Notes about Hinode/EIS
Alessandro Gardini
Institute of Theoretical Astrophysics – University of Oslo
July 2009
The Hinode mission
Hinode (“Sunrise”), formerly known as Solar-B, is a Japan ISAS (Institute of Space and
Astronautical Science) solar mission with national as well as US and UK collaboration. It is the
follow-up mission to the very successful Japan/UK/US Yohkoh1 mission. Using a combination of
optical, EUV and X-ray instrumentation Hinode will study the interaction between the Sun’s
magnetic field and its corona to increase our understanding of the causes of solar variability.
Hinode was launched successfully on September 22, 2006 by an ISAS M-V rocket from the
Uchinoura Space Center in Japan. Initial orbit was perigee height 280 km and apogee height 686
km, inclination 98.3 degrees. Then the satellite manouvered to a quasi-sun-syncronous orbit over
the day/night terminator, which allows near-continuous observation of the Sun for nine month a
year and short nights in the remaining period. The orbital period is about 90 min. On October 28,
the probe’s instruments captured their first images. The nominal mission lifetime is 3 years.
ISAS is part of the Japanese Aerospace Exploration Agency (JAXA). NAOJ acts as domestic
partner, while USA and UK collaborations are performed by NASA and STFC/PPARC
respectively. Norway and ESA also participate to the project.
Science goals of Hinode
To determine the mechanisms responsible for heating the corona in active regions and the quiet Sun
One of the remaining unresolved questions in solar physics is why the temperature of the Sun
atmosphere rises from 5780 K at the photosphere to millions K in the corona. Clearly the second
law of the thermodynamics tells us that the temperature should fall as you move away from a heat
source (at the core of the Sun).
The fact that this happens only until we reach the photosphere suggests that an additional nonthermal heat source is needed to explain and maintain the high temperatures seen in the corona.
Several mechanisms have been suggested and we are confident that the energy must come from the
magnetic field. However, there is as yet no consensus on the details.
Hinode will be able to determine the physical mechanisms responsible for coronal heating using a
combination of spectroscopic and magnetic information to determine how the magnetic energy is
converted in the EUV and X-ray radiation we see in the corona.
To determine the mechanisms responsible for transient phenomena, such as flares and coronal mass
ejections
Flares and coronal mass ejections (CMEs) are the most energetic manifestations of our Sun’s
activity and the most likely to directly affect our environment on Earth. CMEs, in particular when
headed in our direction, can cause major magnetic disturbances when they reach the Earth’s
magnetosphere. The causes of these explosions and expulsions of plasma are not yet fully
understood.
Hinode will provide measurements of magnetic fields, electric currents and velocity fields, which,
coupled with imaging of the corona, will reveal the trigger for both flares and CMEs. With an
understanding of what triggers these events our opportunities for reliable prediction increase
substantially.
1“
Sunbeam”
To investigate the processes responsible for energy transfer from the photosphere to the corona
The energy to maintain the high temperatures of the transition region and corona, and for flares,
CMEs and smaller scale activity observed in these parts of the atmosphere must come from the
magnetic field which originates below the Sun’s convection zone.
To determine how this energy is transferred from below the photosphere up to outer atmosphere we
need to be able to measure changes in the magnetic field simultaneously with the transition region
and coronal response.
With Hinode we will be able to make this connection and understand how the energy is transferred.
Hinode’s scientific instruments
Included in Hinode’s instrumentation is a 0.5 m optical telescope (SOT), an X-ray/EUV telescope
(XRT) and an EUV imaging spectrometer (EIS). The instruments will work together as an
observatory.
Solar Optical Telescope (SOT)
This is the largest optical solar telescope flown in space. Its aperture is 50 cm and angular
resolution achieved will be 0.25” (175 km on the Sun). The field of view of 400”x400”, and it
covers the wavelength range of 480-650 nm. SOT feeds the Focal Plane Package (FPP), which is
composed by two instruments: a spectrograph and a pair of filtergraphs that are used as a vector
magnetograph.
The FPP vector magnetograph will provide time series of photospheric vector magnetograms,
Doppler velocity and photospheric intensity.
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Magnetic lines: Fe I 525.0 nm, Fe I 630.2 nm; Continuum: 524.6 nm; Velocity: Fe I 532.4
nm.
Magnetic sensitivity: B (longitudinal): 1-5 G, B (transverse): 30-50 G.
Temporal resolution: 5 mins.
FOV: 164”x164” (for the FPP magnetograph)
The FPP spectrograph has a spectral resolution of 2.0 nm and provides detailed Stokes line profiles
of intensity and polarization.
X-ray Telescope (XRT)
XRT will provide coronal images at different temperatures, both full disk and partial fields of view.
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Temperature range: 6.1 < log T < 7.5 (1.25·106 < T < 3.16·107 K)
Temperature discrimination: log T = 0.2
Angular resolution: 2”
Temporal resolution: cadence: 2 s (reduced FOV), exposure time: min = 4 ms, max = 10 s
FOV: Full disk (> 30 arcmin)
XRT holds a flare trigger (see EIS below).
EUV Imaging Spectrometer (EIS)
The telescope is an off-axis paraboloid design, with focal length 1.9 m and mirror diameter 150
mm. EIS has a total length of 3 m. Multilayer toroidal gratings disperse the spectrum onto 2
detectors covering 40 Å each (see below).
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Wavelength range: 170-210 Å (17-21 nm) and 250-290 Å (the reflectivity of the
multilayers restrict the performance below 180 Å and above 204 Å)
Angular resolution: 2”. Four slit/slot position are available: 1” and 2” slit, 40” and 266” slot.
Spectral/Velocity resolution: 3 km/s for Doppler velocities, 20 km/s for line widths
Temporal resolution. In spectroscopic mode: < 1 s in dynamic events, ~ 10 s in active
regions. In imaging mode: monochromatic imaging of an active region (~ 4x4 arc min) in 3s
for dynamic events, 10 s othervise
FOV: 360”x512”
EIS provides monochromatic images of the transition region and corona at high cadence using a
slot. High spectral resolution images can be obtained by rastering with a slit.
EIS is able to response to XRT’s trigger. In addition it has internal triggers: a flare trigger and a
brightpoint trigger. It will also be able to have automatic exposure control for the slots.
EIS’s observations
The EIS instrument is described in details by Culhane et al. (2007, Solar Physics, 243, 19-61). It
was built by a consortium led by Mullard Space Science Laboratory (MSSL), and including
University of Birmingham, Rutherford Appleton Laboratory (RAL), NASA MSFC/GSFC, NRL
Hulbert Center for Space Research, University of Oslo, JAXA and NAOJ. In this paragraph, it is
summarized the information concerning observations and proposals.
EIS’s details
Hinode/EIS is a two-channel, normal-incidence EUV spectrometer. Its two channels cover the
ranges 170-210 Å and 250-290 Å, and are designed to observe solar coronal emission lines. The
scientific purpose of EIS is to study the response of the corona to changes observed by Hinode’s
SOT with its vector magnetograph and to compare observations with the full Sun-imaging XRT,
also on board.
EIS has a mirror which is tiltable in the Solar X direction, and is used to build up rastered spectral
images of the Sun in up to 25 spectral ranges. Additionally, EIS has both narrow (one- and twoarcsecond wide) slits, and wider (40- and 266- arcsecond) imaging slots, all with 512 arcseconds in
the cross-dispersion (Solar Y) direction. Under nominal conditions, the 40-arcsecond slot can be
used to make simultaneous, separated, quasi-monochromatic images in up to twelve strong
emission lines, covering a temperature range from He II (log T = 4.9) to Fe XXIV (log T = 7.2).
EIS should be able to make slit images of active regions in 10 s, of quiet Sun in between 30 and 60
s, and of flares in approximately one second.
Without moving the coarse mirror pointing, EIS is able to cover a spatial range in Solar X of 360”.
If the coarse mirror movement is engaged, EIS can off-point (in Solar X) from the centre of XRT’s
and SOT’s pointing by almost one solar radius (890”).
EUV emission lines in the EIS observed bands
The EUV imaging spectrometer (EIS) on board Hinode takes high resolution spectra in the two
wavelength bands 170-211 Å and 246-292 Å, referred to here as the short and long wavelength
bands (SW and LW, respectively). The two bands were chosen as they contain excellent diagnostic
of coronal and flaring plasma; the SW band in particular is the most rich in coronal plasma
diagnostic in the whole EUV region, and EIS is the first satellite-based instrument to observe it in
high resolution.
Constraints due to on board storage and telemetry mean that for EIS observations with a reasonable
cadence and spatial coverage it is not possible to send the complete spectra to Earth. An on board
selection of particular emission lines is thus performed, the reduced telemetry consisting of a set of
windows w pixels wide by h pixels high centered on each of the selected waveleghts. The most
important aspect of designing an EIS observation is to choose the emission lines for the science you
want to do.
The EIS is performing very well in terms of expected sensitivity and spectral resolution, and the
observed spectra are very similar to pre-flight predictions (Del Zanna & Mason, 2005) made using
the CHIANTI database and the expected instrument parameters. See Young et al. (2007), astroph/0706.1857 for figures of the spectra.
A key factor when judging the usefulness of an EIS emission line is the telescope effective area
(EA) at that wavelength, which determines the fraction of incident photons that arrive at the
detector. See Young above, for an overplot of the EA on the spectra. The use of two different
multilayer coatings on the EIS optical surfaces leads to EA curves that are peaked for the two
channels, with less sensitivity at the ends of each wavelength range. This is particularly true for the
SW channel which peaks at 196 Å, very close to the strong 195.12 line of Fe XII, making this the
strongest line observed by EIS in most conditions. The Fe IX 171.02 and Fe X 174.53 lines are
comparable in strenght to Fe XII 195.12 in moderately active solar conditions, but when observed
with EIS they are factors 1000 and 200 times weaker respectively. The EA curve for the LW band
is less peaked, giving more consistent instrument sensitivity across the band. The shape of these EA
curves is an important consideration when choosing emission lines for EIS observation studies: a
weak line near 195 Å can yield more instrument counts than a stronger line at the edges of the EA
curve.
Another important factor when considering emission lines is the degree of blending. Even with the
high spectral resolution of EIS, many lines are blended with other species, particularly in the SW
channel. Examples include two of the three EIS “core lines” -- He II 256.32 and Ca XVII 192.82
– which are included by default in all EIS observation studies. In many cases the contribution of
blending lines can be estimated quite accurately by making use of line ratios that are insensitive to
the plasma conditions.
EUV emission lines reccomended in EIS observation studies:
Ion
Cool lines (log Tmax < 6.0)
He II
OV
O VI
Mg V
Mg VI
Mg VII
Si VII
Fe VIII
Coronal lines (6.0 < log Tmax < 6.4)
Fe X
Fe XI
Si X
Fe XII
Fe XIII
Fe XIV
Fe XV
Hot lines (6.4 < log Tmax < 6.7)
S XIII
Fe XVI
Ni XVII
Fe XVII
Flare lines (log Tmax > 6.7)
Ca XVII
Fe XXIII
Fe XXIV
Wavelength/Å
log Tmax
Band
256.32b
192.90
248.46b
184.12
276.58
268.99
278.39b
280.75
275.35
194.66
4.7
5.4
5.4
5.5
5.5
5.7
5.8
5.8
5.8
5.8
LW
SW
LW
SW
LW
LW
LW
LW
LW
SW
184.54
257.26
188.23
182.16
258.37
261.04
195.12
196.64
196.54
202.04
203.82
274.20b
264.79
284.16
6.0
6.0
6.1
6.1
6.1
6.1
6.1
6.1
6.2
6.2
6.2
6.3
6.3
6.3
SW
LW
SW
SW
LW
LW
SW
SW
SW
SW
SW
LW
LW
LW
256.68b
262.98
249.18
254.87
6.4
6.4
6.5
6.6
LW
LW
LW
LW
192.82b
263.76
192.03b
255.11b
6.7
7.1
7.2
7.2
SW
LW
SW
LW
The three “core lines” that must be included in each EIS observational study are Ca XVII 192.82,
Fe XII and He II 256.32.
Procedure for submitting Hinode/EIS proposals
In order to standardise the proposal submission process for Hinode/EIS, the EIS team has built the
observation study definition software (often referred to as the EIS Planning Tool); it has been
made available through SolarSoft, since this is the most commonly used software package for solar
data analysis. Once the EIS branch of SolarSoft has been included in the local distribution, this is
run by adding it to the instrument list by 1) setssw eis, or 2) setting the environment variable
$SSW_INSTR to include the string “eis”. Then one can type “eis_mk_raster” in SolarSoft IDL to
start the EIS Planning Tool suite. Rasters can be either of scanning type (i.e., scanning a 2D area of
the Sun by moving the EIS mirror) or sit-and-stare observations (i.e., keeping the slit fixed on a area
of the Sun); they are then combined by “eis_mk_study” in one study. Once the study as been
defined, it must be exported to ASCII format so that it can be processed by the submission system
at MSSL. This is done within eis_my_study, once the study has been fully defined, by selecting
“Study, Export …” and saving the file with a .def extension to a convenient location.
All studies submitted must be accompanied by a science case, wich is evaluated by the seven
Science Schedule Coordinators (SSCs) for Hinode; the UK SSC is Prof. Culhane. In order to
evaluate the feasibility of proposed observations, predicted count rates per pixel should be included;
they can be easily derived from tables available in the MSSL web site.
In summary, interested parties have to submit a 500-word plain-text science case as the body of an
e-mail and attach the ASCII study-definition file exported from eis_mk_study. The e-mail has to be
sent to eis_obs@mssl.ucl.ac.uk .
An unsuccessful science case will be returned to the lead author with feedback. Successful
proposals will have the corresponding study tested on the EIS Engineering Model at MSSL, which
acts as an instrument simulator in order to check for any technical problems with running the
observations. If testing stage is passed successfully, the study will be entered into the official EIS
study database, and passed on the monthly Hinode meetings, where observations between the three
coordinated instruments are discussed and provisionally scheduled. The official EIS database will
also be mirrored through SolarSoft so that the general community can be aware of approved and
existing observation programmes.
Guidelines for designing EIS studies
A study is a collection of one or more rasters. When it comes to the timeline planning (i.e., when a
day observations are planned), it is studies that are inserted onto the timeline, not rasters. It is thus
the study definition that is exported to MSSL for testing, not the raster definition. A raster can either
be a raster scan or a sit-and-stare. Rasters are much more complicated to design than studies, since
studies are simply a combination of pre-defined rasters. All the information about spatial area, slit
numbers, emission lines, etc., is contained in the raster.
First of all, the type of observation has to be defined. Questions to be answered are:
 Is spectroscopic information necessary, or is imaging sufficient? This determines the use of
slots (40”, 266”) for images or of the narrow slits (1”, 2”) for spectroscopy.
 It is needed a scan raster or a sit-and-stare?
 How many lines are necessary? Just a few for basic temperature discriminations, or many
for detailed diagnostic?
The nominal EIS data rate is around 50 kb/s, and studies should always give a data rate close to this
value. Data decompression is available; the lossless DPCM compression compresses the data to
around 40% of its normal size, thus giving an effective data rate of about 120 kbit/s. JPEG
compression potentially gives a greater compression factor, but it is a lossy compression scheme,
and test are underways to assess their impact on spectra.
There is a huge number of possible rasters, and developing the best one requires some experience.
Because the ultimate limiting factor is the EIS data rate, four parameters have to be balanced in
order to keep the data rate close to the nominal one; they are:
 the wavelength window width
 the number of pixels along the slit
 the number of wavelength windows
 the exposure time
They are examinated in some detail herebelow.
For EIS, the wavelength windows can have different sizes that have to be set in block of 8 pixels
(i.e., 8, 16, 24, etc.). The EIS team recommends line widths of 32 pixels for most studies; if the
interest is in line broadening or high velocity events in active regions or flares, then 40 or 48 pixels
should be used. Given the high instrument sensitivity in the Fe XII 195.12 line, a window width of
40 or 48 pixels should always be used for this line. Greater than 300 km/s velocities have been seen
in Fe XII even in quiescent active regions. If high cadence is required, then it could be compulsory
to use only 24 pixels in order to reduce the data rate, but high velocity events can be missed out in
this case. Again, there are regions of the spectrum where a much larger window could be chosen in
order to pick up several useful lines in a single window. An example is the region from 256.3 to
258.3 Å where around 10-12 lines are found.
As well as allowing shorter exposure times, as described below, the 2” slit also allows to raster a
region more quickly, since half as many steps are necessary to do it. The downsides are a loose of
spatial resolution in the X direction, and that the spectral resolution is also degraded.
Regarding the number of windows, it has to be noted that even if a study is designed to observe a
loop, or bright point, or some other solar feature, it can actually catch an interesting flare or
transition brightening or similar phenomena. In this case, it would be significant including Fe
XXIV, O V, or some other ion. Thus, unless really big spatial coverage or very high cedence is
necessary, it is reccomended to have always a good selection of emission lines.
There are no hard-and-fast rules to set the exposure time due to the wide variation of the EIS
sensitivity with wavelength, but the following guidelines can be taken in account:
 For quiet sun, good line profiles for Fe XII 195.12 can be obtained in as little as 10 s
exposures. To get good profiles in a number of different lines requires 50-60 s.
 For active regions, 1 second exposures give useful profiles for Fe XII 195.12. For other
strong lines (e.g., the Fe XIII 203/202 density diagnostic) 5-10 s are fine. Some of the strong
lines are likely to saturate going beyond 30 s, but these exposure times are needed to pick up
weak lines such as transition region lines.
Finally, it can be useful remember that using the 2” slit, the count rate doubles, thus the exposure
time has to be halved accordingly.
Data transfer to/from Hinode
At first, Norway contrbuted to the Hinode project through the development and maintenance of
scientific software for the EIS instrument. Later on, ESA and JAXA came to an agreement that
most of the scientific data of Hinode would be transmitted to Earth thru SvalSat, the Norwegian
satellite station in the Svalbard Islands. Moreover, the European data-center for the Hinode data has
been set in Oslo, in the Institute of Theoretical Astrophysics (ITA).
SvalSat belongs to Kongsberg Satellite Services (KSAT), who manages the station, and to the
Norwegian Space Center (NSC). SvalSat is an elegant example about how Norway can exploit its
geographical location to strenghten its position among the nations involved in space. SvalSat,
thanks to its northerly location, is an excellent place to receive data from and send commands to
satellites in polar orbit. At SvalSat it is possible to download data from Hinode on every transit, 15
times a day. For comparison, the Japanese station enters in contact with Hinode on average three
times a day. Keeping in mind that the rate of the data transfer is a major bottleneck for the scientific
usefulness of this kind of satellites, it is easy to understand how much relevant is SvalSat for the
Hinode mission.
SvalSat.
Scientific data are downloaded to the SvalSat’s telemetry antenna at a rate of 4 Mb/s. Meanwhile
“housekeeping” data are transmitted by a reserved channel; these data contain all the information on
the satellite status, like temperature and voltage in electric circuits. A contact lasts on average 10
minutes, thus it is possible to download 2.4 Gb data on every passage.
Data are tansferred from SvalSat to KSAT in Tromsø by optical fiber. They are then sent to JAXA
in Japan using the ordinary internet connection. In Japan data are reduced for the first time, stored
in FITS format files and then sent to data-centers all over the world. The European data-center is set
in Oslo. It has been funded, togheter with the data management system as SvalSat, jointly by ESA
and NSC; this funding is the contribution of ESA to the Hinode mission.
Quick Look and Data Analysis software
The development of the Quick Look (QL) and Data Analysis (DA) software was the first
involvement of Norway in the Hinode mission and still represents a significant part of it. The QL
software is an instrument for the data reduction, while the DA software is devoted to the analysis of
the reduced data.
The data processing chain is segmented in the following steps:
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Planning
Uplink
Observation – Compression – On-board storage (aboard the spacecraft)
Downlink – Store – Decompression – QL
Reformatting – Store in FITS format – QL
Calibration
Distribution
More calibration – Visualization – Analysis
The QL and DA software has been developed by Viggo Hansteen and Øivind Wikstøl at ITA;
it is written using IDL in Object Oriented style, and emphasis has been set on flexibility. The
software can be run either as GUI or from command line.
QL is planned to perform a number of tasks. In the data handling it is able to read data in telemetry
and FITS format; to write data as FITS files, IDL savesets, or JPEG, MPG, GIF, PS, etc images;
and to find data both through catalog searches or browsing the headers. QL will also be able to
display the data information (wavelenght, spatial coordinates, time) as detector views,
spectroheliograms, “whisker” images, maps (intensity, velocity, width), line plots and movies.
Finally, QL is involved in the calibration and data reduction; during calibration, QL can perform the
flat fielding and the measure of the dark current, the removal of cosmic rays, destretching and
rotation of the images and the pointment and alignment with SOT and XRT. In the data reduction,
QL fits the line profile and measures its moments, and calibrates quantities (wavelenght and
radiometric) to physical units.
Requirements for running the EIS Quick Look and Data Analysis software
In order to run the EIS QL and DA software two things are needed:
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IDL version 6.0 (TBD)
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SolarSoft (SSW) with (minimum) the EIS, CDS and CHIANTI packages installed.
The software is designed to run under the most common operating systems. It has been tested on
UNIX, Mac OSX and Windows.
Installation of the EIS QL and DA software
1. Installing SolarSoft
The EIS QL and DA software is distributed through SolarSoft (SSW). SolarSoft is a set of
integrated software libraries (primarily in IDL) and databases that provide a common data analysis
environment for solar physicists. It has a wide range of general IDL routines, as well as instrument
specific ones for many instruments and space missions. Using SSW is highly recommended - both
because of all the useful stuff in there, and it ensures that everyone uses the same (and updated)
version of the EIS software package.
Many institutions have SSW installed and updated regularly. If it is uncertain whether or the
workplace has SSW installed, this needs to be checked with the local computer admin. It has also to
be made sure that the SSW package is updated regularly, preferentially every day.
A SSW installation needs as a minimum the EIS, CDS and CHIANTI packages installed to work
properly.
If there is need to install SSW autonomously (e.g. on a laptop), then it is proper to check the SSW
website and follow instructions from there. SSW has been successfully installed under Windows
XP and Mac OSX by the EIS programmers, and they are willing to share their experience to assist
users if necessary.
Once SSW is installed, it has to be updated regularly. This is very important otherwise users will
eventually end up running an outdated version of the software, which could even have bugs in it.
The details on how to do this are described on the SSW webpage, and this is dependent on which
operating system it is used.
2. SSW Setup on a local machine - UNIX (Mac OSX and Linux)
Assume SSW is installed and updated regularly. Now, SSW has to be set up to run properly. This
again depends on the platform SSW is installed on. In what follows, it will be described the
procedure for a laptop running Mac OSX with tcsh UNIX shell.
Running tcsh, there should be a file called .cshrc on the home directory. In this example the SSW
installation is under the ~/ssw directory. With this setup, the following commands are written in the
.cshrc file:
setenv SSW ~/ssw
setenv SSW_INSTR "CDS EIS CHIANTI"
source $SSW/gen/setup/setup.ssw
The first command sets up the SSW environment variable, and should point to the top diretory of
the SSW installation. The second line lists the instrument packages that are installed and want
access to, and the third command runs the SSW setup script.
Running a different UNIX shell, the commands may differ slightly and need to be put in another
file (e.g. in the .bashrc file for bourne shell).
The SSW setup procedure for Windows XP is not shown.
3. EIS specific setup - UNIX (Mac OSX and Linux)
If the previous steps have been performed correctly, there should be now a working set of EIS
software to perform QL, calibration and data analysis on EIS data. It remains to get acces to the
data, and there are several ways to do it. The EIS data distributed through SSW can be accessed for
example using the ‘eis_cat’ program. Here it is shown how to access data stored locally, and in the
following, it will be assumed that local EIS data are stored in a catalog named ~/solarb/eis/data/. To
access these data, it can be used a widget tool called ‘xfiles’. In order for this program to work
properly, it needs to know where to start looking for data. This information is set in some
environment variables that are usually defined by the SSW setup procedure. They can also be set in
the .cshrc file, for example as:
setenv EIS_INST "ITA-UiO"
setenv EIS_DATA ~/solarb/eis/data
setenv EIS_CCSDS_DATA $EIS_DATA/science/ccsds
setenv EIS_FITS_DATA $EIS_DATA/science/fits
setenv EIS_CAL_DATA $EIS_DATA/cal
This sets up the environment variable "EIS_DATA" and the default catalogs for telemetry
(CCSDS2), FITS and calibration data. Alternatively, information on the storage directories can be
introduced manually in the Xfiles window.
EIS QL Userguide
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Xfiles – The file selection dialog widget
In IDL (SSW IDL) type
SSWIDL> xfiles
to start the QL front widget called Xfiles. This is the user interface to select data files to view.
Xfiles searches for locally stored EIS data, starting from the default directories specified by the
$EIS_DATA environment variables briefly described above. The files selection dialog is the Xfiles
window. A screenshot of the Xfiles interface is displayed below.
The Xfiles window has in its upper left corner an Exit button, which will kill the window and
terminate the EIS QL session. Next is the Select data source button box, where the user can select if
the xiles program should search for telemetry (CCSDS), FITS or housekeeping files. The default is
FITS. The default directories where the program will go to search for files are defined in the
EIS_CCSDS_DATA, EIS_FITS_DATA and EIS_HK_DATA environment variables, defined as a
part of the setup procedure. If these variables are not set, the current directory is chosen, as shown
in the figure.
Below there is a search field based on time, where it is possible to search for rasters based on raster
start and end time. This field is followed by a filter field to show the directory content, and the next
box displays the directory where the file search is performed. All these fields may be edited by the
user. For example, it is possible to type in a different directory name and hit return to list data files
in that directory.
2
Consultative Committee for Space Data Systems
The Xfiles window: the EIS QL local file selection interface (on Linux).
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The files found are listed in the bottom box of the Xfiles window. One file will be selected
by clicking on it. Pressing the Confirm selection button at the bottom, a new window,
Xcontrol, will be opened. Xcontrol - The main QL control window
The Xcontrol window may be described as the QL main control window for the selected raster.
From this window the user is presented with some basic information about the raster, and the basic
displays are selections available from here.
The Xcontrol window: the EIS QL main control window for raster display.
It has two menu items in the top menu bar. Under the File menu, Close terminates the Xcontrol
session for this raster. Under the Options menu, there are two items:
 Display header: Launches a text window displaying the contents of the header.
 Compare EIS/Chianti spectra: This is used to compare an average EIS spectrum calculated
from the raster with a Chianti spectrum of the same wavelength region. An intermediate
window pops up where the user must select the proper emission measure for the Chianti
spectrum. The Chianti spectra have been pre-calculated for the EIS wavelength regions and
the most relevant emission measure regimes. The Chianti spectra are stored in proper files
and distributed as a part of the EIS QL package.
The first box in the left column lists the line windows contained in the raster, as well as line
window information when available (from the planning tool). The line window selection is relevant
and has to be performed for all display modes except the detector display, where full CCD images
are displayed, and the browser.
The second box in the left column is used to select the display mode.
 The detector displays full CCD images (see the Xdetector section below), i.e. count rates as
a function of wavelength (abscissa) and slit position (ordinate). The ordinate translates into
Solar Y position. Count rates are summed over all the exposures.
 The browser opens a general purpose tool.
 The spectroheliogram displays count rates for the selected line(s) as a function of
wavelength (abscissa) and slit position (ordinate), one image for each selected line and
exposure.
 The whisker mode launches one separate display window for each selected line. Here, the
line count rates are displayed as a function of wavelength (abscissa) and exposure number
(ordinate), which translates into time or raster position. If the raster is a sit-and-stare, the
ordinate axis translates into time, if it is a scanning raster, then it is the Solar X coordinate.
 The last display mode available from Xcontrol is the intensity map. It displays data from
one line window as a function of exposure number (abscissa) and slit position (ordinate). In
order to do this, an average over the line profile (wavelength) has been calculated. Again,
the abscissa dimesnion will be Solar X for a scanning raster and time for a sit-and-stare.
The next box contains buttons related to the calibration of the data, described in this document in
the paragraph related to the data analysis. Calibration Def allows defining the calibration options
and the dark current file to be used. Eis Prep calls the eis_prep routine, which actually performs the
task. When the options are not set, eis_prep assumes default values.
In the following box, the Moments button calls the routine to compute moments. Display header
shows the header of the data file.
In the last box, Close exits the window, while Save saves the current version of the calibrated data.
In the top right column, there is a text window (not editable) that presents key parameters
describing the raster, such as if it is a sit-and-stare or scanning raster, number of exposures, number
of slit positions etc.
Next, a figure shows full CCD images, like a reduced version of Xdetector.
This is followed by a row of buttons displaying information about the pointing coordinates and the
CCD temperatures.
At the bottom, a figure shows the pointed region over the Sun. It is sided by a set of labels which
state the calibration status of the data currently in use, as well if they are saved or not saved.

Xdetector - View the data on the detector
The Xdetector GUI is used for viewing EIS data as detector images. If the raster contains lines on
both detectors, two Xdetector windows, one for each CCD will pop up. The Xdetector interface will
look something like the image below, where two line windows each of 256 pixels width are
displayed on CCD A. This raster contains a test pattern from engineering tests at MSSL. The
abscissa in the image is wavelength (here shown as pixel numbers) and the ordinate axis is position
along the slit (again in pixels). Slit position could translate to Solar Y position for real data.
The Xdetector interface: it is used to view data as detector images.
The upper menu bar has to items:
File which has two submenus:
Save as: to save the CCD image as either jpg or ps file.
Close: to close the xdetector window.
Options which has three submenus:
Colour table: to manipulate the color table of the CCD images.
Create animation: if the raster contains more than one exposure, the data can be played as a
movie. This is done for one line at a time, so the function requires a line selection, which is done in
a line selection window that pops up.
Change wavelength scale: to change wavelength scale on the image from pixels to Ångstrøm or
vice versa.
The left column is used for the various display functionalities available for the CCD images.
The Window action box is used to focus on a subimage of the total CCD image.
If the Zoom button is selected, a new window with a zoomed in image of the selected region
appears on the screen. The zoomed region appears on the original Xdetector image as a box, which
is selected using the left mouse button in the image. The magification is presented in a separate
window, and again there is a menu in this window to save the magnified image as jpg or ps.
The zoom region is selected using the left mouse button and appears as a box in the Xdetector
image.
The Xzoom window displays the magnified image of the selected region.
The other two buttons in the Window action box also take effect once a region of the CCD image is
selected using the left mouse button. Average along wavelength has the effect of displaying a line
plot, showing averaged count rates over the selected region in wavelength, for the range of slit
positions selected. Similarly, Average along slit displays count rates as a function of wavelength
averaged over the slit area selected.
The slider labeled Exp # is used to select exposure number when the raster contains more than one
exposure. Slider counter starts at 0, thus when the slider value is 1, the CCD image displays the
second exposure of the raster.
The Plot pixel values drop menu is used to display actual count rates as line plots. The row plot
plots count rates (intensity) as a function of wavelength at a given raster position. The column plot
plots count rates as a function of slit position at a specific wavelength pixel.
The line plots are presented using a tool called Xlineplot. The image below shows an example,
where the row plot has been selected, i.e. count rates as a function of wavelength at a given slit
position. The top slider then may be used to change the slit position for the line plot. The plot can
be narrowed down in wavelength using the Start [x] and Stop [x] sliders. If the Automatic y-scale
button is de-selected, the y-scale of the plot may also be changed manually using two more sliders
that then become active. One may mark each data point with an asterisk using the proper button,
and play the line plot as a movie for all slit positions by clicking the Movie button.
Similarly, the wavelength could be changed, the slit range narrowed down etc. if the plot is a
column plot instead of a row plot.
The Xlineplot tool: in this case a row plot displays CCD count rates.
Again in the left panel of the Xdetector window, the Pixels and Ångstrøm alternative buttons allow
to change wavelength scale on the image from pixels to Ångstrøm or vice versa.
The Real size CCD display button is used to present the CCD image in real size, e.g. 2148 x 1024
pixels.
The Close button terminates the Xdetector window. Each window, if there are two, is closed
separately.
The middle panel of the Xdetector window displays the CCD image, while the right panel displays
the colorbar, to indicate the count rates and colour table used to display the image.

Xraster - Spectroheliograms
The spectroheliogram way of displaying the data is displayed below. It is launched by clicking the
corresponding button in the Xcontrol window, if one or more line windows have been selected from
the list.
The spectroheliogram display mode: in this case a 4 line window raster, with 3 exposures is
displayed. Each image contains count rates as a function of wavelength (abscissa) and slit position
(ordinate).
The spectroheliogram image shows line count rates as a function of wavelength (abscissa) and slit
position (ordinate). The top row shows N number of images for the N exposures in the raster for the
first selected line. Second row shows images for line number 2 etc. In the figure there are 4 lines
selected and 3 exposures in the raster.
The menus are the same as those described for the detector display.

Xwhisker – Whisker images
The whisker display is selected from the Xcontrol display mode menu. Whisker images shows line
countrates as a function of wavelength (abscissa) and exposure number - time or solar X (ordinate). In a sit-and-stare campaign the ordinate axis will be time, and for a scanning raster the
ordinate axis will be solar X position.
Xwhisker
The image then is for one specific slit position. Such images have proven in the past to be
particularly useful for sit-and-stare raster, where the time evolution of the line at one position on the
sun becomes evident. The image displays a whisker plot for a line consisting of 100 exposures.
There are no velocities in these artificial data; in real data, velocities would show up as horizontal
"whiskers" - therefore the naming of this display mode.
The menus available to this mode of display are similar to those described for the detector display.

Xmap - Intensity maps
The intensity map mode of displaying data averages the line over the line profile. An example of
artificial EIS data made by John Mariska from a TRACE3 image of an active region is displayed
below.
Xmap
If more than one line is selected in Xcontrol, several such windows will show up. For a raster
containing M slit positions and N exposures, the map will have N x M pixels. For a scanning raster,
the N exposures (abscissa) translates into solar X position, while for a sit-and-stare the
corresponding axis dimension will be time. The ordinate axis translates to solar Y.
The menus available to this mode of display are similar to those described for the detector display.
EIS Data Analysis Userguide
Here are covered the basics of getting started with the data analysis.
1. How to get the data.
3
“Transition Region And Coronal Explorer” satellite
There are two ways to get the data. Firstly, there are various open archives:
DARTS at ISAS: http://darts.isas.jaxa.jp/hinode/top.do ,
ESA’s archive at the UiO: http://sdc.uio.no/sdc/welcome ,
UK’s archive at MSSL: http://msslxr.mssl.ucl.ac.uk:8080/SolarB/Archive.jsp .
Each of these systems has different ways to search for data. The ESA one has thumbnails, the
MSSL system has a movie maker and thumbnails for both level 0 data (unprocessed) and level 2
data (line shifts and line widths).
Alternatively, it is possible to use the eis_cat routine within IDL. This allows you to search for the
data by remotely searching the EIS catalogue. The routine is referenced in
http://orpheus.nascom.nasa.gov/~zarro/idl/eis/eis.html
2. Types of data which are available.
There are 3 different types of data files with EIS:
Level 0: these are reformatted data which are unprocessed. These data are normally obtained upon
request.
Level 1: these are the calibrated data. To obtain the level 1 data, level 0 data have to be processed
by the routine eis_prep.
Level 2: these data provide a quick look first cut at the velocities. Since they are generated by an
automated process, it is not reccomended to use them for detailed scientific work, but merely to
determine quickly if there is anything interesting in the data set, such as large velocities. These data
can be obtained and looked at online. The level 2 data are produced only for 1” and 2” slit data.
3. How to read and extract data.
The easiest way of accessing data is through the IDL procedure called eis_getdata:
SSWIDL> eis_getdata, myfilename, wd, (iwin = iwin, /all, hdr = hdr, /ccsds)
The inputs parameters are:
myfilename: name of data file (in CCSDS or FITS format. Default is FITS - see CCSDS keyword
described below.)
The output parameters are:
wd: name of the data cube returned. It can be called anything by the user.
The keyword parameters (optional) are:
iwin: this keyword can be set to a series of line window numbers to return those line windows. E.g.
iwin = [0,3,5] returns line windows 0, 3 and 5 from the raster.
all: this keyword is set to return all line windows. If this is not set and iwin is not set, line window 0
is returned as default.
hdr: this keyword is set (e.g.hdr=hdr) to return the header information.
ccsds: this keyword must be set if the data file is a telemetry file (the default is FITS).
Example1: The data file myfilename is a FITS file containing a scanning raster with 200 raster
positions, 512 pixels along the slit, 25 line windows each of 32 pixels width. To return the line
windows 3, 12 and 24 and the FITS header, the instruction is:
SSWIDL> eis_getdata, myfilename, wd, iwin=[3,12,24], hdr = hdr.
The wd cube should now have dimensions: [32, 512, 200, 3] or [nlambda, nslit, nraster, nlines].
Example2: The data file myfilename is a CCSDS file containing a full CCD exposure: 512 pixels
along the slit and 4 line windows each of 1024 pixels width. To return all the line windows and also
the telemetry header, the instruction is:
SSWIDL> eis_getdata, myfilename, wd, /all, hdr = hdr, /ccsds
The wd cube should now have dimensions: [1024, 512, 1, 4] or [nlambda, nslit, nraster, nlines].
4. How to calibrate the data.
The eis_prep routine should cater for most of user calibration needs. It must be set up “setssw cds
eis” in order to run the Cosmic Ray removal. The routine eis_prep can correct for: dark current
subtraction, hot pixels, Cosmic Ray removal, and absolute calibration. The flat field correction has
not been implemented. The calibrated data are saved in a level 1 file, while bad/missing pixels are
marked and saved in the error file eis_er_yymmdd_hhmmss.fits file.
Example: to run eis_prep with default values and save the results, the command is
SSWIDL> eis_prep, myfilename, /default, /save
5. How to display the data.
The EIS instrument is very flexible with many different modes available – hence there are different
ways to look at the data, some of which are described below. The main thing to consider is which
slit/slot was used: 1”, 2”, 40” or 266”.
(a) Showing spectrum data
When using the 1” or 2”, slit rastering is often done. The dimension of the the raster scannng data is
data[lambda,x,y]. To have a quick look at a movie of the spectrum, the instructions are:
SSWIDL> window, 0, xs=xsize[iwin], ys=ysize[iwin]
SSWIDL> stepper, data
(b) Display slot picture
In the case of data taken in the slot mode (40” or 266”), to show the picture in the correct direction,
it is necessary to reverse the left-right (i.e. east-west) direction.
SSWIDL> window, 0, xs=xsize[iwin], ys=ysize[iwin]
SSWIDL> tvscl, rotate(data[*,*,0],5)
(c) Displaying raster data
For raster scanning data, integrating the spectrum along the wavelength direction gives the image of
a particular spectral band.
SSWIDL> intensity=rotate(total(data,1),1)
SSWIDL> window, 0, xs=nexp, ys=ysize[iwin]
SSWIDL> tvscl, intensity
6. How to calculate velocities, widths etc.
There are two ways of calculating intensity, velocity and width from the line profiles. It can either
be done from the QL control window Xcontrol, by clicking the Moments button, or it can be done
using a procedure called eis_getmoments. In the first case, results are displayed using the QL, in the
later, the user will get hands on access to the calculated quantities.
In both cases, the user needs to manually interact in order to define line and continuum pixels. In
the future this could be done automatically from e.g. line IDs, but so far the logic around this has
not been worked out.
The usage of eis_getmoments is decribed in what follows:
SSWIDL> eis_getmoments, file, int, vel, wid, iwin =iwin, all = all, gauss = gauss, moment =
moment, ccsds = ccsds, lamtype = lamtype, line_def = line_def
The input parameters are:
file: name of file (string) containing EIS data (and headers).
The output parameters are:
int: intensity (0th order moment). Normally int[x,y, nwin]
vel: velocity (1st order moment). Normally vel[x,y, nwin]
wid: width (2nd order moment). Normally wid[x,y, nwin]
The keyword parameters are:
iwin: the line number indexes (starting with 0, max 24.)
all: to process all line windows
gauss: to use a Gaussian fit to calculate int, vel, wid. This calls a CDS procedure called xcfit.
moment: (default) to use moments to calculate int, vel, wid
line_def: to use pre-stored pixel definitions for line and continuum. (Not yet implemented) These
are stored in text files, as pixels defining where on the CCD the spectral line and its continuum (to
be subtracted) are.
lamtype: can take two values: lamtype='pix' (default) for pixel wavelength scale or lamtype='wav'
for wavelength scale in Ångstrøm.
ccsds: this keyword must be set if the data file is a telemetry file (the default is FITS).
Example 1: To calculate int, vel wid, from line window 0 from the CCSDS file1= 'file1.dat':
SSWIDL> eis_getmoments, file1, int, vel, wid, iwin=0, /ccsds
Example 2: To get int, vel wid, from the line windows 2 and 3 from the file file2='file2.fits', using
wavelength scale and Gaussian fit:
SSWIDL> eis_getmoments, file2, int, vel, wid, iwin=[2,3], lamtype='wav', /gauss
7. How to remove Cosmic Rays.
There is a routine called eis_despike to do Cosmic Ray removal for one line window and one
exposure. So far there is no script available to process many line windows and several exposures
automatically. eis_despike is based on the SSW procedure called new_spike, which has been
extensively used and tested on CDS4 data.
SSWIDL> eis_despike, wd_in, wd_despike
The input parameters are:
wd_in: line window before CR removal
The output parameters are:
wd_despike: line window after CR removal
The keywords parameters are:
no_neighbours
no_fill
info
For a description of the keywords, see the documentation of new_spike in SSW.
8. How to identify lines in the spectrum.
It is possible to look at the spectrum at each pixel location either with the Quick Look software
Xfiles or using:
SSWIDL> eis_raster_browser,file
To identify the lines it is possible to use CHIANTI through
SSWIDL> xspecplot
Not all lines have yet been identified in the EIS wavebands, so there is most likely some missing
from CHIANTI.
4
“Coronal Diagnostic Spectrometer” onboard SOHO
EIS data files in FITS format
1. The content of the EIS FITS file header.
The list of EIS FITS keywords, including a description, is available at
http://www.astro.uio.no/solarB/software/eis_userguide/fits_kw_list.htm
The page is updated as 11 Sep 2006. An example of an EIS FITS header, is available at
http://www.astro.uio.no/solarB/software/eis_userguide/fitshdr.html
2. How to produce EIS FITS files.
Being Hinode in full operation, this is not something most users of EIS data should have to care
about. Level 0 FITS files are distributed to the users through the various data centers around the
world. Typical users may want to update their FITS files, on the other hand, e.g. after having
performed calibration of the data, as described above. Here it is shown how to produce a FITS file
from a telemetry file (CCSDS) containing data of one EIS raster. This can be done in several ways,
described below. Currently, methods 1 or 2 are recommended, as the GUI is rather buggy.
1) Working with IDL objects:
If has already read a telemetry file, and has the data object and header object defined, the user can
use the IDL procedure called eis_mkfits to produce a FITS file. See the program header for the
documentation on how to use it and what keywords it takes.
2) Using a procedure called eis_mkfits_script:
Typing eis_mkfits_script at the IDL command line will return how to use the script.
If no keywords are given, it will have no action.
If the "file" keyword is specified, as by
SSWIDL> eis_mkfits_script,file='~/solarb/eis/data/science/ccsds/md_data/run01/md_data.0007’
a FITS files will be produced and written to the current working directory.
If the “directory” keyword is specified, as by
SSWIDL>eis_mkfits_script, dir = '~/solarb/eis/data/science/ccsds/run059/'
one FITS file will be produced for each telemetry file listed in the specified directory.
Again, the FITS files will be put in the current working directory.
3) Using the Xmkfits GUI:
Typing xmkfits on the IDL command line starts up the Xmkfits GUI, displayed below.
This is a fairly simple GUI, designed to list telemetry files in a directory, and output FITS files in
another directory (it could of course be the same.) Please note that at this point this program is
under development. It WILL have bugs. Please report them.
The various buttons and text fields have the following actions:
o
The Exit button quits the Xmkfits program.
o
The Create FITS button starts processing the telemetry files listed in the telemetry
data catalog listed in the field below the button. For each telemetry file, a FITS file
will be produced and placed in the specified FITS output catalog.
o
The Telemetry Catalogue text field can be edited by the user. It tells the FITS
production software where to look for telemetry data. The default catalog will be the
one defined by the EIS_CCSDS_DATA environment variable. Editing this field, it
must be hit "return" to do a re-listing of files.
o
The Telemetry data file list lists all files in the directory specified in the above text
field. Note that at this point all files will be listed, not only telemetry files, if there
are others. Also subdirectories and files in them will be listed. To select the file to
make a FITS file from, it must be selected clicking on it and only one file can be
selected at a time. There is currently no "select all" feature.
o
The FITS output Dir/Name text field is where the FITS files will end up. Default
directory is defined by the EIS_FITS_DATA environment variable.
In the Messages field, a message will appear when a FITS file has been produced. It will tell the
name of the file, with full directory reference. A reference article on EIS
The Solar-B Mission and the Forefront of Solar Physics
ASP Conference Series, Vol. 325, 2004
T. Sakurai and T. Sekii, eds.
Solar-B EUV Imaging Spectrometer (EIS)
Hirohisa Hara
National Astronomical Observatory, Mitaka, Tokyo 181-8588, Japan
Abstract. The Extreme-ultraviolet Imaging Spectrometer (EIS) onboard Solar-B is designed and
being developed under international collaboration to execute spectroscopic observations of the
dynamical solar upper atmosphere in two extreme-ultraviolet wavelength ranges: 170-210 Å and
250-290 Å. These bands include strong emission lines that are formed at 104.7-107.3 K. A highcadence spectroscopic observation with EIS is suitable for the studies of coronal heating process
and it will reveal dynamical plasma motions in/around solar flares af all sizes. It is of particular
importance to specify the key physics of coronal heating process and solar dynamic events through
collaborative observations with SOT and XRT on the same spacecraft.
1. Introduction
The Extreme-ultraviolet Imaging Spectrometer (EIS) is one of the three major scientific instruments
of the Solar-B. The EIS instrument was designed was designed and is being developed by
international collaboration of the United Kingdom, the United States, Norway and Japan (Culhane
et al. 2000).
Yohkoh, SOHO and TRACE imaging telescopes have revealed that dynamic events are
continuously occurring in the solar transition region and corona. Such events produce hightemperature plasmas and high-velocity plasma motions in these regions. Although some of these are
observed with a SOHO imaging spectrometer (CDS; Harrison et al. 1995), EIS will be able to track
the change of plasma properties (temperature, density, Doppler velocity, and nonthermal velocity)
as a function of time with better cadence and spectral information due to its large effective area.
2. Solar-B EIS
EIS consists of an imaging spectrometer (STR; lower left in Fig. 1), which is mounted on the
optical bench unit of Solar-B, and an Instrument Control Unit (ICU) that is mounted on the bus
panel of the spacecraft. The EUV optical components and the CCD camera system are set on the
base panel of the spectrometer (Fig. 1). Thin Al filters are accomodated at the EIS aperture and near
the focal plane for the purpose of visible light rejection. The entrance filters are contained in a
vacuum enclosure and are protected from the acoustic load during the launch. Two vacuum doors of
the enclosure are opened in both directions in orbit. A primary mirror focuses the EUV image of the
Sun on the four slits, and a rotary shutter controls the amount of EUV light that goes toward a
toroidal concave grating. The grating refocuses the image of the slit on the two CCDs, each of
which has 2048 x 1024 pixels, dispersing the EUV light in the wavelength direction. All mechanical
actuators and two contamination monitors in the spectrometer are controlled by a Mechanism and
Heater Controller (MHC). Since the spectrometer is directly exposed to space environment, the
STR thermal condition is adjusted by MHC-controlled heaters to the environment of low heat
leakage through MLI (Multi-Layer Insulation). The snout at the aperture is designed to prevent
external contaminants from reaching the pre-filter sulface.
Figure 1. Optical layout of EIS. An overall view of the spectrometer is shown in the inset.
3. EIS Optics and Count Rate
The EIS optical layout is shown in Fig. 1. Since the reflectance in the short EUV wavelength range
from 150 Å to 300 Å is low, the number of reflections is minimized. The off-axis parabolic primary
and toroidal diffraction grating have two sectors, each of which is coated with a Mo/Si multilayer
tuned for high reflectance at the selected wavelengths. The peak efficiency is achieved at 195 Å for
the shorter wavelengths and at 270 Å for the longer wavelengths. The overall property and
performance of EIS are summarized in Table 1.
Telescope
Type
Pre-filter
Primary mirror
Focal length
Off-axis distance
Plate scale at slit
Multilayer coating
Fine mirror scan range
Large FOV translation
Spectrometer
Type
Slit
Focal-plane filter
Normal-incidence telescope with two multilayer-coated primary
mirror
1500 Å Al filter for visible-light rejection
Off-axis parabola of 150 mm diameter
1939 mm
70 mm
9.4 m/arcsec
Mo/Si pairs; highest efficiency at 195 Å and 270 Å
+/- 256 arcsec solar image motion in EW direction
+/- 850 arcsec shift of FOV center in EW direction
Normal-incidence spectrometer with a multilayer-coated grating
and two back-illuminated CCDs
1”,2”,40”,266” widths
Supplementary Al filter of 1500 Å thickness
Concave grating
Multilayer coating
Wavelength coverage
CCD camera
Dispersion
Pixel equivalent width
Pixel scale at CCD
Toroidal surface, holographic and laminar ruling, uniform line
spacing of 4200 grooves/mm
Mo/Si pairs; highest efficiency at 195 Å and 270 Å
170-210 Å on CCD1 and 250-290 Å on CCD2
Back-illuminated CCD
Marconi 42-20, MPP
Pixel size
13.5 m
Format (spatial x spectral)
1024 x 2048 pixels each
Readout-window size
512 x 2048 pixels max
1.65 Å/mm or 0.0223 Å/pixel
34.3 km/s at 195 Å, 26.3 km/s at 284 Å
1.0 arcsec/pixel
Table 1. EUV Imaging Spectrometer overview
The performance of the FM (Flight Model) mirror and FM grating has been measured (Seely et al.
2004). From available information of the occultation of pre-filter frames, the filter transmission, the
reflenctance of primary mirror, the efficiency of grating, and the quantum efficiency of backilluminated CCD, the EIS effective area is calculated as shown in Fig. 2. For a typical active region
and a GOES M2-class flare the expected photon count rates in 1”x 1” area for emission lines of
wide-temperature coverage are calculated in Table 2 (not reported) with the help of the CHIANTI
spectral code (Dere et al. 1997). We expect high-count rates for coronal emission lines (106.0-106.4
K) and low count rates for the transition-region lines (105.4-105.8 K). The accuracy of Doppler
velocity and line width measurement is determined by the number of detected photons. Figure 3
(not reported) shows one-sigma errors in these parameters for an active-region line and a flare line.
A high-accuracy velocity measurement is achieved with EIS in a short exposure observation.
Figure 2. EIS effective area. The reflectance of primary mirror and the grating efficiency are
estimated from FM calibration. Preliminary values are used for the transmittance of filters and the
CCD quantum efficiency.
4. Slit Selection
EIS has four slits at the focal plane of the primary mirror. The slit direction is parallel to the northsouth direction of the Sun when the spacecraft is nominally operated. The length of the slits is long
enough to cover all CCD pixels in the spatial direction, corresponding to 1024 arcsec. The chosen
slit-widths are 1, 2, 40, 266 arcsec. One of them is selected by the rotation of a slit exchanger.
Narrow (Wide) slits of 1” and 2” (40” and 266”) slits are used for line-profile spectroscopy (for
EUV imaging). An EUV image taken with the 40” slit is little contaminated by nearby emission
lines and will almost be a monochromatic EUV image that is broadened by the velocity component
at each point. Images obtained with the 266” slit will be similar to the Skilab S082 slit-less
spectroheliograms that are partially contaminated by nearby emission lines.
5. Field of View and Size of Observing Region
The EIS field of view (FOV) is determined by (1) spacecraft (S/C) body pointing, (2) position of
the primary mirror in the S/C X direction, (3) range of raster scan, and (4) size of detector in the
spatial direction. The spacecraft pointing determines the EIS FOV center when the EIS mechanical
actuators are at their origins. There are two mechanical actuators to change the EIS observing
region. One allows changes in the center of the EIS observing region up to +/- 850” in the S/C X
direction, nominally in the east-west direction, by moving the primary mirror laterally. Since this
motion may slightly affect the stability of spacecraft attitude, it will be used with the agreement of
the SOT and XRT teams. The other allows to change the size of raster area with a scanning motion,
nominally in the east-west direction, by tilting the primary mirror around the S/C Y axis. The scan
range around its origin is +/- 256” at least in the image motion on a slit. As described in the
previous section, 1024 CCD pixels (=1024”) in the spatial direction are fully illuminated by the
EUV light. Thus, the size of EIS FOV is determined by the stroke of two actuators in the direction
perpendicular to the slit times the fixed number of 1024 arcsec in the direction parallel to the slit.
The EIS observing region is selected within the EIS FOV. The maximum length of 512” along the
slit can be observed at a single exposure. This parameter is selected as one of parameters of the
CCD spectral window that is defined in the EIS observation table called the EIS observation
sequence.
Figure 4. (a) Size of EIS wide slits and (b) EIS field of view on an SOHO/EIT Fe XII 195 image
taken near the solar minimum (16 May 1996). See text for details.
Examples are given in Fig. 4. Figure 4a shows the maximum size of observing regions by using
40”- and 266”-width slits in a single exposure, superimposed on solar coronal structures. An area
enclosed by dashed lines in Fig. 4b becomes the EIS FOV for the narrow slits (1” and 2” widths)
when the EIS optical axis is looking near the center of the Sun, as shown in the figure, under the
condition that the two mechanical actuators stay at their origins. In the case of wide slits (40” and
266” widths) the EIS FOV becomes slightly larger in the east-west direction. Observers can set an
observing region within the FOV. A box in the bold line at the Sun center shows the largest raster-
scan area (NS x EW = 512” x 512”) by a scanning motion of the primary mirror alone. Boxes in
bold line near the east and west limbs indicate the raster-scan areas of the same size for cases in
which the maximum stroke of the lateral motion of the primary mirror is used.
6. Data rate of EIS Observations
After readout of the CCD and data transfer to the ICU, the latter picks up spectral windows
according to the information in the EIS observation sequence. Observers can define up to 25
spectral windows on the two CCDs in the observation sequence. The ICU collects data in the
spectral window, reshapes the data as the EIS image packets with the EIS image header, and
transfers the packets to the Solar-B Mission Data Processor (MDP). The data are compressed in the
MDP with a hardware chip. The compressed data are transformed into CCSDS telemetry packets in
the MDP and the packed data go to a data recorder or a transmitter through a format change with a
Data Handling Unit (DHU). The bottle-neck link for EIS data is located between MDP and DHU
and limits the maximum rate for EIS to about 260 kbps. The average data rate for all the science
instruments is determined by the frequency of downlink contacts. If we assume 15 downlink contact
each of 8 min duration, the average data rate becomes about 300 kbps for all the scientific
instruments. If EIS is allowed to use 15% of this rate, 45 kbps becomes the average rate for EIS
observations. Since a maximum EIS observing rate of 260 kbps is possible in the MDP, a high-data
rate observation for a limited duration can be programmed by (1) inserting an observation period
with a lower data rate to maintain the average data rate or (2) negotiating with the SOT and XRT
teams to use a higher rate.
The data rate of EIS images is determined by (1) input data rate from ICU to MDP and (2)
compression ratio (compressed data size/input data size) in the data compression process at the
MDP. The former is related to the size of data in a single exposure and the cadence. These are
controlled by the EIS observation sequence in the ICU. For the latter, there are two types of data
compression. One is a bit reduction from 16 to 12 bits on a pixel basis, and the other is a 12-bit
image compression either by loss-less DPCM (Differential Pulse Code Modulation) or by lossy
JPEG compression. By tuning the Huffman and Quantization tables for the image compression, the
compression error of JPEG can be controlled within the photon noise.
Observation
Slit
Raster step per exposure
Observation duration per orbit
Spectral window: width
: height
Number of spectral windows
Compression
Cadence of exposure
Cadence of raster image
Short-duration data rate
Average data rate
4’ x 4’
Raster
1” slit
1” shift
97 min
16 pixels
256 pixels
14
20%
3 sec
12.8 min
44.8 kbps
44.8 kbps
8’ x 8’
Raster
1” slit
1” shift
25.6 min
16 pixels
512 pixels
25
20%
3 sec
25.6 min
160 kbps
42.2 kbps
0.5’x 4’
Imaging
40” slit
97 min
32 pixels
256 pixels
10
20%
4.5 sec
42.7 kbps
42.7 kbps
4’x 4’
Imaging
266” slit
97 min
256 pixels
256 pixels
4
20%
14 sec
43.9 kbps
43.9 kbps
Table 3. Examples of EIS data rate.
Examples of EIS observations during an orbit and the expected data rate for each example are
shown in Table 3. A data compression of 20% is assumed. The second column shows an example of
a high-rate observation for a short period keeping the average data rate at 45 kbps.
7. EIS Data Analysis
The EIS data for analysis are provided in FITS format after the process of restoration from the EIS
CCSDS packets and following decompression. The EIS analysis software packages are developed
using the IDL software language, which is adopted in the Solar Soft system. Raw EIS image data
will be distributed with basic analysis software packages and calibration data.
Acknowledgments. The author would like to acknowledge the UK and US EIS program scientist
and EIS development team members for their contribution to the EIS program.
References
See the original paper.