzmer and Bothmer, 2018). Remote-sensing observations, on the other hand, allow frequent multi-wavelength measurements of the solar-wind source regions, but can be difficult to interpret. Combining these types of measurements (Landi et al., 2012) is already yielding exciting results in the current early-PSP era (Rouillard et al., 2020). DKIST will make significant contributions to these efforts (Section 6.5). Regular off-disk coronal measurements of the magnetic field and the plasma properties will characterize the solar wind as it emerges from its source up to a height of 0.5 solar radii above the limb. Since the visible and infrared spectral lines observed by DKIST largely result from photoexcitation, their intensities decrease more slowly with height (closer to being in proportion to the electron density rather than its square) than do those of EUV spectral lines, which result largely from collisional excitation (e.g. Landi, Habbal, and Tomczyk, 2016; Del Zanna and DeLuca, 2018). This makes plasma and magnetic-field diagnostics to the outer edge of the DKIST coronal field of view possible. Importantly, lower down, in the chromosphere and low transition region, where current space-based EUV observational efforts are focused, DKIST will enable inference of the vector magnetic field via their spectropolarimetric signatures in lines such as He i 1083.0 nm.
DKIST’s capabilities in combination with those of IRIS and Hinode will facilitate studies of solar-wind acceleration physics from the chromosphere through transition region and into the corona. Through coordinated efforts, these observatories will be able to address critical issues such as how the physical properties of the solar wind change with height, how that profile depends on position from the center to the edge of coronal holes, what defines coronal-hole–streamer boundaries, and how does reconnection at the boundary between closed and adjacent open field in active regions yield the observed wind properties. Plasma diagnostics will allow estimates of the mass and energy flow along magnetic-field lines, and coronal line-width studies will help in understanding wave propagation and damping (Hahn, Landi, and Savin, 2012). Combining magnetic-field, electron-density, and temperature measurements will make possible charge-state evolution modeling of the accelerating solar wind and aid in the development of empirical models of the solar-wind speed between 1 and \(1.5~{\mathrm{R}}_{\odot }\) (Landi et al., 2012). Extrapolation to the freeze-in height, with field extrapolations no longer strictly dependent on the photospheric field, will provide further links to in-situ instrumentation on PSP, SO, and Advanced Composition Explorer (ACE) and constrain the solar-wind source locations (Section 6.5).
Magnetic Reconnection in the Solar Atmosphere
What is the three-dimensional geometry of the magnetic field near reconnection sites? What roles do ion–neutral collisions, current-sheet instabilities, and plasmoid ejection play? How efficiently does the reconnection heat and accelerate the solar plasma?
Magnetic reconnection is a fundamental process that transforms magnetic energy into kinetic and thermal energies in astrophysical plasmas. The magnetic energy is stored, maintained, and amplified over extended periods of time (minutes to hours in magnetic network and hours to weeks in sunspots and active regions) before being released suddenly during reconnection events, with flares occurring on small spatial scales (kilometers) over very short times (minutes). Reconnection events can accelerate plasma in jet-like structures and induce local heating in the chromosphere and above.
In addition to the ubiquitous spicules (Section 5.1), jets are observed in sunspot penumbre (Katsukawa et al., 2007; Tiwari et al., 2018), at the edges of sunspots (Morton, 2012), in light bridges (Toriumi, Katsukawa, and Cheung, 2015; Tian et al., 2018), and in the plage regions surrounding sunspots (Nishizuka et al., 2011). These jets often display morphologies reminiscent of reconnection sites (e.g. Innes et al., 1997; Shibata et al., 2007; Singh et al., 2012), but there are few direct measurements of the local magnetic field and its reconfiguration by reconnection. Localized heating is also observed, for example as Ellerman bombs and UV bursts: point-like brightenings in the wing of chromospheric lines such as H\(\alpha \) and Ca ii (e.g. Vissers et al., 2015; Reid et al., 2016; Rouppe van der Voort et al., 2017; Toriumi, Katsukawa, and Cheung, 2017; Young et al., 2018). These brightenings are often associated with bipolar moving magnetic features near sunspots or colliding bipolar structures in regions of emerging flux, similar to magnetic cancelation events in the quiet Sun (Rouppe van der Voort, Rutten, and Vissers, 2016; Nelson et al., 2017a). Reconnection is again implicated (Figure 11), but again without direct measurements of field reconfiguration. Further, brightening and excess line broadening is observed in transition region and coronal loops, consistent with impulsive reconnection heating, bidirectional flows, and ion-cyclotron turbulence at sites of magnetic braiding of the underlying multi-stranded field (Cirtain et al., 2013; De Pontieu et al., 2014b; Hansteen et al., 2014; Bahauddin, Bradshaw, and Winebarger, 2020).
An important goal of DKIST is to measure the magnetic-field changes associated with reconnection phenomena at high resolution and simultaneously over multiple heights in the solar atmosphere from the photosphere into the chromosphere. Diagnosis of the magnetic field at reconnection sites before and after reconnection events will facilitate reconstruction of the magnetic and thermodynamic history of the plasma in the reconnection volume. It will allow detailed assessment of the total field annihilation, reconnection rate, total magnetic-energy release, and local heating induced. Together with recent in-situ measurements of the underlying microphysical processes in the Earth’s magnetosphere (e.g. Burch et al., 2016; Torbert et al., 2018; Chen et al., 2020; Hesse and Cassak, 2020) and laboratory experiments (e.g. Dong et al., 2012; Gekelman et al., 2016; Olson et al., 2016; Howes, 2018; Takahata, Yanai, and Inomoto, 2019; Seo et al., 2020), such studies will advance our fundamental understanding of the energetics of solar and astrophysical reconnection, determining how the energy is partitioned between bulk flows and random motions, between plasma acceleration and heating (Ji and Daughton, 2011; Coates, 2016; Yamada, Yoo, and Myers, 2016).
A key aspect of DKIST’s contributions to reconnection studies is the particular plasma environment that will be sampled. The lower solar atmosphere (photosphere and chromosphere) is relatively dense and weakly ionized, with an ionization fraction on the order of \(10^{-4}\) at the height of the temperature minimum (e.g. Khomenko, 2017). Weakly ionized plasmas are found over a wide range of astrophysical settings, including the atmospheres of other cool stars, the warm neutral interstellar medium (ionization fraction of \(10^{-2}\): Jenkins, 2013), dense cores of molecular clouds (ionization fraction of \(10^{-7}\): Caselli et al., 1998), and protostellar and protoplanetary disks (ionization fraction \(10^{-10}\) or less: e.g. Armitage, 2019). Importantly, low ionization fractions (in the range of those found in the solar chromosphere) increase reconnection rates. Ion–neutral interactions can lead to increased resistivity and an increase in the effective ion mass, with consequent reduction in the Alfvén speed, steepening of current sheets, heating of the inflow and exhaust regions of reconnection sites, and enhancement of the plasmoid instability (e.g. Zweibel, 1989; Chiueh, 1998; Zweibel et al., 2011; Leake et al., 2012; Mei et al., 2012; Murphy et al., 2012; Zweibel, 2015; Ni et al., 2020). These in turn have important implications for solar phenomena, likely being responsible for increased damping of MHD waves (De Pontieu, Martens, and Hudson, 2001), increased current dissipation and heating of the solar chromosphere (Khomenko and Collados, 2012; Martínez-Sykora, De Pontieu, and Hansteen, 2012), increased flux emergence rates into the corona, reduced Alfvén wave flux from photospheric foot-point motion, and changes in the structure of MHD shocks, prominences, and quiet-Sun magnetic features (see Anan, Ichimoto, and Hillier, 2017, and the references therein). Observational verification of the reconnection implications of partial ionization is more readily achieved in these solar contexts than in more distant astrophysical settings.
Waves in the Solar Atmosphere
What wave modes are present at what heights in the solar atmosphere? What are their sources? What role do waves play in chromospheric and coronal heating? How do the answers to these questions depend on the local magnetic-field structure?
At least two implications promote the use of DKIST to study magnetohydrodynamic waves in the solar atmosphere (e.g. Kostik and Khomenko, 2013): i) MHD waves carry energy into the solar atmosphere and, if dissipated at the correct heights, may provide at least a partial solution to the long-standing coronal and chromospheric heating problems, and ii) the presence of magnetic fields in the chromosphere and corona modify the waves observed, providing a possible opportunity to use them as diagnostics of the conditions there. The chromosphere is a particularly important region of the solar atmosphere, as it modulates wave transmission into the corona; understanding the solar chromosphere is critical to constraining the mechanisms of wave-energy transfer between the photosphere and the corona (see the review by Jess et al., 2015). Moreover, while it is much cooler than the corona, the chromosphere’s relatively high density and the efficiency of the cooling pathways available there mean that high-energy input is required to sustain radiative losses. Typical radiative losses are estimated to be on the order of \(10^{6}\,\text{--}\,10^{7}~\text{erg}\,\text{cm}^{-2}\,\text{s}^{-1}\) in the chromosphere compared to \(10^{4}\,\text{--}\,10^{6}~\text{erg}\,\text{cm}^{-2}\,\text{s}^{-1}\) in the solar corona (Withbroe and Noyes, 1977; Anderson and Athay, 1989). The solution to the coronal-heating problem may well be coupled to or depend on solution of the chromospheric heating problem (e.g. Carlsson, De Pontieu, and Hansteen, 2019).
The possibility that the Sun’s chromosphere and the corona are heated by the dissipation of MHD waves has led to a substantial body of research, starting over 70 years ago with suggestions that acoustic waves generated by convection are responsible for heating the solar atmosphere (Biermann, 1946; Schwarzschild, 1948). With the observational discovery (Leighton, 1960; Leighton, Noyes, and Simon, 1962) and correct theoretical interpretation (Ulrich, 1970) of the resonant solar acoustic oscillations, the diagnostic potential of the solar \(p\)-modes in the study of the solar interior was realized with helioseismology. This was closely followed by an understanding of the observational and theoretical distinction between those modes and propagating atmospheric waves (see Stein and Leibacher, 1974, for an early summary). In the intervening years, studies of chromospheric and coronal waves have leveraged increasingly sophisticated space- and ground-based instrumentation, and the consequent ever-increasing spatial, spectral, and temporal resolution observations, to advance our understanding of chromospheric and coronal wave behavior. Yet fundamental questions about wave-energy transport and wave heating of the solar atmosphere persist (e.g. Erdélyi and Fedun, 2007). What processes dominate wave generation? What modes are excited? How does the energy generated propagate into the solar corona? What is the role of mode conversion? What are the wave-dissipation mechanisms that allow the solar corona to maintain its multi-million-kelvin temperature?
Answering these questions requires the tracking of waves with height in the solar atmosphere, while simultaneously diagnosing changes in wave energy and the corresponding localized atmospheric heating (Jess et al., 2015). In combination with space-based UV observations, DKIST’s unprecedented multi-height spectropolarimetric measurements will be ideally suited to unraveling the nature of the waves, the energy propagation channels accessed, and the atmospheric heating that results.
In a uniform plasma under the continuum approximation, there are three distinct types of MHD wave modes: the slow and fast magneto-acoustic waves and Alfvén waves. Solar observations are often interpreted in terms of these, but this is a dramatic over-simplification, as the manifestation of these waves in the stratified and highly magnetically structured solar atmosphere is much more complex than in a uniform plasma (e.g. Banerjee et al., 2007; Tomczyk et al., 2007; De Pontieu et al., 2007b; Jess et al., 2009; Morton et al., 2012). In simple isolated magnetic geometries such as slabs or flux tubes, modes of the magnetic structures themselves (surface, body, kink, sausage, etc.) can be identified (e.g. Roberts, 1981a,b; Edwin and Roberts, 1982, 1983), but in general, with space-filling magnetic fields, the modes are mixed and coupled, and the waves are subject to resonant absorption, phase mixing, and guided propagation (e.g. Heyvaerts and Priest, 1983; Nakariakov et al., 1999; Bogdan, 2000; Nakariakov and Verwichte, 2005; Aschwanden, 2006; De Moortel, 2009; Jess et al., 2009; Ruderman and Erdélyi, 2009; Goossens, Erdélyi, and Ruderman, 2011; Morton et al., 2011; Wang, 2011; Mathioudakis, Jess, and Erdélyi, 2013; Priest, 2014; Okamoto et al., 2015; Antolin et al., 2015a; Keys et al., 2018). This makes observations difficult to interpret, but also consequently rich in diagnostic potential.
Additional complexities originate with the sate of the solar chromospheric plasma. While the coronal plasma can be treated as a single, low plasma-\(\beta \), fully ionized fluid, the chromosphere is a multi-fluid, partially ionized medium, with a finite spatially varying plasma-\(\beta \), coupled to a radiation field that is out of thermodynamic equilibrium (e.g. Hansteen, Carlsson, and Gudiksen, 2007). Waves in the solar chromosphere are subject to a highly structured (on sub-arcsecond to global scales) evolving field and flow (e.g. Wedemeyer-Böhm, Lagg, and Nordlund, 2009). Since the manifestations and behaviors of magneto-hydrodynamic waves differ in different magnetic and plasma regimes (high \(\beta \) vs. low \(\beta \), structured vs. unstructured (on the scale of a wave) field, partial vs. full ionization, as examples), wave signatures change as the waves propagate upward into the solar atmosphere, and tracking the wave energy from its source to the site of energy deposition is challenging.
Despite these difficulties, significant progress has been made and is anticipated with future observations (e.g. Banerjee et al., 2007). Convective motions in the photosphere excite both longitudinal and transverse perturbations. Of the compressive wave field excited in the photosphere, only waves above the temperature-minimum acoustic cut-off frequency are expected to propagate into the atmosphere above, steepening, shocking, and dissipating as they propagate into the chromosphere (Zhugzhda, Bromm, and Ulmschneider, 1995; Carlsson and Stein, 1997). On the other hand, transverse Alfvénic perturbations can propagate through the chromosphere and into the corona with only weak damping. This allows these waves to reach greater heights but also makes their role in plasma heating somewhat problematic. One possibility is that the waves undergo mode conversion from Alfvénic to compressive, with the compressive motions providing an avenue for dissipation and heating (e.g. Hollweg, Jackson, and Galloway, 1982; Ulmschneider, Zähringer, and Musielak, 1991; Kalkofen, 1997). The height at which mode conversion occurs is then critical to the height-dependent wave fluxes and consequent heating of the upper atmospheric layers.
Observations of compressive wave motions support this general picture with some modification. High-frequency power (near three minutes in period) dominates the chromospheric wave field as expected, but it does so only in limited inter-network regions devoid of strong photospheric or chromospheric-canopy fields (Vecchio et al., 2007; Vecchio, Cauzzi, and Reardon, 2009). In plage regions, compressive wave power in the chromosphere reaches a maximum well below the acoustic cut-off frequency (e.g. Centeno, Collados, and Trujillo Bueno, 2009), and low-frequency magneto-acoustic waves propagate upward from the photosphere in regions surrounding network elements (Jefferies et al., 2006; Vecchio et al., 2007). The presence of magnetic field thus appears to dramatically influence the frequency of the waves propagating into atmosphere. Several mechanisms have been proposed to facilitate this (Roberts, 2000, 2006): inclined magnetic field can effectively act as a wave-guide to reduce the cut-off frequency (Michalitsanos, 1973; De Pontieu, Erdélyi, and James, 2004) and, even for vertically oriented field, nonadiabatic waves can be evanescent outside of magnetic flux tubes yet propagate within them (Roberts, 1983; Centeno, Collados, and Trujillo Bueno, 2006; Khomenko et al., 2008; Centeno, Collados, and Trujillo Bueno, 2009). These mechanisms can be distinguished by careful measurement of the magnetic-field structures supporting the waves and wave temperature–velocity phase relations (Kostik and Khomenko, 2013, but cf. Heggland, De Pontieu, and Hansteen, 2009). More generally, MHD wave properties, and the propagation characteristics and amplitudes of waves of different frequencies, depend in detail on the magnetic structure of the region, radiative energy exchange, and the ionization state of the plasma (Khomenko et al., 2018).
Direct detection of Alfvén-wave perturbations is even more challenging than detection of compressive waves, but Alfvén waves have been successfully observed in the chromosphere (De Pontieu et al., 2007b; Jess et al., 2009), transition region (De Pontieu et al., 2014a), and solar corona, both as time-varying non-thermal line widths (see review, Mathioudakis, Jess, and Erdélyi, 2013) and directly as linear-polarization and Doppler velocity fluctuations (Figure 12, Tomczyk et al., 2007, but cf. Van Doorsselaere, Nakariakov, and Verwichte, 2008). The spectrum of the motions observed in the corona has similarities with that of the solar \(p\)-modes, suggesting the waves have their origin in the solar photosphere (Tomczyk and McIntosh, 2009; Morton, Weberg, and McLaughlin, 2019), although other evidence indicates that they may be more closely related to the ubiquitous Alfvén waves observed in spicules (De Pontieu et al., 2007b; McIntosh et al., 2011a). Observational evidence for mode coupling between fast and slow modes has also been reported. Upward-propagating transverse motions coupled to longitudinal motions subsequently dissipated were identified in quiet-Sun network bright points by careful cross correlation between wavelet-identified wave packets at multiple heights (McAteer et al., 2003; Bloomfield et al., 2004), and evidence for photospherically generated longitudinal magneto-acoustic oscillations propagating upward (Freij et al., 2014) before undergoing mode conversion to predominantly transverse motions has been found in observations of spicules (Jess et al., 2012).
While these interpretations are compelling, to fully disentangle the signal of the different wave modes as they travel upward through the complex solar atmosphere from the photosphere to corona requires simultaneous intensity, polarimetry (magnetic field), and Doppler measurements at many heights in order to deduce the background magnetic field, thermodynamic state of the plasma, and wave perturbations. Such measurements are required for different magnetic-field regimes in order to assess the range of behaviors and quantify the wave fluxes with height. The high-resolution limb spectropolarimetry needed to achieve these measurements is difficult because the low photon counts off the solar limb limit the accuracy of the Stokes profiles deduced at the required cadence and because the low photon counts make the use of adaptive optics, needed to achieve the high resolutions required, challenging. DKIST is poised to meet these challenges. Its large aperture and coronographic capabilities will allow for spectropolarimetric measurements of chromospheric and coronal waves on unprecedentedly small spatial scales and with high cadence.
Flux Emergence into the Non-eruptive Solar Atmosphere
How does magnetic-flux emergence impact energy storage and release in the chromosphere and corona? How are the underlying magnetic-reconnection geometries and heights reflected in observations of small-scale flux emergence/cancelation events?
Magnetic-flux emergence allows for mass, energy, and magnetic field to flow from the solar interior through the photosphere and into the chromosphere and corona above. The impact of flux emergence depends on the amount of field emerging, its spatial distribution, and the pre-existing structure of the atmosphere into which it is emerging. A small bipole will have vastly different impact when emerging into a coronal hole, into the quiet Sun, or adjacent to a \(\delta \)-spot active region. The emergence of magnetic field through the solar photosphere and its reconnection with pre-existing field at different heights thus leads to a variety of dynamic phenomena on different temporal and spatial scales: global magnetic-field restructuring (Török et al., 2014), flaring active regions (Toriumi and Wang, 2019), emerging flux regions or arch filament systems (e.g. Centeno et al., 2017; Su et al., 2018), Ellerman bombs (Figure 13), and quiet-Sun Ellerman bomb-like brightenings. Larger-scale field reconfiguration and destabilization was considered in Section 4.2. The discussion in this section is focused on flux emergence on scales below those of active regions and the local response to that emergence. Tracking the consequences of flux emergence on these small scales allows us to understand the fundamental energy storage and release mechanisms that may be responsible for heating the chromosphere and corona.
Ellerman bombs (Ellerman, 1917) are point-like brightenings seen in the wings of chromospheric lines (such as those of H i and Ca ii). They are often associated with bipolar moving magnetic features around well-developed sunspots or as colliding bipolar structures in emerging flux regions (e.g. Rutten et al., 2013; Reid et al., 2016). The atmospheric heating and associated bi-directional flows observed are thought to be caused by magnetic reconnection near the temperature minimum (e.g. Matsumoto et al., 2008), with rapidly evolving flame-like features in the wings of the Balmer \(\alpha \) line suggesting reconnection of small-scale fields when observed at high resolution (Watanabe et al., 2011). While Ellerman bombs primarily occur in active regions, and are sometimes associated with arch-filament systems (e.g. Zachariadis, Alissandrakis, and Banos, 1987; Georgoulis et al., 2002; Ma et al., 2015, but see Rutten et al., 2013, who suggest these are distinct phenomena), suggesting that larger-scale flux eruption may underlie their occurrence (Pariat et al., 2004), recent observations indicate that smaller, shorter-lived, and lower \({\mathrm{H}}\alpha \)-wing-intensity-contrast events also occur in the quiet Sun. These quiet-Sun Ellerman-like brightenings (Rouppe van der Voort, Rutten, and Vissers, 2016; Shetye et al., 2018) are similar to Ellerman bombs, but with typical size scales of \(\approx0.25\,\text{--}\,0.5\) arcseconds and lifetimes of less than a minute compared to arcseconds and minutes for Ellerman bombs proper (e.g. Roy and Leparskas, 1973; Kurokawa et al., 1982; Vissers, Rouppe van der Voort, and Rutten, 2013). Bright flame-like emission in the wings of \({\mathrm{H}}\alpha \), similar to that observed for Ellerman bombs, suggests a common reconnection origin, but the heating profile and the characteristics of the magnetic-field evolution may imply a somewhat different reconnection scenario (Rouppe van der Voort, Rutten, and Vissers, 2016). Similarly, UV bursts appear to be reconnection events that differ from both Ellerman bombs and quiet-Sun Ellerman-like brightenings in their magnetic topology, atmospheric penetration height, and energy, with the plasma in these events heated to transition-region temperatures (Nelson et al., 2017a; Young et al., 2018; Ortiz et al., 2020).
This full range of small-scale flux-emergence events is well suited for study with the planned DKIST instrument suite. High-resolution, multi-thermal, moderate field-of-view observations can resolve the finely structured thermal properties of the plasma with height, and high-sensitivity spectropolarimetric observations can be used to determine the height-dependent vector magnetic field and Doppler velocities. These quantities in turn will allow estimates of the local electric fields and energy fluxes. Electric-field measurements, derived from time series of vector magnetic-field and Doppler-velocity maps (e.g. Fisher, Welsch, and Abbett, 2012; Kazachenko, Fisher, and Welsch, 2014), are important for determining the rate of electromagnetic-energy transport into the solar atmosphere: the Poynting flux through the photosphere. To date, such measurements have been made only in strong-field regions due to the limited reliability of vector magnetic-field deductions in weak-field regions (Kazachenko et al., 2015). Using DKIST’s unprecedented vector magnetic-field measurement capabilities, electric-field determination can be dramatically improved, and when combined with transition-region and coronal observations can be used to address a range of open questions about the net transfer of magnetic energy into the solar atmosphere: How much magnetic energy reaches the chromosphere? Why are active region cores the sites of the hottest and most dense coronal loops? Is there a measurable correlation between the input of energy at the photosphere and consequent emission in the chromosphere, transition region, and corona? Is the injected energy dissipated immediately, or stored with some typical latency time? Importantly, with DKIST these questions can be addressed as a function of solar activity to uncover the underlying energetics of the solar cycle.
A recent study of a UV burst with the Sunrise balloon-borne telescope (Smitha et al., 2018) revealed dynamic substructure on scales of 75 km, and likely smaller, within a chromospheric heating site. DKIST will not have UV capabilities, but the He i D3 and 1083.0 nm lines may serve as useful proxies when studying plasma at transition-region temperatures (Libbrecht et al., 2017). Coordinated observations with space-based UV assets are also anticipated, and numerical models will play a critical role in data interpretation. Radiative magneto-hydrodynamic simulations of magnetic reconnection during flux emergence show reconnection events similar to Ellerman bombs and other burst events (Danilovic, 2017; Hansteen et al., 2017). The occurrence of these in simulations of emerging active regions suggests a possible role for an underlying emerging large-scale twisted loop structure (e.g. Isobe, Tripathi, and Archontis, 2007; Archontis and Hood, 2009). Detailed modeling of expected DKIST spectropolarimetric measurements in the photosphere and chromosphere will enable careful comparisons between the simulated and observed plasma flows and magnetic-field evolution at the Ellerman-bomb sites (Socas-Navarro et al., 2006; Kondrashova, 2016) to determine if the existence of such large-scale structures is implicated by observations of the Sun. Additionally, direct detection of the brightness temperature excess at Ellerman-bomb sites and in the surrounding atmosphere will be possible with coordinated observations using DKIST and radio instruments such as the ALMA. These will allow careful assessment of magnetic-reconnection heating efficiency in the partially ionized chromospheric plasma, and help clarify the overall importance of Ellerman bombs and other burst and localized brightening events to the energy budget.
Multilayer Magnetometry and Magnetic-Field Extrapolation
How does the magnetic-field change with height and evolve in time through different layers of the solar atmosphere? How do we best use multi-layer magnetic-field observations to constrain chromospheric/coronal field extrapolations?
The stratified solar atmosphere is threaded by magnetic field. Most existing solar instruments employ one or a few spectral lines at a time and thus probe the magnetic field over a limited range of heights in the atmosphere. Moreover, the rapid decrease of the magnetic-field intensity with height implies very weak chromospheric-polarization signals, making their measurement challenging with existing facilities (Schad, Penn, and Lin, 2013). A common need underlying much of the first critical science proposed for DKIST is multi-line high-sensitivity spectropolarimetry. In meeting this need, DKIST will enable simultaneous multi-height measurements of the solar atmosphere that will revolutionize our understanding of the coupling between atmospheric layers.
The solar chromospheric plasma is highly dynamic, inhomogeneous, and out of local thermodynamic equilibrium. Magnetic fields play a central role in its behavior. Observations of the upper chromosphere, particularly near active regions, are dominated by intricate filamentary structures called fibrils. These fibrils are seen in images taken in the cores of strong chromospheric lines, such as H\(\alpha \), Ca ii K, the Ca ii IR triplet (e.g. Hansteen et al., 2006; Cauzzi et al., 2008; Pietarila et al., 2009), and the He i 587.6 and 1083 nm lines (Schad, Penn, and Lin, 2013). They are assumed to be aligned with the magnetic field, but fine-scale chromospheric filamentary structure is detected even above sunspots, where the photospheric umbral field is strong, thought to be largely vertical, and quite uniform; umbral flashes show filamentary fine structure and apparent associated horizontal magnetic field, at or below the scale of current spatial-resolution limits (e.g. Socas-Navarro et al., 2009). Is this an indication of complex field structure or an indication that using fibrils as an indicator of field direction is problematic?
The assumed alignment of chromospheric fibrils with the magnetic field is an important tool in active-region field extrapolation (Wiegelmann et al., 2008; Jing et al., 2011; Yamamoto and Kusano, 2012), and accurate field extrapolation is critical to assessments of the free energy available for solar flares and eruptions, active-region flaring potential, and stability, and models of chromospheric heating. However, ever since the early conjecture by George Ellery Hale that fibrils around sunspots reflect lines of magnetic force (Hale, 1908b), a conjecture made even before his momentous measurement of the field using the then recently described Zeeman effect (Hale, 1908a, see Harvey, 1999, for a brief history), the observational evidence for a direct association between fibrils and the local magnetic-field direction has remained sparse. Because of the small amplitude of the polarization signals within the primarily horizontally oriented (relative to the solar surface) chromospheric fibrils (\(<0.1\%\) in linear polarization), conclusions are wide ranging. For example, attempts to directly measure the alignment between the thermal and magnetic structure of super-penumbral fibrils (Figure 14) have yielded results ranging from often, but not always, aligned (de la Cruz Rodríguez and Socas-Navarro, 2011) to aligned within \(\pm 10\) degrees with no evidence for misalignment (Schad, Penn, and Lin, 2013). A recent Bayesian statistical analysis (Asensio Ramos et al., 2017) finds penumbral and plage fibrils to be well aligned but with non-negligible dispersion. That study concludes that higher signal-to-noise observations are needed to discern whether the misalignment seen in some simulations, particularly those that include ion–neutral coupling (Martínez-Sykora et al., 2016), is compatible with that seen on the Sun. Understanding the three-dimensional connectivity of active-region chromospheric field to the photosphere (at sites such as the outer foot points of super-penumbral fibrils; Schad, Penn, and Lin, 2013) requires high-resolution, high-sensitivity spectropolarimetric measurements. DKIST will enable these.
The connectivity of the magnetic field through the atmosphere is an important issue outside of active regions as well. In the quiet Sun, the photospheric magnetic field is organized by supergranular motions into strong flux concentrations on the network scales and mixed-polarity inter-network magnetic field on the scale of granulation. The field expands above the photosphere into the chromosphere and corona, and the presence of the weak small-scale inter-network magnetic field has a considerable effect on the overall field geometry with height, which deviates significantly from a simple funnel-expansion model (Schrijver and Title, 2003; Aiouaz and Rast, 2006; Martínez-Sykora et al., 2019). This is critical because the magnetic field forms the underlying channel for energy transport into the solar chromosphere and corona, playing an important role in the acceleration of the solar wind (e.g. Gabriel, 1976; Aiouaz, Peter, and Lemaire, 2005; McIntosh et al., 2007; Tian et al., 2008, see also previous sections in this Research Area). Beyond idealized potential or force-free field extrapolations, the variation in field strength and topology with height is typically poorly known.
Force-free extrapolations can be improved. Because of the availability of photospheric magnetograms, field extrapolations generally depend on photospheric boundary conditions, but these boundary conditions are inconsistent with the force-free assumption because both gas pressure and gravity play important roles at photospheric heights. Employing chromospheric-field measurements instead, measurements made at sufficiently great heights, where the magnetic field is much more dominant and is consequently configured much closer to a force-free state (Zhu et al., 2016), can significantly improve the reliability of field extrapolations (Fleishman et al., 2019); when combined with photospheric measurements, even very limited chromospheric-field measurements allow significant improvement (Fleishman et al., 2019). Additionally, careful comparison between independent extrapolations using photospheric- and chromospheric-field measurements can aid in determining relative line-formation heights and in resolving the 180-degree field ambiguity (Yelles Chaouche et al., 2012), and reliable multi-height magnetic-field measurements with DKIST will not only contribute to more reliable extrapolation of that field but will strengthen deductions of the local field at the measurement site.
Magnetic Topology, Helicity, and Structures
What role does the near conservation of helicity play in the structuring of the solar corona and coronal mass ejections? Are measurements of helicity useful indicators of imminent eruption? Do vortex tubes exist and do they act as portals for MHD wave propagation and energy transfer in the quiet solar atmosphere? What are the magnetic-field properties of vortex structures in the lower solar atmosphere?
Magnetic helicity is a property of the field that helps describe its topology, whether it is twisted or linked, writhes or is sheared (e.g. Moffatt, 1969; Berger, 1999; Moffatt, 2014; Blackman, 2015; Moffatt, 2016). It is strictly conserved in ideal MHD (Woltjer, 1958) and during two-dimensional reconnection, and approximately conserved after three-dimensional reconnection (Taylor, 1974; Berger, 1984; Hornig and Rastätter, 1997). Magnetic helicity cascades to larger scales (e.g. Frisch et al., 1975; Pouquet, Frisch, and Leorat, 1976; Pouquet and Patterson, 1978; Alexakis, Mininni, and Pouquet, 2006) and is converted from one form to another (between twist and writhe for example) as it moves to larger scales (e.g. Pevtsov et al., 2014; Knizhnik, Antiochos, and DeVore, 2017; Scheeler et al., 2017; Zuccher and Ricca, 2017). This means that as magnetic fields reconfigure in the solar atmosphere, magnetic helicity is lost only slowly.
Many solar magnetic structures contain self-helicity (internal twisting) and/or mutual helicity (tangling about each other), with helicity observed on the Sun on scales ranging from the largest global to the smallest quiet-Sun magnetic fields (e.g. Pevtsov and Balasubramaniam, 2003; Welsch and Longcope, 2003, and the references therein). The intense magnetic-field structures that form within and rise through the Sun’s convection zone, and are thought to be responsible for sunspots and most solar activity, are likely highly twisted. Untwisted, such tubes would lose their integrity as they ascend. Observations of sunspots show that they rotate as they emerge (e.g. Evershed, 1909; Brown et al., 2003). That rotation is likely associated with an underlying large-scale twisted flux tube rising through the photosphere (e.g. Sturrock et al., 2015), and the helicity that enters the solar atmosphere on all scales from below has important implications for the structure and behavior of the field there.
Some coronal loops appear to be tangled about each other forming a braided pattern (Parker, 1983; Cirtain et al., 2013; Pontin et al., 2017), and the degree of coronal-loop braiding overall has been used to estimate the role of small-scale reconnection in coronal heating (Schrijver, 2007; Knizhnik, Antiochos, and DeVore, 2017). The appearance of braided structures, however, depends critically on the details of the field-line windings within them, with loop substructure difficult to distinguish in observations (Berger and Asgari-Targhi, 2009; Pontin et al., 2017; Li and Peter, 2019), so careful high-resolution spectropolarimetric observations are vital. Further, contrary to expectation, coronal loops have quite uniform width along their length (Klimchuk, 2000; Watko and Klimchuk, 2000). Explanations for the observed lack of expected field expansion rely on loop substructure, either to provide magnetic tension (e.g. López Fuentes, Klimchuk, and Démoulin, 2006) or to allow fine-scale interchange reconnection that enables cross-field loss of the hot loop plasma (Schrijver, 2007; Plowman, Kankelborg, and Longcope, 2009). Untangling these observationally, by leveraging DKIST’s coronal magnetic-field measurement capabilities, is important in understanding these thermodynamic structure of the corona and its maintenance.
Beyond loop substructure and heating, the accumulation of magnetic helicity in the corona appears to be of key importance to coronal mass ejections. Magnetic-helicity accumulation accompanies the magnetic-energy build-up that precedes the loss of stability when a coronal mass ejection is initiated (Zhang and Low, 2005; Zhang, Flyer, and Low, 2006; Yeates and Hornig, 2016). The precise stability implications of the helicity accumulation are still somewhat uncertain (Amari et al., 2003; Phillips, MacNeice, and Antiochos, 2005), and some measures of helicity may be more reliable instability indicators than others (Pariat et al., 2017), but independent of the exact triggering mechanisms, coronal mass ejections associated with filament eruptions often reveal large-scale helical magnetic structures that partially unwind during the eruption (e.g. Kurokawa et al., 1987; Xue et al., 2016). Coronal mass ejections may thus play an essential role in relieving the solar atmosphere of accumulated helicity (Zhang, Flyer, and Low, 2006), and detailed observational assessment of the coronal helicity budget and its role in coronal mass ejection initiation are crucial. Moreover, highly twisted structures typically have a high magnetic energy, and that twist can lead to local instability (e.g. the kink instability), reconnection, flaring, and small-scale eruptions in addition to coronal mass ejections proper. DKIST will significantly enhance our ability to deduce the magnetic helicity in pre-and post-eruptive magnetic fields.
Prominences (or filaments on the disk) are cool plasma structures (at chromospheric temperatures) embedded into the hot corona (e.g. Parenti, 2014; Gibson,