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| EIS Science Team Report |
The purpose of these studies is to make it easier to make critical decisions on the design. The arguments are based on "Do we have enough counts/temporal resolution/spatial resolution/spectral resolution to achieve our scientific objectives?"
Two reports are given below discussing the observing possibilities for two major flare topics - chromospheric evaporation (George Doschek) and reconnection (Peter Cargill).
Title: Chromospheric Evaporation in Flares
Author: George Doschek
Justification:
Although the idea that the coronal mass of flares is supplied by chromospheric evaporation is generally accepted, there are many problems in understanding the process, both experimentally and theoretically. For example, simple 1D models of evaporation fail to fit the data (Mariska et al., ApJ, 419, 418, (1993)). The sites of evaporation are difficult to determine, mainly because BCS is uncollimated and obvious evaporation mass flows are not seen in SXT flare images (e.g., Doschek et al., ApJ,, 431, 888 (1994)). Flares that exhibit the largest evaporation signature on the profiles of X-ray lines seem to be morphologically complex even in SXT images, and from recent TRACE images of flares they would look incredibly complicated (e.g., Doschek et al., ApJ, 459, 823 (1996)). Finally, the bright flare loop tops are not easily explainable by any simple flare loop model (e.g., Feldman et al., ApJ, 424, 444 (1994)).
EIS Capability:
EIS is capable of observing the hot flare plasma (Fe XXIV, Fe XXIII, Ca XVII), of determining the footpoint locations of flare loops (He II), and of observing Doppler motions as a function of position in the flare. These three observations allow the sites of evaporation to be determined as a function of time in the flare. They also allow the magnitudes of the motions to be determined as a function of position and time, and also give some idea of the temperature of the upflowing plasma. These results would provide a quantum leap in our observational understanding of energy transport in flares.
EIS Wavebands and Spectral Line Parameters:
The necessary bandpasses for this work are either NRL1 or the baseline band. Only these bands contain the hot flare lines. The footpoint locator line in NRL1 would have to be the O V transition region line seen during the impulsive phase of flares (Widing, Skylab data).
From X-ray and EUV Skylab spectra of flares we can assume an electron temperature of 12-20 MK for the Fe XXIV line and a turbulence that decreases monotonically from about 160 km/s to 60 km/s in the decay phase. For a disk flare, there will be a stationary line component and a blueshifted upflow component. I assume the same temperature for each component with some justification (see Doschek et al. ApJ, 459, 823 (1996)).
If I adopt a temperature of 15 MK, the FWHM of the stationary component Fe XXIV line at 192 A will vary from 0.185 A to 0.0958 A. The thermal width is 0.0711 A. The Fe XXIV line at 255 A will have widths that are 255/192 = 1.33 times larger.
Upflow speeds of the blueshifted component range from a few km/s up to 800 km/s, with an average centroid shift of say 250 km/s. For a count rate estimation, let us assume an upflowing component FWHM of 290 km/s or 0.31 A (from Doschek et al., PASJ, 44, L95 (1992)). This is for a strong blueshifted signature event. For the more average event Mariska et al. (see above reference) found a centroid shift of about 50 km/s. Let us furthermore assume a FWHM of 0.31/3 = 0.1 for this more average event.
The intensity of the blueshifted component relative to the stationary component decreases monotonically with time. For a strong signature event, at flare onset it can dominate the total line profile. For an average event it is about 20% as strong as the stationary component at flare onset. Note that only the strong signature events produce spectral line signatures that begin to agree with numerical simulations. These constitute only about 10% of the total number of events (Mariska et al. (1993)).
The Ca XVII line will also exhibit dynamical signatures, but the main line for dynamics should be Fe XXIV, both because it is hot and it is much stronger than than the Ca XVII line.
Observing Strategy:
Plans for coordination between the Solar-B instruments and pointing the instruments at specific targets have not yet been worked out. In the absence of any guidance I assume that the XRT will provide a flare flag similar to SXT. One can envision a flare program where the partial FOV instruments are pointing over a highly sheared neutral line in an active region that is prediced or known to be producing flares. When a flare occurs in the active region the brightest pixel location is provided by XRT and EIS is moved to the flare location. At the flare location EIS can obtain Fe XXIV images using the slot and thereby function as the Skylab S082-A instrument, or it can raster the flare region like CDS or SUMER.
Time Resolution:
The rise phase of a flare is the significant factor in this evaporation study. It varies from about 1 minute to as much as 20-30 minutes. Blueshifts have been seen in a 30 minute rise phase flare (Doschek et al., ApJ, 345, 1079 (1989)). For the intense blueshift events observed with Yohkoh (Doschek et al. (1996)) the rise times were only a few minutes. Therefore time resolution on the order of 10s is required to adequately observe the rise phase.
Count Rates -- Cassegrain:
Count rates are given for an approximately M3 flare observed on August 9, 1973 from Skylab.
NRL1: The predicted photon detection rate in the 192 A Fe XXIV line at flare maximum is 53,000 photons/s/ spatial pixel. For the Cassegrain, the dispersion is 10.5 mA/pixel. Consequently, the number of pixels contained in the FWHM of the line, which is only the stationary component at the peak of the flare, is about 0.0958 x 2/0.0105 = 18.25 pixels. Thus, the average count rate per pixel is roughly 53,000/18 = 2,900 counts per dispersion pixel/s. This count rate is quite large and EIS can therefore provide unprecedented soft X-ray spectral resolution near times of flare maximum.
We desire to start measuring blueshifts when the total flare flux is very roughly two orders of magnitude below the peak flux. For an intense blueshift event, assume that the blueshifted component dominates. The the average count rate per pixel is 53,000/100/(0.31 x 2/0.0105) = 9 counts per dispersion pixel/s. In 10 s we have about 100 counts per pixel which is 10% counting statistics. Thus EIS is easily sufficiently sensitive to observe intense blueshifts at flare onset. For weak blueshifts the predicted average count rate is 53,000/100/3/(0.1 x 2/0.0105) = 9 counts, about the same as for the prominent blueshifts. This comes about because although the blueshifted component is never dominant (hence the factor of 3 above), it also is not as broad in wavelength and therefore the FWHM in pixels is less. The two effects tend to compensate each other.
Conclusion: EIS can do the evaporation problem using NRL1 and Cassegrain optics.
Baseline: The baseline observes the 255 A line in first order. The predicted peak count rate is 5,565 photons/s/spatial pixel. The dispersion I take as 19 mA/dispersion pixel, although this is strictly for NRL2 at 320 A. The predicted count rate at peak is 5,565/(0.0958 x 1.33/0.019) = 830 counts per dispersion pixel/s, which is a lot of counts.
For flare onset, the count rate is 5,565/100/(0.31 x 1.33/0.019) = 2.6 counts/dispersion pixel/s. In 10 s there are 26 counts, which is not too large. However, one can observe the upflows well into the rise phase. So for simply identifying the source of upflows with footpoint regions as seen in He II 256, the baseline band should be adequate.
Conclusion: EIS can do the evaporation problem using the baseline waveband and Cassegrain optics.
Title: Flares -- Coronal Reconnection
Author: Peter Cargill (ICSTM)
Justification:
Since Petschek proposed his famous mechanism in 1964, magnetic reconnection has been seen as a viable mechanism for solar flare energy release. However, it was not until the launch of Yohkoh that good evidence for reconnection in the corona was obtained from the SXT and HXT instruments. Even so, the Yohkoh results could not produce the essential evidence for reconnection namely the appropriate mass motions. The reconnection process involves the conversion of magnetic energy into thermal and kinetic energy, in perhaps roughly equal parts. In addition, reconnection involves the flow of cool, unreconnected plasma into the reconnection site. Thus, spectral lines in the appropriate temperature ranges should be Doppler shifted. EIS provides the first opportunity to observe explicitly such motions. This is in fact quite a difficult observing sequence to construct. The inflow into the reconnection site is likely to be parallel to the surface of the Sun, while the high speed outflow is directed either downward or upward. So one would appear to be restricted to observing one or the other, though tilting of the reconnection site could alleviate this. The inflows would be best observed on the limb and the outflows on the disk. The latter could thus be combined with the previous entry on evaporation.
Possible Strong Lines: (based on statements in Harrison/Mason report)
Range 1: Fe XXIV (192.02: possible blending)
Range 2: Fe XXIII (263.76: possible blending) and Fe XXIV (255.10)
Range 3: Fe XXI (335.9) and Fe XXII (349.3) Strength????
Range 4: not being considered
Range 5: not being considered
Range 6: Nothing.
Temperatures: Adopt 2.5 10**7 as a generic value.
Emission Measures and velocities:
These are related, as noted by Klimchuk (1998). Adopt his approach. Assumes that reconnection produces jet at Alfven speed based on reconnection field. For typical coronal values of parameters, (density 10**10, 100G field, 10**31 erg flare in 100 secs) this gives for a 1" by 1" pixel:
V = 2000 k/s, EM=1.5e45 cm**-3
These velocities are rather big. I feel more comfortable with a smaller coronal velocity (say 25% of this). You then get
V = 500 k/s, EM=5e46 cm**-3
One could expect to also see turbulent line-broadening with such velocities. Note that the range of velocities and emission measures that need to be considered is very uncertain -- perhaps 2 orders of magnitudes.
Strategy:
This is a very challenging observation to make. Given the launch of Solar-B at solar minimum, flares will still be present, but not as common as one would like them to be. Success is a case of being lucky enough to be in the right place at the right time.
The obvious implementation would be to be positioned looking at a active region that was likely to flare (it would make sense to organise this in conjunction with the optical telescope so that any flare-causing photospheric disturbances would be seen.) There would need to be some sort of flare flag to start the sequence. When this was detected, EIS would go there directly. The problem is with knowing whether the brightest region is associated with the reconnection jets, or is displaced somehow. The former case is the more favourable. One can then afford to raster the slit over a fairly small area, hence reducing any problems with low count rates. One could envision wanting to make observations for 100 secs (the impulsive phase).
It is worth noting that the best chance to see jets is with disc flares -- on the limb, the flows will probably be up or down, with a small line of sight component. On the disc, one can expect to be looking along (or at least partially along) the line of sight, so that one would be looking down through the jet. This would probably increase the above emission measures quite significantly. Another important point about disc flares is that if one believes the standard Yohkoh picture of the hot X-ray source below the reconnection site, rastering the slit over a small diatance may work.
If the jets are displaced from the peak emission, things are more difficult, though the above arguments suggest that this may not be the case. One then has to try and scan as big an area as possible in the hope of seeing them. In this instance, one would find the bright location and then scan over say 1' in each direction within 100 seconds.
Another issue is the strength of the moving to stationary components. Clearly both will be present, and estimates are needed of their relative importance. In addition, the thermalisation of the jets will lead to a broad range of velocities being present rather than just that of the jet. It is probable that the lines will exhibit rather length wings rather than a clear Doppler shifted component. Of course detection of this would in itself represent a significant achievement.
Summary:
There is a real chance to make a major breakthrough in the study of flares with EIS on Solar-B. However, the predicted count-rates are based on our very ill-defined picture of how magnetic reconnection operates in flares. One can imagine the emission mesaures in particular changing by an order of magnitude or more in either direction, far greater than the factor of 3 between a baseline design with 2" pixels, and the sum of 4 1" Cassegrain pixels (I think these are the numbers I got from Louise). One can imagine a situation where either jets appear clearly, or are not seen at all, independent of the chosen spatial resolution. Thus these possible observations should not be driving the instrument design. Rather they should be viewed as a wonderful potential additional benefit from an observatory designed to study the connection between the photosphere and corona. A spatial resolution as close as possible to that of the optical telescope would appear to be desirable.
Title: Active Region Heating
Author: L.K. Harra-Murnion (MSSL), S.A. Matthews (MSSL) and H. Hara (NAOJ)
Scientific Justification:
There has been much recent work investigating active region heating with Yohkoh, SOHO, TRACE and ground-based observatories. It has been found that there are transient brightenings occurring frequently at coronal temperatures (Shimizu) and that these brightenings do not provide enough energy to heat the active region. It has also been observed that there are small flows in coronal loops of up to 10 km/s. Measurements of line broadening have been used to determine whether some form of wave-heating exists. There have been indications of Alfven wave heating in measurements of line broadening with height in an active region observed in Fe XIV using the Norikura Observatory (Hara and Ichimoto, 1999). SOHO-CDS has shown that cooler loops are also exist at the same heights as coronal loops. These cool transition region loops are extremely dynamic. There is a strong suggestion that the line-broadening is related to the dynamical behaviour of the loops. There is also a qualitative suggestion that, although the cool loops appear very different in appearance, that they are a route to the heating of the hotter coronal material.
EIS can for the first time will provide information on the velocity field in a wide range of temperature lines. This will provide us with the possibility of understanding the heating mechanism(s).
The velocity resolution should be on order of several km/s for the coronal lines. It is not so crucial for the transition region lines are the observed flows are much higher. The temporal resolution should be no more than 5 minutes for as large a raster area as possible. The desired temporal resolution would be 1 minute (or less!!).
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Raster Area: as large as practical - 4 X 4 arcmins
Raster step: 2 arcsecs
Raster Locations: 120
Exposure time: 2 secs
Duration of Raster: 4 mins + unknown for overheads
Number of rasters: for several hours on an active region on the limb, and then again for an active region on the disk.
Line Selection (this covers the strongest transition region and coronal lines in all the four wavelength regions):
| Ion | Wave | Log T | Ct/s Cass | Cts/s OAP |
| Fe XII | 186.8 | 6.2 | 13 | 443 |
| Fe XII | 193.5 | 6.2 | 231 | 3039 |
| Fe XI | 188.2 | 6.1 | 51 | 913 |
| He II | 256.3 | 4.7 | 38 | 635 |
| Fe XVI | 262.9 | 6.3 | 27 | 519 |
| Fe XIV | 264.8 | 6.2 | 39 | 698 |
| Mg VI | 270.4 | 5.6 | 11 | 162 |
| Si VII | 275.3 | 5.8 | 13 | 254 |
| He II | 303.8 | 4.7 | 813 | 5679 |
| O IV | 306.6 | 5.2 | - | 11 |
| Mg VI | 314.7 | 5.6 | - | 37 |
| Si VIII | 319.8 | 5.9 | 39 | 630 |
| Fe XVI | 335.3 | 6.3 | 358 | 6638 |
| Mg VI | 400.7 | 5.6 | 31 | 330 |
| Ne VI | 401.9 | 5.6 | 61 | 653 |
Telemetry: 8 lines x 120 (spatial) bins x 24 (spectral) bins x 12 bits =276,480 bits per exposure. At 64 kb/s would be 4.3 secs.
Supporting Observations: The most important aspect for the disk observations of the magnetic information at the resolution provided by the optical telescope on solar B. It is crucial to have a cool line which will permit alignment. One of the biggest mysteries about the cool loops is how they coexist with the hotter coronal loops. If they are completely different structures this should be apparent in the magnetic data.
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What's new?
Fast temporal resolution alongwith good spectral resolution for the line broadening measurements, alongwith high resolution magnetic information.
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For this active region study, the best single wavelength range is the baseline. The reason is because you get a wide coverage of temperature from log T=4.7 to 6.3. If two wavelength ranges were permitted then the EIS400 and NRL1 would be preferred, since they give higher count rates.
Also, a preferred set up would be high spatial and time resolution for the transition region (since the structures are filamentary and dynamic), and for the corona, high velocity resolution is preferred over spatial and temporal. Matt and Alec think that this option can be easily implemented onboard, which would increase the amount of telemetry available.
Title: Active Region Science
Author: John Mariska
Overview
Unquestionably, the most important issue in the area of active region science is, how is the corona in active regions heated? Clearly, no single observing sequence in any of the wavelength bands being considered for EIS can solve this problem.
Heating theories for the active region corona essentially boil down to loop heating models that involve either waves or magnetic reconnection. In the case of wave heating models, one would expect the propagation and dissipation of the waves to result in oscillatory phenomena in the coronal loops. One could consider looking for evidence of the oscillations in temporal measurements of line intensities, widths, Doppler shifts, and even loop position along the slit for a slit placed perpendicular to the direction of the loop.
In the case of reconnection models, the possibilities are equally diverse. If heating is the result of a small number of relatively large reconnection events, then it may make sense to search for inflows or outflow jets through Doppler shift or line width measurements. If, instead, heating is the result of large numbers of small reconnection events, then it makes more sense to think in terms of characterizing the average properties (line intensities, widths, and Doppler shifts) of many relatively isolated loops to look for properties that can be compared with the predictions of models or used to develop new models.
Since theories for reconnection jets suggest that they will be low emission measure features, looking for the signatures of single large active-region reconnection events would require the highest possible throughput. For all the other heating scenarios, one is attempting to deduce the average properties of active region loops. In this case throughput translates into the issue of the minimum integration time one can use. This then determines the range of accessible temporal frequencies. Wave heating theories suggest that observing time scales of a few seconds to perhaps a few tens of seconds will be necessary.
Spatial resolution is a final key issue for all active region observations. TRACE observations of active regions at coronal temperatures clearly show that many loops have very small cross- sectional sizes. Moreover, many of these thread-like structures often exist in close proximity to each other. Any attempt to observe the average properties of "isolated" active region loops or to observe the temporal history of the properties of such a loop could produce nearly impossible to interpret data if there are too many separate emitting structures in the line of sight.
Possible EIS observations
One way to approach these issues is to ask, what are the shortest integration times for the strongest coronal lines in each wavelength channel that will provide reasonable characterization of the line intensity, width, and Doppler shift? Below, I list the five strongest lines in each channel and the count rates and times to detect 1000 photons in an active region in the line for each instrument configuration.
NRL1 (184-204)
| Ion | Wave | T | Cass Ph/s | t(s) | TR Ph/s | t(s) |
| Fe XI | 188.232 | 6.10 | 42 | 24 | 605 | 2 |
| Fe XII | 192.393 | 6.15 | 62 | 16 | 687 | 1 |
| Ca XVII | 192.819 | 43 | 23 | 470 | 2 | |
| Fe XI | 193.521 | 6.10 | 191 | 5 | 2022 | 0.5 |
| Fe XII | 195.118 | 6.15 | 350 | 3 | 3600 | 0.3 |
Base (240-290)
| Ion | Wave | T | Cass Ph/s | t(s) | TR Ph/s | t(s) |
| Fe XVI | 262.984 | 6.30 | 27 | 37 | 396 | 3 |
| Fe XIV | 264.780 | 6.25 | 35 | 29 | 495 | 2 |
| Fe XIV | 270.507 | 6.25 | 26 | 38 | 372 | 3 |
| Fe XV | 284.160 | 6.25 | 76 | 13 | 1765 | 0.6 |
| He II | 303.780 | 4.75 | 37 | 27 | 1639 | 0.6 |
NRL2 (303-340)
| Ion | Wave | T | Cass Ph/s | t(s) | TR Ph/s | t(s) |
| Si XI | 303.324 | 6.20 | 36 | 28 | 269 | 4 |
| He II | 303.780 | 4.75 | 623 | 2 | 4633 | 0.2 |
| Mg VIII | 315.039 | 5.95 | 32 | 31 | 320 | 3 |
| Si VIII | 319.826 | 5.95 | 34 | 29 | 391 | 3 |
| Fe XVI | 335.410 | 6.30 | 299 | 3 | 1942 | 0.5 |
EIS400 (385-421)
| Ion | Wave | T | Cass Ph/s | t(s) | TR Ph/s | t(s) |
| Mg VI | 400.666 | 5.75 | 35 | 29 | 284 | 4 |
| Ne VI | 401.136 | 5.70 | 26 | 38 | 208 | 5 |
| Ne VI | 401.926 | 5.70 | 67 | 15 | 544 | 2 |
| Mg VI | 403.310 | 5.75 | 49 | 20 | 406 | 2 |
| S XIV | 417.660 | 26 | 38 | 371 | 3 | |
Summary thoughts
There are many unanswered questions associated with coronal heating in active regions. I believe that we should not tie ourselves to the ones now most in vogue. Rather, we should try to come up with an instrument concept and wavelength bands that will allow us to search for the consequences of all the possible active region coronal heating theories. The classic TRENDY design has coronal lines in each wavelength band under consideration with count rates of better than 100 per second. The cassegrainian design only achieves this goal in the NRL1 and NRL2 bands. If only one band were selected, the count rates would double for the both designs and at least one coronal line in each band would exceed 100 counts per second in each of the possible bands. In terms of overall temperature coverage, the baseline band (240-290) would provide the best coverage, but would not provide the strongest lines for active region loop studies.
While maximum count rates would be achieved with TRENDY, I believe that the lessons of TRACE, SUMER on SOHO, and all of the inferences that have been made about filling factors from earlier missions argue strongly for going to 1 arcsec spatial pixels if at all possible. This is even more important if we are forced to one wavelength band and it does not provide robust density diagnostics for use in determining filling factors.
Overall, the above discussion seems to point toward going to a cassegrainian design to achieve spatial resolution and selecting only one wavelength band to achieve maximum throughput.
Title: Quiet Sun
Authors: Richard A. Harrison, Dave Pike, Helen Mason and Peter Young
1. Introduction
The aim of this study is to consider quiet Sun structure and activity and define a set of observational requirements which must be met to achieve a significant step forward. These should be considered in the light of the various wavelength and resolution characteristics being considered for the EIS instrument on Solar-B.
2. Scientific Case
A scientific case, specifically to address the interpretation of quiet Sun transient activity - blinkers, explosive events etc... - is given below:
Coronal heating and solar wind acceleration are ultimately powered by the kinetic energy of the convection layers. Magnetic fields carried by the convection cell flow migrate to the cell boundaries (the network) and the combination of flux concentrations and newly merging flux in these regions probably drives magnetic transient events which provide the acceleration and heating processes. If this picture is true, the resulting, globally distributed 'disease' of mini-exposive events may hold the key to coronal heating and solar wind acceleration mechanisms. SOHO has been used to observe transient events in the network which fit this picture, i.e. the CDS blinkers and the SUMER explosive events. The former are EUV flashes especially in the few hundred thousand K region lasting typically 10 minutes, with a few tens of thousands distributed over the disc at any time. No significant velocities have been associated with these events. However, the explosive events are UV velocity events seen in the network, with a similar number on the Sun at any time, with speeds of up to 150 km/s. The two classes of network transient events maybe related, but attempts to show this have, thus far, failed. This is most likely due to an averaging effect due to the size of the CDS 2 arcsec pixels and the lack of temperature coverage for SUMER. In addition to this relationship, we need to invesigate further the effects of these events in the corona and the related magnetic activity in the regions below. This would require a good temperature coverage but would demand good spectral resolution to identify wave broadeing as well as flows emanating from the blinker/explosive event site. Such observations are essential to determine the ultimate influence such events have on the surrounding atmosphere.
3. Observational Requirements
From the science case, we can make some comments on the observational requirements - leaving the wavelength selection to the next section:
Spatial resolution - 1 arcsec or better per pixel * Must be better than CDS
Temporal Resolution - Must be able to raster over areas of several tens of arcsec squared in minutes. If we assume a 1 arcsec slit with 30 steps, this means we need exposure times of order 1-3 sec giving good statistics in the prime lines of interest.
Spectral Resolution - Must be able to determine flows (wave activity) of order 10 km/s or better.
4. Wavelength Selection
First point to note is that the blinkers are not too clear in He I or He II compared to, say, O III, O IV and O V and the gap in temperature between the He lines and the coronal lines is so large that this study would demand some intermediate temperatures.
In addition, CHIANTI is an extremely useful tool, but we know that (a) the solar atmosphere is extremely variable and (b) the atomic data in CHIANTI cannot be considered to be accurate enough to predict the intensities of all lines. Thus, much of the work below is based on lines observed in the Sun rather than those predicted by a model.
4.1 184-204 A Range (NRL 1)
This band is fine for coronal diagnostics and coverage, and is good for pre-flight calibration. However, the temperature range is very poor with only weak transition region lines. This is not good for the quiet Sun research, UNLESS it is one of two bands with the other band providing good transition region lines. The coronal lines themselves are bright and thus lend themselves to the resolutions mentioned above, but this is not sufficient.
4.2 240-290 A Range (Baseline)
This band is fine for pre-flight calibration and it does have a reasonable temperature range. The transition region coverage is not brilliant, but it is there. The quiet Sun studies would have the higher temperatures in the Fe IX 244.92, Fe XIII 251.94, Fe XIV 264.78, 274.20, and Fe XVI 262.98 lines. The cooler and intermediate lines are covered by He II 243.03 (256 is blended), Mg VII 278.41, 280.74, for example. CDS observations would suggest that you would need exposures of quite a few seconds to obtain reasonable statistics in these lines. Everything is driven by the weaker lines, which appear to have counts of order 0.1-0.2 count/s per pixel. This being the case, we would require exposures of 50-100 s for 30% counting stastics. One could improve on this by summing intensities between pixels for the weaker lines, but that defeats the object a little.
4.3 303-340 A (NRL 2)
This band has only limited pre-flight calibration possibilities, with only one set of lines at around 313 A available. On paper it has a good temperature range, from the He II at 303 A to the Fe XIV at 334.17. For coronal studies, the band is not so bad, with a few lines from Fe XIII (320.80), Fe XIV (334.17) visible (using CDS) even outside active regions and some lines (e.g. Fe XVI 335.40) which become bright in active areas. For transition region analyses you are dependent on the lines from Mg VIII and Si VIII in the 310-320 A range. The CDS intensities, folding in a reasonable expectation of increased EIS sensitivity, would suggest that these lines would be in the 3-6 count/s per pixel level. Thus, we could live with few second exposures - if 30% to 20% counting statistical errors are OK.
4.4 385-421 A (EIS 400)
This band is OK for calibration pre-flight. It covers the transition region temperatures well and the intensities from CDS and the CHIANTI calculations suggest that we could use few second exposures. The big problem is the lack of coronal lines. We need this to assess the full impact of the quiet Sun events and their ability to heat plasma to coronal temperatures. Thus, this band would need to be matched to a second band with higher temperature lines.
5. Conclusions/A Sample Observation
For EIS quiet Sun studies a combination of the EIS 400 band and the Baseline or NRL1 bands would be ideal. 1 arcsec resolution is critical and spectral reolutions of order 10 km/s would be required. We would need to consider exposure times of order a few seconds or less, to drive rasters of minutes at the worst. A sample raster is given below:
Study Details:
Raster Area: 1x1 arcmin
Raster Step: 1 arcsec
Raster Locations: 30
Exposure Time: 3 sec.
Duration of Raster: 3x30 plus 10% overhead = 100 sec. (*Overhead = CCD readout time etc...)
Number of Rasters: Minimum 500 (= 833 min)
Line Selection: Range of temperatures using bright lines to keep raster repeat times low. Ensure they are well separated to reduce blends and enable good velocity studies. Also, some density capability. e.g.:
From NRL1 -
Fe IX 171
Fe X 174
Fe X 177
Fe XI 180
Fe XI 188
Fe XII 186 } Density pair
Fe XII 193 }
Fe XIII 202 } Density pair
Fe XIII 203 }
Fe XIV 211
From Baseline -
He II 243.03 (256 blended),
Mg VII 278.41 } Density pair
Mg VII 280.74 }
Fe IX 244.92
Fe XIII 251.94
Fe XIV 264.78
fe XIV 274.20
Fe XVI 262.98
From EIS 400 -
Mg VI 399
Ne VI 399
Ne VI 401
Ne VI/Mg VI 403
Na VIII 411
Ne V 416
C IV 419
(Say 16 lines in total in final selection)
Bins Across Line: 25 (To cover about +/- 250 km/s)
Telemetry/Compression: 16 lines x 25 bins x 120 bins x 12 bits(?) = 576,000 bits per exposure. At 64 kb/s would take 9 s. To keep within 3 s exposure time, would require compression factor of 3. (* This assumes we take the slit only, in the dumb-bell case.)
Solar Feature Tracking: YES (* Must be able to make steps of less than pixel size to avoid jumpy movies).
Supporting Observations: 1. Magnetic Mapping of source region - This would be from the Solar-B magnetic imager. Observations of the same region with a similar cadence and similar or better spatial resolution would be ideal, in order to map the magnetic changes to the transient EUV activity.
2. Context Mapping of Coronal Structures - The EUV images would be relatively small, in order to have reasonable temporal resolution. Larger area coronal maps would be useful to assess the local and remote magnetic structure and activity.
6. What is New?
This problem has been tackled using a combination of CDS and MDI data from SOHO - with some success. However, to perform the same kind of operation with significantly better temporal, spatial and spectral resolution would be essential. In particular, the association of EUV blinkers and UV high velocity events must be sorted out.
Title: Heating of the Quiet Sun
Author: Ken Dere
The source of the energy, that is dissipated in the corona and radiated away into space, lies in the solar photosphere. The motions of the dense photospheric plasma continually buffet and convect the footpoints of coronal field lines and raise the energy of the coronal plasma in the process. As suggested by Dere (1994) an important aspect of this problem involves the continuous recycling of the network flux by the emergence of intranetwork bipoles.
There are a number of observations consequences to this scenario that allow one to test its relevance to coronal heating. When magnetic elements collide with other elements of the opposite polarity, the result is magnetic reconnection or cancellation, first described by Martin (1977). Another consequence of the reconnection process is the generation of explosive events seen in the wings of transition region lines (Dere 1991). Experience with the HRTS indicates that high spatial resolution are needed to detect these events. The effect of this process in the coronal will be the transformation of the magnetic field line topologies and the work done by the photospheric motions drives the field further away from a potential field state to a more energetic state.
Observing Program:
The test of such a model will require the full observational capabilities of the Solar-B mission. First, a region of quiet sun containing a few supergranular cells will be selected and placed withing the field of view of the magnetograph. The duration of the observing program should be on the order of 2 days or on the order of the length of time in which the supergranular flux is recycled by the intranetwork flux emergence. A magnetogram cadence of about 1 minute is probably adequate and the magetograms should be obtained at high sensitivity to the weak intranetwork fields. The sequence of magnetograms should show in detail the continuous emergence and transport of the intranetwork fields and their reconnection or merging with the pre-existing network fields.
Images of the corona over an area roughly 3 times larger than the magnetogram target will be obtained with X-ray imager to show the field line connectivity of the corona within the target region and to other regions. Again a cadence of 1 minute might be reasonable but a higher cadence might be needed to see small scale transients.
In order to raster a meaningful area of the quiet sun at a good cadence, the EIS spectrograph raster will be confined to an area of the size of a single supergranular cell (35000 km or 50"). First, this area will be observed through the EIS slot to obtain a quick image of the coronal and transition regions structures. Then, the area will be rastered in 1" increments. The explosive events should be measured with a S/N ratio of 10. If they are at about 10% of the peak line intensity, then an integration time of 15s at a count rate of 65 c/s should achieve that necessary signal of about 1000 detected photons.
Similarly, the structures and velocities in the corresponding coronal structures can be investigated with the Fe XII lines. The sum of the Fe XII 193 and 195 lines will yield a detected photon rate of about 17/s. Detecting 250 photons in the line profile should result in a velocity resolution of about 1 km/s and can be achieved with an integration time of 15s. Transient flows may well also be a signature of coronal loops that develops new connectivities.
Integration times of 15s are needed in both the NRL1 and NRL2 bands and allow the supergranular cell to be rastered in 750s or 12.5 m. At the end of the raster, another observation through the EIS slot would also be useful.
Interspersed through the 2 day long observing period, it would also be useful to obtain information on densities in the coronal structures in order to test for the existence of subresolution heating events such as nanoflares. In order to obtain 10% statistics in the Fe XII 186 A line, it would be necessary to integrate for about 200s. Consequently, it would be necessary to perform a relatively coarse raster to maintain an observing cadence or to place these time consuming rasters at the beginning and end of the observing period.