These notes are a follow-up from the Solar-B EIS UK science meeting held at MSSL 2-3 June 1998 (minutes, science discussion). This is the output of actions 4 (Wavelength Ranges) and 5 (EIS Applications / studies).
This document will develop as the contributions are submitted and incorporated. A draft with representative content is expected to be ready by the beginning of July, and a final draft for discussion at the UK science meeting at IC on 28 July 1998.
Please submit further contributions to me, Matthew Whyndham, at mwt@mssl.ucl.ac.uk and Len Culhane at jlc@mssl.ucl.ac.uk.
24/6/98 - 9/7/98
Solar-B: Study of Potential Wavelength Ranges
---------------------------------------------
A report from RAL and Cambridge, June 1998.
The following wavelength ranges were identified at the Solar-B EIS
meeting on 3 June 1998:
Range 1 170 - 210 A
Range 2 240 - 290 A
Range 3 308 - 381 A
Range 4 517 - 633 A
Range 5 1334 - 1428 A
The current 'favoured' range is Range 2.
Ranges 3 and 4 are similar to the CDS NIS bands.
Ranges 1 and 5 have been suggested as potentially useful.
The following notes consider the merits and disadvantages of each range.
*************************************************************************
* A complete line list and intensities (CH, QS, AR, FL) for *
* each wavelength range, obtained from the CHIANTI package *
* are given on the RAL directory /disk2/mason/solarb_jun98/*.CH *
*************************************************************************
Range 1 - 170-210 A
-------------------
- OK for pre-launch calibration. The usual hollow cathode
lines which would be available in or near this range are:
169.5-175.6 Al III
204.3-208.9 Ne IV
Fairly good possibilities for in-flight calibration checks
(see Landi et al. 1998, A&A, submitted -- GIS calibration).
- Spectral range observed previously:
NRLA Flare spectrum - Dere, 1978, Ap. J., 221, 1062.
Flare (TR lines) - Widing, 1982, Ap. J.,
SERTS AR spectrum - Thomas and Neupert, 1994, Ap. J., 91, 461.
SERTS-95 AR spectrum - Brosius, Davila, Thomas, 1998, ApJ, submitted
CDS GIS1 spectrum - Harrison et al. Solar Phys. 170, 123, 1997.
- Contains a range of useful iron lines (+ a few weak O V and O VI)
Below is a list of intensities (calculated from CHIANTI)
for QS, NeTe = 1e14. The units are erg/cm^2/s/st/A.
Only I > 1e2 are given (plus O V, O VI).
LogT QS (I>1e2)
171.07 Fe IX 5.9 1e3
172.17 O V 5.4 2e1
172.94 O VI 5.5 3e1
173.08 O VI 5.5 6e1
174.52 Fe X 6.0 7e2
177.24 Fe X 6.0 4e2
180.41 Fe XI 6.1 6e2
183.94 O VI 5.5 2e1
184.12 O VI 5.5 3e1
184.54 Fe X 6.0 1e2
188.23 Fe XI 6.1 3e2
188.30 Fe XI 6.1 1e2
192.39 Fe XII 6.2 2e2
192.75 O V 5.4 5e0
192.80 O V 5.4 1e1
192.90 O V 5.4 3e1
193.51 Fe XII 6.2 4e2
195.12 Fe XII 6.2 6e2
202.04 Fe XIII 6.2 1e3
209.62 Fe XIII 6.2 2e2
- Te coverage:
The O V and O VI line are the only useful transition
region line and they are normally relatively weak.
Their intensity would increase by an order of magnitude
or more in QS and AR brightenings.
The main temperature range is LogT = 5.9 to 6.2,
with the addition of 5.4 to 5.5 if the O V, O VI
lines are observable.
This temperature coverage is not very good for coronal holes.
The extension of the wavelength range down to 165 A would permit
to include the Fe VIII lines (168.17, 168.55, 168.93) at LogT=5.7.
- AR/Flare lines:
Fe XXIV at 192.02 is expected to be strong.
Other flare/AR lines are Ca XVII (192.82), Fe XVII (204.65),
Fe XX (171.68, 173.43, 201.01), Fe XXI (178.77, 180.84, 187.89),
Fe XXII (184.18), Fe XXIII (173.32 and 180.04).
The extension of the wavelength range down to 165 A would include
the flare lines Ni XXVI (165.34) and Fe XXIII at 166.69 at
LogT=7.3 and 7.1 respectively.
- Ne diagnostics:
Examples of useful iron ions (plus others):
Log Ne range
Fe X 175.27/174.52 (8 - 10)
Fe XI 179.76/181.14 (8 - 10)
Fe XII 186.87/193.51 (8 - 10)
Fe XIII 202.04/203.79 (8 - 10)
Fe XXI 187.89/178.77 (10 - 12)
Thus, it seems that density measurements at coronal temperatures
are well covered. Some of the lines are weak.
There are some very, very weak O IV lines scattered over
this wavelength range which have the potential of being
density diagnostics (9 - 12).
Just outside the range, the ratio Fe XIV (211.32/219.13)
could be a useful line for density diagnostics (8 - 10).
Fe XXI provides possibilities of density
diagnostics in flares or nano-flares.
- Abundance studies:
There are some very, very weak NeV, NeVI lines, which are
unlikely to ever be seen and no Mg TR lines.
There are no suitable Ne/Mg ratios for abundance determination.
It might be possible to study the Ca/A and Ni/Fe ratios in AR
and flares (Ni XIV, XV, XVI are visible in AR's).
The Fe/Si abundance could be studied if the wavelength
range were extended to include Si IX (227.00).
However both Fe and Si are low FIP elements.
- Velocity studies:
The Fe IX line is isolated and very strong, so would be
useful for velocity studies at coronal temperatures.
Fe XXIV should be useful for flares.
- Other comments:
This wavelength range is deficient in TR lines.
There are some very, very weak NeV, NeVI lines, which are
unlikely to ever be seen and no Mg TR lines.
There are no suitable Ne/Mg ratios for abundance determination.
Coverage at coronal/AR/flare temperatures is excellent.
Range 2 - 240-290 A
-------------------
- Not bad for pre-launch calibration. The usual hollow cathode
lines which would be available in or near this range are:
267.1-267.7 Ne III
282.5-283.9 Ne III
Not very good possibilities for in-flight calibration, due to the
lack of density insensitive line pairs spanning the whole wavelength
range (see Landi et al. 1998, A&A, submitted -- GIS calibration).
- Contains a range of useful lines, mainly iron, silicon, sulfur,
and magnesium, see previous observations:
NRLA Flare spectrum - Dere, 1978, Ap. J., 221, 1062.
Flare (TR lines) - Widing, 1982, Ap. J.,
SERTS AR spectrum - Thomas and Neupert, 1994, Ap. J., 91, 461.
CDS GIS2 spectrum - Harrison et al. Solar Phys. 170, 123, 1997.
Below is a list of intensities (calculated from CHIANTI)
for QS, NeTe = 1e14. The units are erg/cm^2/s/st/A.
Only I > 1e1 are given (plus Mg V, Ne V, Fe XVI).
LogT QS (I>1e1)
240.70 Fe XIII 6.2 4e1
241.75 Fe IX 5.9 4e1
242.85 S XI 6.2 2e1
243.03 He II 4.7 2e1
243.79 Fe XV 6.3 5e1
244.92 Fe IX 5.9 2e1
246.00 Si VI 5.5 3e1
246.21 Fe XIII 6.2 9e1
248.46 O V 5.4 4e1
249.12 Si VI 5.5 2e1
249.18 Ni XVII 6.4 1e1
251.07 Fe XVI 6.3 1e1
251.96 Fe XIII 6.2 2e2
252.20 Fe XIV 6.2 3e1
253.79 Si X 6.1 2e1
256.20 Fe X 6.0 6e1
256.32 He II 4.7 7e1
256.38 Si X 6.1 7e1
256.42 Fe XIII 6.2 2e1
256.68 S XIII 6.4 1e2
257.15 S X 6.1 4e1
257.29 Fe X 6.0 1e1
257.38 Fe XIV 6.2 1e2
258.37 Si X 6.1 9e1
259.50 S X 6.1 8e1
261.05 Si X 6.1 7e1
262.98 Fe XVI 6.3 2e1
264.23 S X 6.1 1e2
264.78 Fe XIV 6.2 1e2
265.02 Fe XVI 6.3 2e0
270.00 Mg VI 5.6 1e1
270.40 Mg VI 5.6 2e1
270.51 Fe XIV 6.2 1e2
271.99 Si X 6.1 5e1
272.64 Si VII 5.8 1e1
274.21 Fe XIV 6.2 1e2
275.35 Si VII 5.7 6e1
276.58 Mg V 5.5 7e0
277.03 Mg VII/Si VIII 5.8 2e1
277.26 Si X 6.1 4e1
278.02 Ne V 5.5 1e0
278.40 Mg VII/Si VII 5.8 4e1
281.40 S XI 6.2 6e1
284.16 Fe XV 6.3 8e2
285.58 S XI 6.2 3e1
288.43 S XII 6.3 4e1
- Te coverage:
The temperature range is very healthy, LogT = 4.7 to 6.4, the lowest
temperatures being provided by He II 243 (4.7), O V 248 (5.4),
Mg V 276 (5.5), Mg VI 270 (5.7), Si VII 275 (5.7), Mg VII 278 (5.8)
and a bunch of lines in the 5.9-6.0 region. Much depends on the
intensity of these lines - the Si VII 275, Mg VII 278 and He II
243 lines should be bright enough, but we also need the O V or Mg V
lines. The CDS GIS2 range starts at 256 A allowing us to check on
the Mg V and Mg VI lines.
- AR/flare lines:
Fe XXIII (263.76) and Fe XXIV (255.10) are strong
in solar flares.
The other flare-like lines of Fe XVII (254.87), Fe XX (232.89),
Fe XXI (242.07, 246.99, 270.57), Fe XXII (247.16, 253.16),
Ni XXIV (264.83) should also be observable. There are some
weak lines from Ca XVII (244.06).
At slightly longer wavelengths, Ni XVIII (291.98)
would be useful.
At shorter wavelengths the strong Ni XXVI (233.76) and Ni XXV (238.82)
lines would be covered, also the weaker Ni XVIII (233.76, 236.34)
and Ca XVII (232.83, 238.31) lines.
- Ne diagnostic lines:
possibilities for Ne diagnostics:
Log Ne range
Mg VI 268.99/270.39 (Ne < 8.5)
Mg VII 280.74/278.40 (7 - 10)
Si X 261.06/258.37 (8 - 10)
S XI 246.89/281.44 (7 - 10)
Fe XIV 264.80/274.20 (8 - 10)
Fe XXII 235.17/247.19 (Ne > 11.5)
If the wavelength range were increased to 300 A, the following
density diagnostics would also be available:
Si IX 258.10/any 290,292,296
Fe XII 291.05/258.43
S XII 288.40/299.52
The Ne diagnostics cover several coronal lines and
one transition region pair.
- Abundance studies:
The is a good range of elements for abundance studies
- Velocity studies:
Fe XV 284 could be useful for velocity studies.
- Other comments:
The lines in this wavelength range are weaker than 170-210A.
There is a good temperature coverage, including TR and
flare lines, with some density diagnostic capability.
Good possibilities of element abundance studies.
Range 3 - 308-381 A
--------------------
- The hollow cathode calibration lines from Ne III allow a measure
at either end of the range, but there is nothing in the
middle. This could be covered by invariant ratios. So, the
calibration should be OK. The usual hollow cathode
lines which would be available in or near this range are:
308.6 Ne III
313.1-313.9 Ne III
379.3 Ne III
Good possibilities for in-flight intensity calibration checks.
(see Landi et al. Solar Physics 1997, 175, 553)
- The line list is given in Table 4.3 of the CDS Blue Book (Version
6, August 1995). It is the NIS1 range of CDS. It contains a
range of iron, magnesium and silicon lines. See also
R.A. Harrison et al. Solar Phys. 170, 123, 1997.
SERTS AR spectrum - Thomas and Neupert, 1994, Ap. J., 91, 461.
- Te range:
The temperature range of the brighter lines in this band, i.e.
in table III of Harrison et al. (1997) is LogT = 5.9 to 6.3.
There are 'cooler' lines such as Mg V (353.1 A), Mg VI (349.2 A),
Mg VII (367.7 A), Ne V (359.3 A) and O III (373-374 A)
but they are weak, they are seen in transition region brightenings.
- AR/Flare line:
The flare-like lines of Fe XXI 335.9 A and Fe XXII 349.3 A lie
in the range. The Ni XVIII 320.6 is visible in active regions.
- Ne diagnostics:
Density diagnostics:
Log Ne range
Mg VIII 315.0/317.0 (7 - 9)
Si IX 349.9/345.1 (7 - 9)
Si X 356.0/347.4 (8 - 10)
Fe XII 338.3/364.5 (9 - 12)
Fe XIII 359.7/348.2 (8 - 11)
Fe XIV 353.8/334.2 (9 - 11)
Fe XV 321.8/327.0 (9 - 11)
Abundance Studies:
Reasonable range of elements.
Velocity Studies:
No obvious candidates, several possibilities.
Further comments:
The power of this range lies in the number of density
diagnostics and the detailed temperature coverage of the
corona (particularly 5.8 < log T < 6.5). The main disadvantage
is the lack of any strong transition region lines. For CDS
this problem was solved by simultaneously observing the 517 -
633 A band (Range 4 - see below), which contains many strong
transition region lines.
Range 4 - 517-633 A
--------------------
- The hollow cathode calibration lines for this band are:
537.00 He I
584.3 He I
These two are well separated so it is workable.
Not very good possibilities of in-flight intensity calibration
checks (see Landi et al. Solar Physics 1997, 175, 553)
- The line list is given in Table 4.4 of the CDS Blue Book. Again, it
is the NIS2 range of CDS. See also the Solar Physics paper.
- Te coverage:
The temperature of the brighter lines of this range (Table
II of Harrison et al. 1997) is logT = 4.3 to 6.3, with a
good spread (4.3, 4.5, 4.7, 5.0, 5.2, 5.4, 5.5, 5.6, 5.8, 6.0,
6.2, 6.3). Indeed, this is a superb spread of temperatures for
a set of bright lines in one band!
- AR/Flare lines:
The flare-like lines of Fe XIX 592.1 and Fe XXI 587.9 become
strong at high temperatures. Two lines of Fe XX are also
expected but have uncertain theoretical wavelengths.
- Ne diagnostics:
Density diagnostics:
Log Ne range
O IV 625.9/554.5 (9 - 12)
If second order lines have significant intensity in this
wavelength range, then it may be possible to use the following
ratios:
Log Ne range
S XII 299.5/288.4 (9 - 11)
Fe XIV 264.8/274.2 (9 - 11)
- Velocity Studies:
There are some good, isolated lines for velocity studies,
O V (629.7) being an obvious example.
- Further comments:
This region has excellent temperature coverage, but very few
density diagnostics. It is a useful spectral region for
studying the dynamics of the corona and transition region.
Range 5 - 1334-1428 A
---------------------
- Intensity calibration??
It is intended to cover 1334-1428 A (1st order), 667-714 (2nd order)
There is a paper by Feldman et al (1997), Ap. J. ???
This range was covered by SUMER.
Below is a list of intensities (calculated from CHIANTI)
for QS, NeTe = 1e14. The units are erg/cm^2/s/st/A.
Only I > 1e0 are given.
First Order:
1334.52 C II 4.3 1e3
1335.66 C II 4.3 2e2
1335.71 C II 4.3 2e3
1343.51 O IV 5.2 1e0
1349.38 Fe XII 6.2 1e0
1371.29 O V 5.4 7e0
1393.76 Si IV 4.8 3e2
1397.23 O IV 5.2 1e1
1399.78 O IV 5.2 2e1
1401.16 O IV 5.2 1e2
1402.77 Si IV 4.8 1e2
1404.80 S IV 5.0 2e0
1404.81 O IV 5.2 9e1
1406.02 S IV 5.0 1e1
1407.38 O IV 5.2 2e1
1412.84 Fe II 4.2 1e0
1416.90 S IV 5.0 6e0
1417.23 Si III 4.5 1e0
1423.83 S IV 5.0 1e0
1424.72 Fe II 4.2 1e0
Second Order: (* ID by Feldman et al, not yet in CHIANTI)
667.735 Fe XI 6.1 1e0
671.015 N II 4.4 2e0
671.385 N II 4.4 6e0
671.410 N II 4.4 1e0
671.629 N II 4.4 1e0
671.772 N II 4.4 1e0
672.000 N II 4.4 1e0
681.72 *Na IX
684.998 N III 4.9 2e1
685.515 N III 4.9 5e1
685.817 N III 4.9 1e2
686.336 N III 4.9 2e1
690.519 C III 4.8 1e1
691.193 N III 4.9 2e0
691.396 N III 4.9 1e0
692.731 Si IX 6.0 1e0
694.14 *Na IX
696.622 S V 5.2 3e0
700.245 Ar VIII 5.6 5e0
702.332 O III 5.0 8e1
702.821 O III 5.0 7e1
702.891 O III 5.0 6e1
702.897 O III 5.0 9e1
703.848 O III 5.0 9e1
703.854 O III 5.0 3e2
706.060 Mg IX 6.0 4e0
712.671 S VI 5.3 1e0
713.812 Ar VIII 5.6 2e0
- Te coverage:
Good low temperatiure TR lines (4.2 - 6.0), but very poor coverage
of coronal temperatures. Fe XII is the only coronal line
(1st order) and this can only usually be observed on the limb.
There are some other coronal lines (around 10^6K) in 2nd order.
Feldman identified these in SUMER spectra.
- Flare lines:
The only flare line which is observable in this wavelength range is
Fe XXI (1354.1).
- Ne diagnostics:
The O IV lines around 1400 provide opportunities for Ne
deterination, however S IV and second order lines (O III)
can be problematic.
- Velocity studies:
One of the major advantages of going to longer wavelengths
is the opportunity to measure line profiles. However, care
must be taken with weak lines and second order blends.
- Further comments:
Temperautre coverage above 10^6K is poor.
This wavelength range is excellent for line profile studies.
It is also good for opacity studies.
CONCLUSION
----------
Range 1 - OK for calibration lines; good coronal diagnostics.
Poor temperature range; weak transition region lines.
Range 2 - OK for calibration; reasonable temperature range, with
some transition region coverage; Some coronal
and TR density capability.
PROBABLY THE BEST COMBINATION FOR EIS.
Range 3 - OK for calibration; superb for coronal temperatures and
coronal density diagnostics; not very good for the
transition region.
Range 4 - OK for calibration; wonderful temperature coverage;
very poor density capabilities.
WOULD PUT THIS AS THE SECOND BEST FOR EIS,
because of temperature coverage.
Range 5 - OK for calibration? Poor high temperature coverage,
but excellent for TR studies; spectral line profiles.
| Network Dynamics | Richard Harrison |
| Active Region Cool Loop Dynamics | Louise Harra-Murnion |
| Coronal Holes | Len Culhane |
| Particle Acceleration Issues | George Simnett |
| EMERGING FLUX -- CORONAL RESPONSE | Peter Cargill |
| Flares - Mass Motions in Coronal Lines | Peter Cargill |
| Flares - Reconnection inflow/outflow | Peter Cargill/Saku Tsuneta |
| Flares - Plasma Dynamics/evaporation | Len Culhane |
| Flares - Non-thermal Line Broadening | Louise Harra-Murnion |
| Abundance Anomalies | Helen Mason |
| CMEs - Role of Reconnection in the Onset | Richard Harrison |
| Diffuse Corona - Streamer Dynamics | Len Culhane |
| Loop Heating | [Eric Priest] [Suggestion for Eric by JLC] |
Title: CORONAL MASS EJECTION ONSET STUDIES
Author: Richard Harrison (RAL)
Justification:
The CME onset eludes our detailed understanding;
we have many models but lack detailed observation of the CME source
region at the time of a CME onset. This is partly due to the fact that
coronagraphs actually occult the onsets of the CMEs they observe! In
addition, since coronagraphs (i.e. CME observations) are tuned to the
plane of the sky, the source regions are near the solar limb. Such
regions are very difficult to observe because of foreshortening and
occultation by the limb itself. Even if the source regions are observed
at the time of a CME onset, to obtain detailed diagnostic information
of the CME onset process one requires plasma diagnostic information
(density, temperature, velocity) on spatial and temporal scales of
about a few arcsec and less than a few minutes, over a region of at
least several arcminutes by several arcminutes. A good range of
temperatures must be observed because the eruption processes include
both chromspheric and coronal plasma. Supporting coronagraph (to
identify CME) and ground-based (to see any prominence/filament eruption)
would be essential. The bottom line is that a spectrometer in the EUV is
required. This is an activity being performed using the CDS instrument
on SOHO. Whilst much headway may be made using the CDS observations,
the fine-scale flow patterns witnessed in prominences and the desire to
image over large areas with cadances of minutes are beyond the
capability of CDS. Thus, we anticipate great rewards from a close
investigation of CME onsets using EIS and the Solar-B payload.
Study Details:
Raster Area: 4x4 arcmin (* as large an area as is practical)
Raster Step: 2 arcsec
Raster Locations: 120
Exposure Time: 1 sec. (*Assuming a CDS count of 10 per sec
for the weakest line and an EIS
sensitivity 10x CDS and a desire
for 10% counting stats).
Duration of Raster: 1x120 = 2 min.
Number of Rasters: Open, but ideally should monitor region
for many hours each day as it aproaches limb.
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. He II 243.03 (256 is
blended), Mg VII 278.41, 280.74, Si VII 275.38,
Fe IX 244.92, Fe XIII 251.94, Fe XIV 264.78,
274.20, Fe XV 284.16, Fe XVI 262.98 (10 lines).
Bins Across Line: 25 (* To cover about +/- 250 km/s).
Telemetry/Compression: 10 lines x 25 bins x 120 bins x 12 bits(?)
= 360,000 bits per exposure. At 64 kb/s would take
5.63s. Require compression factor of about 5.
(*Assumes we take the slit only,
not the wider parts of the dumbell).
Solar Feature Tracking: Not required (near to limb)
Supporting Observations:
The EUV observations cannot identify a CME so we
require coronagraph support. In addition, the events
are undoubtedlt magnetic in nature so detailed magnetic
mapping would be ideal. Also, the EUV spectroscopy
can only monitor a relatively small region; a larger
area coronal mapper/imager wouldbe useful. So the
supporting devices would be:
- Coronagraph (to identify CME!)
- Magnetic Mapping of source region (Solar-B)
- Context Mapping of Coronal Structures (Solar-B)
What is New?
The temporal capability of EIS combined with
the ability to map flows of the EUV plasmas
provides a new approach to this important problem.
Notes:
------
- The basic driver is a large area, quick raster in a number of
emission lines over a good range of temperatures with some density
capability.
- The line selection contains a transition region and coronal density
diagnostic, but the weaker Mg VII line is rather weak. However,
summing successive rasters to obtain reasonable counts in the
weakest lines may be a useful approach.
- The observation assumes that we can select data from the slit
only and not the slots.
===========================================================================
Title: DYNAMIC EVENTS IN THE NETWORK
Author: Richard Harrison (RAL)
Justification:
Coronal heating and solar wind acceleration are ultimately powered
by the kinetic energy of the convection layers. Magnetic fields
carried in the convection cells 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 about 3000 distributed over the disc at eny time. No
significant velocities are associated with these events. However,
the explosive events are UV velocity events seen in the network,
maybe 30,000 on the Sun at any time, with speeds of up to 150 km/s.
The two classes of network transient events maybe related.
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. The superior EUV spectral resolution
and temporal resolution (afforded by the better sensitivity) would
allow EIS to take sigificant steps beyond the initial studies
made using CDS. The basic requirement is temporal resolutions
of under a minute over rastered areas of several tens of arcsec
(larger than a cell), with good temperature coverage.
Study Details:
Raster Area: 1x1 arcmin
Raster Step: 2 arcsec
Raster Locations: 30
Exposure Time: 1 sec. (*Assuming a CDS count of 10 per sec
for the weakest line and an EIS
sensitivity 10x CDS and a desire
for 10% counting stats).
Duration of Raster: 1x30 = 30 sec.
Number of Rasters: Minimum 500 (= 250 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. He II 243.03 (256 is
blended), Mg VII 278.41, 280.74,
Fe IX 244.92, Fe XIII 251.94, Fe XIV 264.78,
274.20, Fe XVI 262.98 (8 lines).
Bins Across Line: 25 (* To cover about +/- 250 km/s).
Telemetry/Compression: 8 lines x 25 bins x 120 bins x 12 bits(?)
= 288,000 bits per exposure. At 64 kb/s would take
4.5s. Require compression factor of 4.5.
(*Assumes we take the slit only,
not the wider parts of the dumbell).
Solar Feature Tracking: YES (* Must be able to make steps of less
than pixel size to avoid jumpy movies).
Supporting Observations:
- Magnetic Mapping of source region (Solar-B)
- Context Mapping of Coronal Structures (Solar-B)
What is New?
The temporal capability of EIS combined with
the ability to map flows of the EUV plasmas
provides a new approach to this important problem.
Notes:
------
- The basic driver is a quick raster on quiet Sun in a number of
emission lines over a good range of temperatures with some density
capability.
- The line selection contains a transition region and coronal density
diagnostic, but the weaker Mg VII line is rather weak. However,
summing successive rasters to obtain reasonable counts in the
weakest lines may be a useful approach.
- The observation assumes that we can select data from the slit
only and not the slots.
Questions Relating to EIS Performance from these Studies
========================================================
- For these two Studies I would want to use the slit, not the slots.
Thus, rather than waste telemetry, I would select on the
pixels covered by the slit data. I assume that this is no
problem.
- I have assume a 'final' full instrument sensitivity something
like 10x CDS. CDS has the two reflections off the telescope,
the scan mirror and the grating, as well as the detector
efficiency to worry about. EIS has the off-axis mirror,
the grating and detector. I think I remember quotes of
about 5-20x better than CDS. The count rates and, thus,
exposure times will depend on this.
- I have assumed that solar feature tracking is an option, using
the mirror mechanism. It would be nice if such a system
had step sizes less than the pixel sizes to avoid
jumpy movies. For the CDS smaller rasters this is
particularly obvious. I guess pointing should be automatically
updated between rasters.
- The two studies need a compression of order 5 to avoid the telemetry
driving the raster frequency.
- I have assumed a 2 arcsec slit and 2 arcsec steps in the rasters.
- I have assumed a 21 mA pixel size in the wavelength range.
- Other notes are given in the two Studies.
From: lkhm@msslac.mssl.ucl.ac.uk
Subject: Cool loops
Date: Thu, 09 Jul 98 16:33:19 +0100
Cool Loop Dynamics
------------------
Scientific Justification:
Recent results from CDS on SoHO and TRACE have confirmed the existence of
a separate class of loop system to the more familiar hot coronal loops -
that of cool 'transition region' loop systems. The cool loops exist as high
as the hot coronal loops but behave very differently.The differences are
summarised as follows;
1) There are fewer cool loops than hot loops
2) The cool loops are more sharply defined
3) There is no 'diffuse' cool component
4) Cool loops are extremely dynamic
The heating mechanism involved in the cool loop is different to that in the
hot loops. CDS has identified the loop structures but to understand them
further we must have higher time resolution (the limit on CDS imaging
for this type of observation was approx 10-15 mins), and we must be able
to observe the flows in the photosphere using OT, and then relating
that to what is happening in the cool loops. Measurements of line broadening
have been used to try and understand the heating mechanism in the corona.
To pinpoint the heating mechanism we need to understand the non-thermal
line broadening. There are some hints that the higher non-thermal velocities
in the transition region are merely due to multi-directional flows.
^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^
Line Selection: He II, Mg VII, Si VII, Fe XIII, Fe XVI
Pointing: Active region (once at the limb and once in disk center)
Number of exposures: Run for approximately 2 hrs in each location
The time duration of each observation should be minimised.
The minimum observing region should be 2 X 2 arcmins. The optical telescope
is necessary to observe the motions in the photosphere below the hot
and cool loops to search for different types of motions which produce
the dramtically different loop systems.
^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^
What's new?
Fast temporal resolution alongwith good spectral resolution for
the line broadening measurements.
Title: EMERGING FLUX -- CORONAL RESPONSE
Author: Peter Cargill (ICSTM)
Justification:
While the role that emerging photospheric flux plays in driving
coronal processes is widely accepted, recent observations
from the HAO/ASP, Hawaii magnetographs and SOHO/MDI have
in fact revolutionised this field. The ground based instruments
has indicated that flux can emerge in an already-twisted state,
so that free energy is already available to power dynamic
phenomena while MDI has shown the remarkable fact that flux
is recycled roughly every 72 hours. Given this, flux emergence
is likely to be a prime target for the Solar-B optical telescope.
Given the pivotal role flux emergence plays, it is essential
that its effect on the upper solar atmosphere be measured in
conjunction with its emergence. In particular, it is essential
to pin down the role that magnetic reconnection plays in the
interaction of new flux with pre-existing coronal field.
The best characterisation of magnetic reconnection are associated
plasma motions and these are best measured by obtaining spectra
of a few carefully chosen lines. The temperature choice is critical
for this observing sequence, since we expect the temperature
attained in reconnection to depend on the magnetic field strengths,
where the reconnection takes place, and how fast the new flux is
emerging. It is probably preferable to have better coverage at
transition region temperatures, but some trial and error will
be required to obtain the correct lines. Thus this run will
probably need to be repeated. It is anticipated that such a
sequence would have the optical telescope picking on a likely
region for flux emergence, and staring at it for as long as is
practical. We probably want a smaller FOV than OPT to get time
resolution so will need advice on where to point.
Probably want slot for context.
Strategy:
6 - 8 lines covering TZ and corona. He II, OV (if strong
enough), some of coronal Fe lines. Maybe a flare line in
case there is strong heating. Raster area should be only
part of FOV -- 1x1 arcmin say. With 1 sec exposure times
(see RAH's entries for arguments for this) can get 30
second raster duration, about same as resolution of
optical telecsope. Velocity coverage -- +/- 300 k/s
is probably a good start, though this can be increased
if required. Need slot for context. Again, similar
resolution to OPT will do as a start.
Coordination:
This sequence requires very close coordination with OPT.
XRT probably less necessary, but no harm in having it.
What's new:
The ability to observe flows in the transition
region and corona in response to well-resolved
observations of newly emerging flux
Title: FLARES -- MASS MOTIONS IN CORONAL LINES
Author: Peter Cargill (ICSTM)
Justification:
Solar flare energy release is believed to involve coronal
energy release and subsequent energy transport to the
photosphere. The photospheric response, a massive expansion
of heated plasma into the corona, is generally referred
to as chromospheric evaporation. The phenomenon was first
detected by SMM and P78-1, but with no spatial resolution.
The BCS instrument on Yohkoh has observed numerous cases of
blue-shifted plasma, but of course did not have any
spatial resolution. Solar-B provides the first opportunity
to infer the location of these mass motions. In addition
to blue shifts, line broadening has also been observed
in coronal lines during flares. It has been suggested
that this is turbulence associated with coronal reconnection
and as such could be responsible for coronal particle
acceleration. An observing sequence must scan a wide
area and do so quickly. The problem is that flares
tend to occur in a variety of places. The best bet
is probably to make this part of a sequence where SOT
studies a likely looking active region. The velocities
inferred in these phenomena are quite large, so a
large number of bins should be used. Unfortunately I
don't see a way around a large raster area since flare
loops can be quite long and we don't want to miss the
action. So I suggest same FOV as SOT. Can we get exposure
time < 1 sec? Are the lines strong?
Strategy:
Fe Flare lines (XXI and XXIII) are esential. Cooler coronal
line and TZ line to see depth of heating in upper chromosphere.
Density diagnostics would be nice, but appear to be restricted
to cooler plasma, so probably not worth it. Need many bins
in each line since we probably need to worry about velocities
up to say 800 k/s. FOV is a problem. We probably want maximum
in both slit and slot to be sure to catch loop footpoints. But
this is a problem for the time resolution and needs some thought.
Coordination:
Coordination with OPT and XRT essential.
What's new:
First spatial determination of evaporation and turbulence in flares.
Title: FLARES -- CORONAL RECONNECTION
Author: Peter Cargill (ICSTM) and Saku Tsuneta
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.
Strategy:
Reconnection outflows are expected to be hot, thus require the
Fe XXI and XXIII lines. There is no harm in including a couple
of cooler Fe lines, but TZ lines would not appear to be useful.
For spectral resolution, we would need to measure up to
1000 k/s, so perhaps 100 bins are needed. For the inflows, a
range of coronal lines (Fe IX - XVI) would work. Density
diagnostics would be nice and this could be done with some of the
Fe lines. The velocties expected are smaller than in the outflow
region, so +/- 250 k/s would work. The FOV is a problem,
since we want a large one in order not to miss the flare,
but a high time resolution in order not to miss the peak
of the reconnection. For context, we could probably use
XRT rather than the slot.
Coordination:
Disk campaign for reconnection outflows would be in
collaboration with OTP and XRT
Limb campaign would involve only XRT, with OPT giving
advice on likely-looking regions for a flare.
What's new:
The first accurate observations of mass flows associated
with reconnection in flares. First estimate of reconnection
rates in flares.
From: lkhm@msslac.mssl.ucl.ac.uk
Subject: Line Broadening
Date: Thu, 09 Jul 98 16:35:07 +0100
Flares - Non-thermal Line Broadening
-------------------------------------
Scientific Justification:
Using the data from the Yohkoh spacecraft we have been able to
increase our understanding of the various mechanisms involved in solar
flares. The time behaviour of the hard X-ray bursts (which is traditionally
viewed as the flare start) and the turbulence (as measured from excess
spectral line broadening) was analysed. Interestingly, it was
found that the peak of the turbulence occurred well before the peak of the hard
X-rays (Alexander et al, 1998).
Since turbulence occurs before the start of the flare
this requires a major reassessment of what is actually triggering
solar flares. Also, small flares occur more frequently than their larger
counter-parts. We found that the turbulence of these small flares
was surprisingly as large as that previously found in major flares
(Harra-Murnion et al., 1997). This suggests that the flare trigger is the
same in flares irrespective of their magnitude. The BCS onboard Yohkoh
has provided us with many clues to the understanding of the solar flare
trigger. However the major drawback is that it is a full Sun instrument
with no spatial resolution. Efforts have been made to search for the
region of highest turbulence by using the limb to occult the footpoints of
the flaring loops (Khan et al, 1995, Mariska et al, 1996). The results
were not conclusive.
^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^
Line selection: concentrate on the higher temperature lines initially - then
try a do a similar study for the cooler lines as well.
Pointing: Flaring active region
Number of exposures: as many as necessary to follow the preflare, impulsive
and gradual phase of the flare
As with Peter's mass motion study the information from the optical telescope
is crucial to follow the photospheric motions. The OT could also be used
to determine a suitable region to observe. The XRT should be used especially
if there is an overlap in the one of the lines used.
Another possibility is alternating between the slit and the slot at the same
location (a nodding motion) so that spatial and spectral info could be
obtained quickly and hence reduce the need for the large rastering???
^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^^
What's new?
The main issue is that this will provide the first opportunity to observe flares
with spatially resolved spectral information - hence the possibility to
resolve the location of energy release.
STUDYING ABUNDANCE ANOMALIES WITH EIS
The exact specifications of the EIS have not been finalised at the
time of writing (29-Jun-98) and so this document discusses necessary
conditions for investigations of the FIP effect.
The FIP Effect
Spectroscopic observations of the solar atmosphere, together with in
situ measurements of the solar wind have revealed that element
abundances often deviate from their values in the photosphere. A
common feature of such abundance anomalies is that elements with a
high first ionisation potential (FIP) are underabundant relative to
those with a low FIP, although the magnitude of the discrepancy varies
from feature-to-feature on the Sun. A useful summary of spectroscopic
studies of the FIP effect is presented in Feldman (1992).
Instrument Requirements
The key requirement, in terms of studying abundance anomalies, is that
the wavelength range that will be used for EIS contains emission lines
of both low and high FIP elements that have some overlap in
temperature.
A secondary requirement is that a range of consecutive ions of low and
high FIP ions are observed, to help minimise errors in the atomic data.
As an example, the Normal Incidence Spectrometer (NIS) of CDS/SOHO
observed lines of Ne IV - VII and Mg V - X (where neon is a high FIP
element and magnesium a low FIP element), for which there is
significant overlap in temperature around logT = 5.4 - 5.7. See, for
example, Fig.5 of Young & Mason (1998).
One difficulty in studying element abundances is that low FIP ions
give rise to many strong lines at temperatures logT > 5.7, but few
below this temperature, whereas the opposite is true for high FIP
ions. Two high FIP elements that do give rise to EUV lines at coronal
temperatures are argon and sulphur, but the lines are often rather
weak and so high instrument sensitivity and spectral resolution are
required to measure these lines accurately.
References
Feldman U. (1992). Physica Scripta 46, 202.
Young P.R. & Mason H.E. (1998). Proceedings of the ISSI workshop
`Solar Composition and its Evolution - from Core to Corona'.
If you have comments or suggestions, email me at mwt@mssl.ucl.ac.uk