The Solar-B science teams have identified key problem areas to be addressed by this mission. Among those in which the XRT plays a major role are:

1. Flares & coronal mass ejections. How are they triggered, and what is their relation to the numerous small eruptions of active region loops? What is the relationship between large-scale instabilities and the dynamics of the small-scale magnetic field?

2. Coronal heating mechanisms. How do coronal loops brighten? Are wave motions visible, and are they correlated with heating? Do loops heat from their footpoints upward, or from a thin heating thread outward? Do loop-loop interactions contribute to the heating?

3. Reconnection & coronal dynamics. Yohkoh observations of giant arches, jets, kinked and twisted flux tubes, and microflares imply that reconnection plays a significant role in coronal dynamics. With higher spatial resolution and with improved temperature response, the XRT will help clarify the role of reconnection in the corona.

4. Solar flare energetics. Although Solar-B will fly after solar maximum, there will still be many flare events seen. The XRT is designed so that it can test the reconnection hypothesis that has emerged from the Yohkoh data analysis.

5. Photosphere/corona coupling. Can a direct connection be established between events in the photosphere and a coronal response? To what extent is coronal fine structure determined at the photosphere?

The Yohkoh analysis and new results from SOHO and TRACE have clarified the directions in which solutions to the dual problems of structure and stability of the corona may be found, and indicate the types of observations which need to be made in order to address these problems. The corona is found to be highly non-uniform in spatial structure, in temperature structure and also as a function of time [9]. There are a number of conclusions arising out of this work which have implications for the type and design of XRT to be flown on Solar-B:

1. The corona is structured on small scales perpendicular to B, and on large scales along the field. This means that observations must be made with both high spatial resolution and a large field of view.

2. The corona is highly variable on very short timescales. Observations must therefore be made with short exposure times and with a high cadence rate.

3. The corona is multi-thermal, meaning that different structures are seen at only slightly different temperatures, i.e., a small change in temperature sensitivity can lead to a large change in what is observed. Observations must be made over a wide range of temperatures (3-10 MK) with good temperature discrimination, and especially with the ability to detect coronal plasma at about 4-6 MK, where the peak of the differential emission measure distribution is located [e.g. 10].

4. CMEs have an identifiable on-disk soft X-ray signature [11], and events seen in the outer corona can now be traced to an origin on the disk [12]. It is therefore possible that, with higher spatial resolution, large field of view and higher data cadence, the initiation of these events can finally be studied.

In response to these requirements, we are building a full-Sun grazing-incidence X-ray telescope (XRT) that will address the questions raised by the Yohkoh, SOHO and TRACE observations. The instrument is discussed in detail in Section 2, along with technical items such as the resolution, field coverage, required pointing stability and accuracy, and quantity of data returned. In the remainder of this section we discuss the scientific objectives of the Solar-B XRT, the ways in which the scientific requirements lead to the need for specific observations, and the manner in which the XRT will address these objectives.

(i) Coronal Heating

Few problems in astrophysics have proved as resistant to solution as the coronal heating problem. Observational constraints on wave fluxes (based on line broadening measurements) have for all practical purposes completely eliminated the classical acoustic wave heating models. Theoretical models have focused in recent years instead on the various ways in which energy may be transported to the corona, and there dissipated, through the mediation of magnetic fields. Virtually without exception, these models have in common the feature that the actual dissipation of energy transported to the corona occurs in spatially highly localized regions, although there may be either a few or many such regions distributed throughout a typical coronal loop.

Figure 1.1. TRACE image of an active region in Fe IX/X, 19 May 1998.

The interaction between the magnetic fields and the fluid at the photospheric level causes two classes of disturbance, depending on whether the driving timescale is longer than the Alfven transit time across a coronal structure (the DC models) or shorter (AC models):

1. Periodic motions of flux tubes generate MHD waves which propagate upward and may dissipate their energy in the corona. The dissipation is likely to involve phase-mixing: the development of fine-scale structure in the wave's velocity field due to density and/or magnetic-field inhomogeneities [13;14; 15; 16].
2. The random walk of flux tubes produces DC field-aligned electric currents, which may dissipate resistively; this applies only to ``closed'' structures in which magnetic stresses are able to build up over time Parker [17 18)] proposed that the random footpoint motions lead naturally to the formation of ``tangential discontinuities,'' which correspond to thin current sheets; van Ballegooijen [19, 20] described this process in terms of a cascade of magnetic energy to small spatial scales. The current sheets may be distributed more or less randomly within the corona, or may be preferentially located at the interfaces between the flux tubes [21].

Data from the TRACE satellite [22] show directly that there is structure present in the corona at 1 arcsec resolution (Fig 1.1). In particular, we see fine "threads" of hot plasma in the cooler lines, such as Fe IX/X and Fe XII, but not in Fe XV. Although the explanation for this is not clear, it is apparent that some emphasis on cooler material for the Solar-B is appropriate at this resolution. Priest et al.[23] used Yohkoh SXT observations to determine the temperature profile along a large loop. Comparison with models shows that a heating function localized either near the footpoints or near the apex does not fit the observations well, whereas a uniform heating function provides a better fit. The model has been extended to analysis of a loop arcade [24], showing that a constant heat flux for all loops does not provide a good fit, whereas a heat flux varying with B² does. In order to extend the range of applicability of such models, we require observations that can better isolate coronal structures, and that can also observe them with lower errors over a broader range of temperatures.

The Nanoflare Model. Observations of rapid hard X-ray fluctuations [25] and variable emissions from the chromosphere-corona transition region [26] have led to the suggestion that the corona is heated by nanoflares: small-scale reconnection events which release part of the magnetic free energy stored in a coronal loop [27,28,29,30,31]. The energy release likely occurs as an avalanche of such reconnection events [32,33]. For the nanoflares to be energetically important they must be more frequent than predicted by extrapolation of the observed flare energy distribution [34,1]. The dynamical response of a coronal flux tube to impulsive nonoflare heating has been studied by a number of authors (e.g., [35,36,37, 38, 39]). After the initial heating phase, the plasma is extremely hot (T > 10⁷ K) but not very dense. Electron thermal conduction causes ``evaporation'' of chromospheric plasma, leading to a gradual increase of coronal n_e and decrease of T at the loop top. The n_e increase continues until T drops to a few MK, at which point radiative losses become important and the coronal n_e reaches a maximum. Further cooling causes mass to drain out of the tube and return to the chromosphere. Each such heating and cooling cycle requires some tens of minutes [30]. Cargill & Klimchuk [31] have used nanoflare models to interpret observations of active region loops obtained with the Yohkoh SXT. They have found that for hot loops (T > 4 MK) small filling factors can fit the data (f < 0.1), although for cooler loops (T ~ 2 MK) the nanoflare model cannot reproduce the observed temperature and emission measure for any value of the filling factor (also see [40]). Judge et al[41] have studied the correlation between density sensitive line ratios and Doppler shifts of O⁺⁴ emission lines seen with SUMER on SOHO. If it is assumed that the observed correlations are due to wave motions, then they are consistent with downward propagating compressive waves. Detailed models by Wikstol et al [42] show that such waves are a natural result of nanoflare heating.

Specific Objectives. We will combine XRT, FPP and EIS data to study how coronal loops are heated. Specifically, we will determine the emission measure EM(T) as a function of time and position, determine the small-scale structure within the loops, and correlate the observations with magnetic structures seen in the photosphere. Some of the key questions are:

1. What is the emission measure EM(T) of coronal loops in the temperature range 1--10 MK on arcsec spatial scale?

2. How does the EM distribution vary with position along the loop? What is the nature of the pressure gradients found by Kano & Tsuneta [43], and what do they imply about the heating mechanism?

3. How do coronal loops evolve in time? How does EM(T,s,t) evolve as a function of temperature, position and time? Can we confirm that mass is injected by chromospheric evaporation?

4. Can we obtain better observations of the fine structures and temporal variations associated with nanoflare heating? The variations in coronal T and in coronal n_e should have observable effects on the T and EM distributions at 1" resolution, which should be detectable by XRT.

5. Is there a relationship between coronal heating events and spicules seen in the chromosphere [44]? Spicules may be a chromospheric response to nanoflares in the corona [45,46]; coordinated observations of the XRT, EIS and optical instruments will determine whether such a relationship exists.

6. Can we detect MHD waves in coronal loops? At high cadence we will search for intensity fluctuations associated with compressional waves, and for undulation of fine threads associated with transverse waves.

(ii) Flares & CMEs Traditionally, there are three different types of large-scale eruptive phenomena occurring in the solar atmosphere, namely coronal mass ejections (CMEs), prominence eruptions, and large two-ribbon flares. It has become increasingly clear with time that these phenomena are closely related and may be different manifestations of a single physical process. The opening of the field lines in the active region by the CME leads to the formation of flare ribbons and loops, appear to move through the chromosphere and corona, and these motions provide some of the best evidence for magnetic reconnection in the solar atmosphere. Doppler-shift measurements show that the motions of the flare loops and ribbons are not due to mass motions but rather to the upward propagation of an energy source in the corona, as required by the reclosing of open field lines by reconnection [47].

 

Figure 1.2 Flare Model

Figure 1.2 is a diagram showing one proposal for how reconnection occurs during the gradual phase of large flares and CMEs. It is based on models [48,49,50,51]; on simulations of magnetic reconnection [52]; evaporation [53,54,55,56,57]; and condensation [58]. According to this scenario, flare loops are created by chromospheric evaporation on field lines mapping to slow-mode shocks in the vicinity of a neutral line. Conduction of heat along the field lines causes them to dissociate into isothermal shocks and conduction fronts as shown in the figure. The shocks annihilate the magnetic field in the plasma flowing through them, and the thermal energy which is thus liberated is conducted along the field to the chromosphere. This in turn drives an upward flow of dense, heated plasma back towards the shocks, and compresses the lower regions of the chromosphere downward. Until the advent of Yohkoh, virtually all the evidence for reconnection on the Sun was indirect. However, the high resolution and sensitivity of Yohkoh SXT made it possible to see the reconnection region directly for the first time. The detection of a cusp-type geometry at the top of flare loops along with the detection of a nonthermal X-ray source in the same region now provides some of the best evidence that a reconnection site does actually occur in the corona.

To determine whether the reconnection process occurs in the manner proposed in Fig. 1.3, one must observe the changes in shape of reconnected field lines with time. Because flare plasma on reconnected field lines undergoes an enormous temperature variation from 10⁷ K to 10⁴ K, no single instrument has been able to track continuously the plasma as it cools. The XRT will make a major advance by observing at arcsecond resolution the cooling of the X-ray loops down to a temperature of 10⁵ K, an order of magnitude better than achieved by the Yohkoh SXT.

Coronal Shock Waves. There is still considerable debate about the number, origin, and structure of shock waves produced by coronal mass ejections (CMEs). In interplanetary space, only a single, fast-mode shock wave is seen in front of the ejecta (magnetic cloud) thrown out by the CME, but when and where this shock originates is not yet known with any precision. Indirect evidence for the existence of shock waves in the lower corona is provided by ground and space observations of Moreton waves and radio observations of metric type-II bursts, but it is far from clear whether these shocks are the same as those seen in interplanetary space [59,60]. The improved resolution and sensitivity of the XRT will allow better detection of coronal shock waves. In fact, the XRT should be able to observe directly the 3-D structure of shock waves as they propagate through the corona. A Mach 2 fast-mode shock propagating across the magnetic field increases the plasma density by a factor of ~2 and the temperature by a factor of ~80 (for a plasma b = 0.01), but immediately behind this heated region lies a rarefaction wave which progressively reduces both n_e and T with increasing distance from the shock. Thus, the combined shock-rarefaction wave has a unique density-temperature signature which the XRT will be able to detect with a sensitivity and resolution that has not heretofore been possible.

These observations will help resolve three long-standing scientific issues. First, determining the precise region in the corona where the shock originates will tell us the location, and extent of the driving force of a CME. For example, if the shock originates from a volume which is larger than any active region or prominence, then we will know that the j x B forces which drives an eruption is not created by a local magnetic instability within the active region or the prominence. Second, knowing the shock strength as a function of its location in the corona can help determine whether the proposed relation between CME shocks and prompt (< 30 min) energetic particles (> 1 Mev) by [61] is correct. Finally, if interplanetary shocks are distinct from the shock waves generated near the surface, then the XRT should be able to detect signs that such multiple shocks actually exist and determine when and where they are created and dissipated relative to one another.

Flare and CME Energetics. Most models for eruptive flares and coronal mass ejections are based on the principle that the energy which drives them comes from magnetic energy stored in coronal currents [62]. The currents may form when a flux-tube emerges from the convection zone or when the footpoints of a pre-existing arcade are sheared. Since magnetic helicity is a well preserved quantity in the corona (see [63] or [64])), only part of the stored magnetic energy can be released during a confined flare. Large eruptive events can remove helicity from the corona by ejecting flux ropes[65], but this mechanism of helicity shedding is severely constrained by the fact that in a magnetically dominated medium the fully open field has maximum energy. Consequently, the field cannot be opened by an MHD instability, ideal or otherwise.

Figure 1.3  Schematic view of magnetic field expansion from photospheric elements into larger coronal structures.