**[8. Scientific Payload]{.underline}**

The strawman payload discussed in this chapter contains sufficient
detail for the system design but may not reflect the finally selected
suite of instruments.

**8.1 Measurement requirements**

The Helios, Ulysses, Yohkoh, SOHO and TRACE missions provide the
expertise and experience for designing the payload of the Solar Orbiter.
It must be state-of-the-art, withstand the considerable thermal load at
45 R~s~, comply with the general requirements of a low-mass, compact and
integrated design, make use of on-board data compression/storage and
require a modest data transmission rate. The payload meeting the solar
and heliospheric science objectives encompasses two instrument packages:

-   *In-situ* heliospheric instruments (see Table 8.1 below): Solar wind
    plasma analyser measuring electrons and ions, magnetometer for DC
    and AC fields, particle detector for solar energetic ions and
    electrons, interplanetary dust detector, neutral particle detector,
    radio instruments. Heritage for such instruments is from Helios,
    SOHO, Ulysses, Wind, ACE.

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-   Remote-sensing solar instruments (see Table 8.4 below): EUV/X-ray
    full-Sun and high resolution imager; high-resolution EUV
    spectrometer covering selected emission lines from the chromosphere
    to corona; high-resolution visible light telescope and magnetograph;
    EUV and visible-light coronagraph; solar neutron and γ-ray detector
    and a radiometer. Heritage for such instruments is from Yohkoh,
    SOHO, TRACE.

    **8.2 *In-Situ* heliospheric instruments**

    **8.2.1 Solar Wind Plasma Analyser (SWA) *Still old prop.
    Version!***

The SWA will measure separately the three-dimensional distribution
functions of the major solar wind constituents: protons, alpha particles
and electrons. The basic moments of the distributions, such as density,
velocity, temperature tensor, and heat flux vector will be obtained
under all solar wind conditions and be sampled fast enough to
characterise fully the fluid and kinetic state of the wind. In this way
we will be able to determine possible non-gyrotropic features of the
distributions, ion beams, temperature anisotropies and particle
signatures of wave excitation and dissipation.

The SWA will measure protons, alpha particles and electrons. In
addition, measurements of representative high-FIP elements (the C, O, N
group) and of low-FIP elements (such as Fe, Si or Mg) will be carried
out, to obtain their abundances, velocities, temperature anisotropies
and charge states, and in order to probe the wave-particle couplings
(heavy-ion wave surfing) and determine the freeze-in temperatures (as a
proxy for the coronal electron temperature). Also, pick-up ions of
various origins, such as singly-ionised ions stemming from
interplanetary dust, will be measured.

The SWA will unambiguously separate ion species and provide M/Q
determination with an energy resolution of at least 5 % in ∆(E/Q) /(E/Q)
and range between 10 eV/Q and 30 keV/Q. The angular resolution will be
better than 5° in azimuth and elevation angle. The electron analyser
will be capable of resolving the strahl (heat flux) and cover a large
fraction of 4π in solid angle, with a pitch-angle resolution better than
22.5° and an energy range between a few eV up to 10 keV. To cope with
fast wave-particle effects to be expected in the velocity distributions,
the sensors will have time resolutions of about 1s or even shorter. The
specifications can be met by existing state-of-the-art plasma
instruments, using time-of-flight techniques with post-acceleration for
particle discrimination.

**8.2.2 Plasma Waves Analyser (PWA) *Still old prop. Version!***

The PWA will identify the various plasma waves and kinetic modes
comprising the high-frequency part of the fluctuation and turbulence
spectra. The instrument will cover a broad band in frequencies,
extending from some Hz into the MHz range. The expected field strength
may range between a few μV/m and mV/m, and perhaps up to a V/m for the
convection electric field. The magnetic field strength is expected to
vary between a few nT and μT, with large differences between the
longitudinal and transverse components with respect to the mean magnetic
field and the solar wind flow direction (substantial Doppler shifts are
to be expected).

The electromagnetic fluctuations are measured with search coil
magnetometers and electric field dipole antennas, minimised in size and
arranged in a compact configuration. The PWA will perform on-board
processing of the data and deliver selectable wave forms, to identify
nonlinear coherent fluctuations. To resolve the vector components of the
electric and magnetic fields, which is scientifically highly desirable
and required to determine wave modes unambiguously, the search-coils are
to be mounted on short booms required for magnetic cleanliness reasons.
More than one electric field antenna is needed, which will be short
enough to adjust to the variable Debye length of about one meter.

**8.2.3 Magnetometers (MAG)** RS, June 5, 2000

The magnetometer will measure the interplanetary magnetic field
components over a broad dynamic range from a few nT up to the μT range,
and for frequencies well below the local heavy-ion or proton
gyrofrequencies. The magnetometer may consist of a triaxial identical
sensor system. A compromise has to be found between the requirements of
the spacecraft and payload systems and the demand of the magnetometer
for magnetically clean conditions.

The heritage of MAG includes dozens of similar instruments flown on
various magnetospheric, planetary and deep space missions. Expertise has
been built up in several European countries, e.g. the UK, Germany,
Russia. After all, the MAG sensor and electronics for the Solar Orbiter
requires no new development. State of the art instruments can be built
for less than 1 kg. Further mass savings appear possible.

MAG might consist of a 3-axis fluxgate magnetometer with the sensor
mounted at the end of an extendable boom (length tbd) which is
positioned in the shadow of the Orbiter. Other mounting locations such
as on the high-gain antenna structure appear possible, provided MAG
stays in permanent shade. The MAG proposed for the STEREO mission might
serve as a guideline. It is in turn based on the design used for the
NEAR and ACE missions. It measures fields in 8 ranges (automatically
selected) from 4 nT to 65536 nT with a resolution of 30 ppm, from DC up
to AC frequencies of 500 Hz. The parameters can be adjusted according to
the particular needs o the Solar Orbiter mission. The DPU design might
resemble the ones developed, e.g., for the AMPTE, Ulxsses and Rosetta
missions. It could be shared with either the plasma instruments or the
wave instruments.

The mass of the sensor package will be less than 1 kg, the average power
consumption is about 1 W, and the required telemetry is of the order of
200 bits/s.

Magnetic cleanliness of a spacecraft as required by a magnetometer on
board has been a cost driver for some missions in the past. However,
more recent experience has shown that there are in fact acceptable
low-cost engineering solutions possible. A reasonable magnetic
cleanliness program can be performed, if the MAG experimenters work in
close contact with the all other teams in order to ensure proper
engineering at an early phase and to make other precautions without
impacting the mission schedule.

Similar schemes were applied in the past for several successful
missions.

**8.2.4 *Energetic Particle Detector (EPD) Still old prop. Version!***

The EPD measurements on the Solar Orbiter will allow to study the
sources, acceleration and propagation of solar energetic particles in
association with coronal and interplanetary shocks. Energetic pick-up
particles originating from outgassing or sputtering of near-Sun dust
should be measured as well in conjunction with the SWA measurements. The
EPD will determine chemical and charge composition of ions in a wide
energy range, from about the typical solar wind energies of a few keV to
several 100 MeV/nucleon for protons and heavy ions. Electrons should be
measured from 10 keV to 10 MeV.

The combination of electrostatic E/Q-analysis with time-of-flight
E/M-determination and subsequent direct energy measurement in a solid
state detector has been employed in many EPDs in the past. Ongoing
research in solid state detector technology is promising and will
certainly lead to further improvements in energy thresholds, and noise
level as well as to size reduction. It is essential that all
measurements are done fast, typically at 1 s, with complete angular
coverage to resolve the pitch-angle distributions. These requirements
can readily be met by a low-mass and compact sensor assembly with
different look directions and modest aperture sizes. We propose to place
several detectors at different positions onboard the spacecraft, in
order to obtain a sufficient field of view and angle-of-incidence
coverage without using a scanning platform.

**8.2.5 Dust Detector (DUD) *Still old prop. Version!***

The DUD will analyse interplanetary dust particles with respect to their
composition and orbital characteristics. The spatial distribution of
particles with masses ranging between 10^-16^ g and 10^-6^ g should be
determined. These objectives can be reached with an instrument of
Ulysses, Galileo or Cassini heritage. The minimal dust experiment may
consist of a multiple sensor assembly, with each sensor looking in a
different direction to resolve incidence angles. A low-mass
dust-detecting element could consist of a thin polarised polyvinylidene
fluoride foil with conducting electrodes on opposite sides. The particle
impact irreversibly depolarises this foil in the penetration crater
volume and thus creates a fast (ns) current pulse as input to the
electronics. The calibrated detector signal carries information about
mass and velocity of the impact particle. Its mass can be determined by
the foil thickness setting a threshold value. The differential dust
particle fluxes are directly obtained from the differences of the
integral fluxes measured with foils of various thickness. Full chemical
analysis and mass resolution of the important elements H, C, N, O of
carbonaceous material and the metals like Na, Mg, Al, Si, Ca and Fe
contained in chondritic silicates is desirable.

**8.2.6 Radio Instruments *Still old prop. Version!***

*[Radio Spectrometer (RAS)]{.underline}* The RAS will measure the solar
and interplanetary radio waves in the frequency range from 100 kHz to 1
GHz, with a sweep period of 1 s and a high spectral resolution (∆f/f ≈
0.01). The RAS will observe plasma processes associated with energetic
electrons from the low corona up to 50 R~s~. It will probe the plasma at
distances ranging from the solar surface to the spacecraft location,
thereby connecting the low-altitude coronal regions observed by the
optical instruments with the near-Sun heliospheric conditions specified
by the in-situ measurements. The instrument could be divided into four
sub-spectrometers according to a frequency ratio of 1:10. Each
sub-spectrometer may consist of 128 channels, leading to a total data
rate of 4 kb/s.

#### 

#### *Coronal Radio Sounding (CRS)* A passive radio science experiment can be carried out using the available radio links. The S/C will pass behind the Sun (superior solar conjunction) on many occasions, making it possible to investigate the solar corona by radio sounding at solar distances much less than the perihelion distance down to at least 2 R~s~. Integrated line-of-sight parameters such as electron content (densities), Faraday rotation, scintillations and angular broadening can be recorded. The CRS experiment with a two-way radio link via the S/C high-gain antenna (ranging capability) would require a radio subsystem with dual-frequency phase coherent downlinks at X-band and Ka-band. Linear polarisation of the downlink signals would enable Faraday rotation measurements as an option. The RF power for both radio links (X and Ka) planned for this mission is sufficient for this investigation. A two-way dual-frequency coherent radio link (X-band uplink; X-band and Ka-band simultaneous coherent downlinks) is considered as the optimal configuration for a sufficiently stable link and for detecting signatures of CME events traversing the radio ray path. If only a one-way radio link (S/C to ground) is feasible due to operational constraints during solar conjunction, the transmitted radio signal must be stabilised by an onboard ultra-stable oscillator (frequency stability 10^-11^ to 10^-12^ at 3 s integration time, 200g , 3 W). 

**8.2.7 Neutral Particle Detector (NPD)**

## [Instrument characteristics]{.underline}

The Lyman**-α** line of neutral hydrogen is used in spectroscopic
observations on SOHO to determine the solar wind velocity profile in the
corona up to several solar radii. The NPD on the Solar Orbiter will
measure primarily neutral hydrogen and thus render possible the
detection of the interplanetary \"neutral solar wind\". The NPD will
also measure energetic neutral atoms emitted from various coronal
sources, and thus enable images of these coronal emission regions to be
constructed from rays of neutral atoms with different velocities. As a
baseline, the instrument parameters could be: Energy/mass range: 0.6 to
100 keV/nucleon, moderate velocity resolution; mass resolution is a
secondary objective. The field of view (for the imaging capability) of
the NPD should be centred at about 17° off the spacecraft-Sun-centre
line .

The instrument must be capable of detecting neutral atom particle
populations originating in the vicinity of the Sun. Therefore, the
instrument entrance system must block radiation emerging from the solar
disk. The UV scattered in the solar corona, i.e. Ly**-α**, must be
suppressed by a factor of more than 10^-10^. The entrance system must
reduce ions and electrons of the solar wind plasma by a factor of about
10^-8^ while transmitting neutral atoms. The instrument should detect
neutral atoms in the typical solar-wind and supra-thermal velocity range
between about 350 and 5000 km/s. The neutral atom flux is expected to be
about 100-1000 atoms /(cm^2^ s) at 0.21 AU, but it could be up to
(estimated) 10^6^ atoms/(cm^2^ s) in a CME.

## [Instrument design]{.underline}

A feasible neutral atom detector may consist of two sections: The first
compartment accommodates the baffle plates keeping out solar radiation
and scattered UV and the deflection unit for charged particles. The
second section contains the time-of-flight (TOF) double-coincidence
detector. Mid of the two sections is the pinhole/slit that defines the
optical properties of the neutral atom detector.

The direct solar radiation is blocked by a collimator. For the
suppression of scattered light within the collimator, the plates are
covered with saw-teeth like structures. The collimator consists of,
alternatively positive and negative, high-voltage biased plates and
therefore acts a charged-particle deflector. The collimator plate
length, distance and potential determine the cut-off of the
charged-particle transmission (about 100 keV/q). UV originating from the
solar corona, i.e. Ly**-α**, is transmitted through the collimator. The
grid in front of the pinhole reflects the UV. The grid consists of 2
gratings with a grating constant matched to the Ly**-α** peak (gratings
are perpendicular to account for different polarisation angles of the UV
light). The UV photons are scattered according to diffraction theory.
The grids are more transparent for particles than for UV photons (ratio
up to 10^5^ per grating, in dependence on the UV wavelength).

The neutral atoms pass through the pinhole (or slit) and a 0.5 to 1.1
**μ**g/cm^2^ carbon foil. As they passed the carbon foil, there is some
probability to release secondary electrons which are accelerated towards
the START MCP (multi-channel plate). The atoms travel through the
3-cm-wide TOF drift region and hit the rear STOP MCP. The time of flight
for 1 keV/n-100 keV/n atoms is about 8 to 80 ns. The position resolution
is 1-D for the START MCP and 2-D for the STOP MCP. The flight pass of
the neutral atom can be calculated from these measurements by an
inversion of the instrument function. The detector determines the
velocity (energy/nucleon) and the flight pass (imaging) of the neutral
atoms.

![](media/image1.png){width="4.825in" height="3.1006944444444446in"}

Figure 8.1 The instrument schematics of the Neutral Particle Detector.

An alternative instrument design approach could consist of a solar disc
baffle, the pinhole, the charge particle deflection system and the TOF
unit. While this approach has superior qualities subduing the solar disc
radiation, the length of the charged particle deflection system is very
limited and the TOF section and the detectors are much larger and
therefore require more mass.

## [Resources, mass budget and envelope]{.underline}

\- Volume envelope: 10 x 10 x 20 cm^3^ including mounting lugs.

\- Mass: Housing, collimator, TOF structure: up to 300 g MCPs; MCP
mounting: 150 g

Electronic boards: HV supply, TOF and interface to DPU: 350 g

Connectors: 50 g

Instrument mass: 0.85 kg

DPU share: 150 g

Total mass: 1 kg; total mass with 20% margin: 1.2 kg

\- Power: less than 2 W

\- Data/DPU: It is envisaged to use a common DPU with the solar wind
plasma instrument. The required tasks of the DPU are command transfer,
data handling and the provision of memory for the in-flight look-up
tables and software.

\- Field of view: 17° off spacecraft-Sun-centre line

(due to the up to 10° aberration for a 1 keV/n particle at 0.21 AU)

\- Telemetry: less than 0.5 kb/sec

\- Possible accommodation: near the dust detector

### Heritage

Neutral atom detectors have been flown on SOHO, CASSINI and IMAGE.
Time-of-flight instruments are for example flown on ULYSSES, SOHO and
ACE. The later two made extensive use of the shared DPU concept for
different sensors.

**Table 8.1: *In-situ* heliospheric instrumentation**

+----------+-----+-----------+-------+----+----------+-----+--------+
| Name     | A   | Me        | Speci | Ma | Size     | Po  | Tel    |
|          | bbr | asurement | fica- | ss |          | wer | emetry |
|          | e-v |           | tions |    | cm x cm  |     |        |
|          | iat |           |       | kg | x cm     | W   | kb/s   |
|          | ion |           |       |    |          | att |        |
+----------+-----+-----------+-------+----+----------+-----+--------+
| Solar    | SWA | Thermal   | 0-30  | 6  | 20 x 20  | 5   | 5      |
| Wind     |     |           | k     |    | x 20     |     |        |
| Plasma   |     | ions      | eV/Q; |    |          |     |        |
|          |     |           |       |    |          |     |        |
| Analyser |     | and       | 0-10  |    |          |     |        |
|          |     | electrons | keV   |    |          |     |        |
+----------+-----+-----------+-------+----+----------+-----+--------+
| Plasma   | PWA | AC        | μ     | 3  | 20 x 20  | 3   | 4      |
|          |     | Electric  | V/m - |    | x 20     |     |        |
| Wave     |     |           | V/m   |    |          |     |        |
| Analyser |     | and       |       |    |          |     |        |
|          |     | magnetic  | 0.    |    |          |     |        |
|          |     | fields    | 1nT - |    |          |     |        |
|          |     |           | μT    |    |          |     |        |
+----------+-----+-----------+-------+----+----------+-----+--------+
| Magne    | MAG | DC        | to    | 1  | 10 x 10  | 1   | 0.5    |
| to-meter |     | magnetic  | 500   |    | x10      |     |        |
|          |     | field     | Hz    |    |          |     |        |
+----------+-----+-----------+-------+----+----------+-----+--------+
| E        | EPD | Solar and | Ions  | 4  | 20x 20   | 3   | 2.5    |
| nergetic |     | c         | and   |    | x20      |     |        |
| Particle |     | osmic-ray | elec  |    |          |     |        |
| Detector |     | particles | trons |    |          |     |        |
|          |     |           | .     |    |          |     |        |
|          |     |           | 01-10 |    |          |     |        |
|          |     |           | MeV   |    |          |     |        |
+----------+-----+-----------+-------+----+----------+-----+--------+
| Neutral  | NPD | Neutral   | 0.6 - | 1  | 10 x 10  | 2   | 0.5    |
| Particle |     | hydrogen  | 100   |    | x 20     |     |        |
| Detector |     | and       | keV   |    |          |     |        |
|          |     |           |       |    |          |     |        |
|          |     | atoms     |       |    |          |     |        |
+----------+-----+-----------+-------+----+----------+-----+--------+
| Dust     | DUD | Inter     | Mass  | 2  | 10 x 10  | 1   | 0.5    |
| Detector |     | planetary | (g):  |    | x 10     |     |        |
|          |     |           |       |    |          |     |        |
|          |     | dust      | 10^-  |    |          |     |        |
|          |     | particles | 16^-1 |    |          |     |        |
|          |     |           | 0^-6^ |    |          |     |        |
+----------+-----+-----------+-------+----+----------+-----+--------+
| Radio    | RAS | Coronal   | 100   | 4  | 10 x 10  | 10  | 1      |
| Spec     |     | radio     | kHz - |    | x 10     |     |        |
| trometer |     |           |       |    |          |     |        |
|          |     | waves     | 1 GHz |    | antennae |     |        |
+----------+-----+-----------+-------+----+----------+-----+--------+
| Radio    | CRS | Wind      | X     | (U | 5 x 5 x  | 3   | 0      |
|          |     | density   | -band | SO | 5        |     |        |
| Sounding |     |           |       | 2  |          |     |        |
|          |     | and       | Ka    | 00 |          |     |        |
|          |     | velocity  | -band | g) |          |     |        |
+----------+-----+-----------+-------+----+----------+-----+--------+

**8.3 Remote sensing solar instruments**

**8.3.1 General considerations**

An integrated ensemble of high resolution remote sensing instruments is
proposed for the Solar Orbiter. Emphasis is put here on the scientific
requirements and their feasibility. All the instruments will resolve the
small-scale dynamical processes in the solar atmosphere. The combination
of instruments proposed here will provide a complete set of
multi-wavelength measurements required to understand the Sun's
magnetised atmosphere -- from below the photosphere up into the extended
corona.

Critical issues like heat input into the instruments and their
mechanical structures, mass and power requirements have been
investigated. The feasibility of the instruments has been demonstrated.

**8.3.2 Visible-light imager and magnetograph (VIM)**

**Scientific Drivers**

The purpose of the VIM is to determine the magnetic field boundary
conditions for the MHD processes observed by other remote sensing
instruments. It will observe the morphology, dynamics, and strength of
the magnetic elements and flux tubes at the photospheric level with a
resolution that is consistent with the resolution of the EUV telescopes.
It will also provide images, Dopplergrams and magnetograms from the side
of the Sun which is invisible from Earth.

VIM will be a vector magnetograph. Having vector capabilities is of
fundamental importance to understand the nature of photospheric fields
(Are they vertical unipolar fields? Do they harbour complex neutral
lines with horizontal sheared fields?). Having vector capabilities is
also the only way in which quantitative inferences of the magnetic field
in the transition region and corona can be made (from force-free or full
3D MHD extrapolations). The importance of such data for the
understanding of coronal processes (like e.g. the initiation of CMEs)
has been amply demonstrated by SOHO.

VIM will also produce line-of-sight velocity maps by observing two
points at each side of the spectral line. These maps can be used,
through local helioseismology techniques, to understand subsurface
flows. The internal structure and dynamics of the near polar regions of
the Sun is of paramount importance and perhaps *the* key to our
understanding of the solar cycle.

The key scientific goals of the VIM are:

-   To resolve solar magnetism down to its fundamental length scale (100
    km).

-   To provide measurements of the \`magnetic carpet\` which drives
    chromospheric and coronal activity.

-   To observe and accurately quantify for the first time the surface
    polar magnetic field of the sun.

-   To unveil the small scale photospheric dynamo.

-   To provide the first magneto- and Dopplergrammes of the far side of
    the Sun.

-   To measure rotation and flows near the Sun's poles using techniques
    of local area helio-seismology and thereby provide crucial
    constraints on solar dynamo theories.

**Instrument Concept**

The optics of the VIM include

-   the 25 cm diameter High Resolution Telescope (HRT),

-   the 2 to 5 cm diameter Full Disk Telescope (FDT), and

-   the Filtergraph Optics (FO).

Both telescopes are used alternatively and share the filtergraph optics
with the only science detector of the VIM. While the HRT delivers an
image of a part of the solar surface with a resolution at its
diffraction limit near 0.5 arcsec, the FDT is capable to adjust its
focal length such that the solar disk fills entirely the detector at all
orbital phases. The HRT will provide high cadence filtergrams, Doppler
maps and magnetograms for limited times during the phases of closest
approach and polar proximity, while the FDT provides low cadence
filtergrams and magnetograms during all orbital phases.

The filtergraph consists of a two Fabry-Perot interferometers with fixed
gap widths and a narrow band interference prefilter. The whole assembly
is mounted on a rotating stage in a telecentric beam. Rotation of the
etalons and the filter provides tuning of a 5 mĂ passband around the
magnetically sensitive Fe I 630.2 nm line, also allowing for Doppler
shifts of up to 40 km/s resulting from the spacecraft orbit.

*High Resolution Telescope*

The HRT is a Gregorian reflector with 250 mm entrance aperture and an

effective focal length of 4250 mm. **A broad-band interference filter
covers the entire telescope and reflects most of the incident sunlight,
with the exception of the spectral range of interest.** A polished field
stop at the primary focus reflects most of the sunlight which is
incident on the primary mirror outside the spacecraft. A 1.2 mm diameter
hole in the field stop transmits the science field of view. The size of
the field stop is sufficient to permit pointing of the telescope
anywhere on the solar disk. An external cover shields the telescope
during off-sun pointing conditions.

The F ratio of the parabolic primary is F/2. The elliptical secondary
mirror has a focal length of 134 mm and provides a magnification of 8.5
for the F/17 telescope focus. The secondary is actively mounted with
three translation degrees of freedom for telescope alignment control and
pointing.

Two flat mirrors fold the beam in the back focus area of the primary.
The first mirror provides tip-tilt control in order to maintain the
image position within 0.01 arcsec relative alignment in the presence of
3 arcsec rms spacecraft pointing error. The second mirror is a dichroic
which reflects the light longwards of 600 nm towards the telescope
focus, while transmitting the shorter wavelengths towards the image
position and alignment sensor. The concept for the alignment sensor is
described below.

The field of view covers a square with 10 mm side. It contains only a
few percent of the total incident power and presents no further thermal
problem.

Figure 1: perspective view of the optical layout of the high resolution
telescope and filtergraph.

The optics of the HRT are made from SiC which combines excellent
thermostability with lightweighting capability. Astrium has just
completed under contract with the KIS a 360 mm all C/SiC mirror which
meets optical surface quality specifications. The possibility to
fabricate the entire telescope structure from SiC is under
investigation.

*Full Disk* *Telescope*

The FDT is mounted on the same optical bench as the FO. It consists of
four singlet lenses in a zoom configuration, a modulator and two folding
mirrors. The second folding mirror can be moved into the optical beam of
the HRT just prior to the focal plane in order to direct the FDT light
towards the filtergraph.

The axial position of two elements of the Zoom optics is changed during
the orbit in order to adjust the focal length between 233 mm and 888 mm.
A fixed aperture stop results in variable diameter of the entrance pupil
which changes from 50 mm for the long focal setting to 20 mm for the
shortest setting. An interference filter at the entrance transmits the
spectral region which is used by the filtergraph and the image position
control system and reflects the rest outside the spacecraft.

Figure 2 shows the annotated optical layout of the FDT from a similar
point of view as the HRT in Figure 1. Note that the layout is
\"horizontal\" (in the XZ plane) compared to the vertical (YZ plane)
layout of the HRT proper. A separate modulator is necessary for the FDT
in order to perform polarization modulation within the centered optical
path before the first oblique reflection.

Figure 2: perspective view of the optical layout of the full disk
telescope.

The FDT has its own pointing alignment control system, which consists of
a simple limb sensor and the same combination of tip-tilt mirror and
dichroic as the HRT. It shares the control electronics with the HRT.

*Filtergraph Optics*

The FO consists of two singlet lenses which combine the functions of
collimator and camera with a 1:2 magnification. Two off-axis mirrors in

between provide a 1:1 pupil relay function and an F/66 telecentric

intermediate image plane. Two 50 mm diameter Fabry-Pérot etalons at the
telecentric image plane provide the required spectral resolution of
typically 50mĂ. If necessary, a separate narrow band interference filter
will block the secondary transmission maxima of the combined etalons.
About ±1° of rotation shifts the passband of the two filtergraphs by
±1Ă, which is sufficient to scan the spectral line and to account for
the S/C velocity effects. The principle proposed here is the same as
used in the TESOS filtergraph which operates at the Vacuum Tower
Telescope of the Kiepenheuer-Institut on Tenerife.

The FO are mounted together with the FDT on a common optical bench.

**Figure 3** shows a plane view of the filtergraph optics. The etalons
are not included but their position indicated.

*Figure* *3:* *optical arrangement of filtergraph optics.*

*Overall Configuration*

The combination of all optical subsystems is shown in Figure 4. The HRT
is the large cylinder on top of the Y and Z views. The FDT is to the
lower left of the Z view and near the bottom of the X view. The FO are
to the lower left of the Z view and near the top of the X view.

![](media/image5.png){width="4.35in"
height="2.1743055555555557in"}![](media/image6.png){width="1.707638888888889in"
height="2.1770833333333335in"}

Figure 4: plain views (top left: Y, top right: Z, bottom: X) of the
combined VIM. The HRT is arranged above the FDT and the FO, both of
which are mounted on a common optical bench. The optics fit a 300 mm x
400 mm x 1600 mm envelope.

*Polarization Modulation Package*

The Polarization Modulation Package (PMP) will allow VIM to provide
vector magnetograms of the region being observed. The PMP will produce
the modulation of the intensity at the CCD as a function of the input
polarization state. These intensity changes of the CCD measurements will
be used to recover the Stokes vector of the solar light.

The PMP will be located before any mirror reflection of the light beam.
In this way the polarization measurements will be free of the
instrumental polarization problems induced by the reflections. In the
present configuration, VIM will use two PMPs, one for the HRT and one
for the FDT. Each PMP will be composed of a rotating retarder and
followed by a fixed linear polarizer. The retarder follows the design of
the Solar-B instrument, which is based on the Advanced Stokes
Polarimeter (ASP) used at Sacramento Peak observatory. The retarder will
be positioned at 8 different positions per half rotation, each separated
by 22.5 degrees. At each of these positions the CCD will be read and the
8 readouts combined in different ways to determine the final Stokes
parameters (I,Q,U,V). Depending on the light flux budget, a number of
such independent determinations of the Stokes vector will be added on
four memory buffers to reduce the noise to the level demanded by the
scientific objectives (of around

10^-3^ times the continuum intensity level). The retarders will have a
retardance at the VIM nominal wavelength of 123 degrees. This value
ensures that the three Stokes parameters, Q, U and V, are all measured
with similar and large efficiencies. The rotating retarder will be
designed with a wedge system that eliminates beam wobbling and fringe
patterns. The exact waveplate rotation frequency needs to be defined,
but similar devices proposed for other S/C have rotation frequencies of
around 0.5 Hz. In this case the CCD (which must be synchronized with the
rotation of the waveplate) will be read at 8 Hz.

There are two options to simplify the actual design of the PMP. Both
will have to be studied further in a Phase A study. First, instead of
using two PMP, only one could be used located after the M4 mirror (that
selects the light from either the high-resolution or full-disk telescope
\[HRT and FDT\]). In this case, fixed half-wave retarders should be
inserted in each beam to compensate the instrumental polarization
produced by the mirror reflections. The second option to simplify the
PMP, will be the use of electro-optical devices instead of rotating
retarders. Liquid Crystal based retarders offer the possibility to
produce the modulation of the Stokes parameters without having to rotate
any mechanism. This Liquid Crystal retarders have been tested
successfully on the ground for solar polarization studies. Their
advantages are the use of a simple voltage signal to produce the Stokes
modulation and the lack of moving parts (saving mass and power). Their
disadvantage is their sensitivity to UV light. If Liquid Crystals are to
become an option, this last issue needs to be carefully studied.

*Optical Alignment and Image Stabilization*

The optical alignment of the VIM is an important and critical issue. Due
to large temperature gradients during the mission, it is important to
have reliable and accurate in-flight alignment capabilities. To this
end, we propose a wavefront control system which is capable to detect
low order modes of wavefront deformations in the telescope. A wavefront
sensor measures the actual state of the optical alignment and generates
an appropriate error signal. A control system converts this error signal
into actuation signals which are used to drive the position of the
secondary mirror M2 and the tip-tilt mirror (M3).

The detection principle is based on a correlation tracker which produces
tip and tilt error signals. However, instead of sensing the position of
a single image which is derived from the entire pupil of the telescope,
seven subapertures sense the local wavefront tilt in two zones of the
pupil. The information derived from the seven independently analyzsed
images of the same solar scene suffices to determine the coefficients of
a Zernike function decomposition of the wavefront error up to the third
radial degree. The coefficients for tip and tilt, defocus, and Seidel
coma are used as error signals for the control system. The system is
capable to measure and compensate axial and lateral displacement of the
secondary mirror M2 as well as dynamic image displacement errors, i.e.
it includes the functionality of a correlation tracker.

Figure 5: Optical configuration of wavefront sensor

Figure 5 shows a sketch of the wavefront sensor which is located in the
vicinity of the secondary focus. The sensor consists of a collimating
lens, a hexagonal lenslet array with seven - one central and six
peripheral - elements which cover the pupil of the telescope, and lenses
which reimage the lenslet array focal plane onto a matrix array detector
(CCD camera). The image scale of the pupil on the array is chosen such
that the illuminated areas of the lenslets are equal, taking also into
account the central obstruction. All lenslets are partially illuminated.
Figure 6 shows the configuration of the lenslet array.

![](media/image8.jpeg){width="2.1666666666666665in"
height="2.077777777777778in"}Figure 6: Hexagonal 7-element lenslet array
of wavefront control unit. The dashed lines indicate the position of the
pupil image.

*CCD* *cameras*

A variety of CCD detectors exist on the market that are fully adequate
for the VIM. The HRT and FDT telescopes have been designed for a CCD
with 2k x 2k pixels of around 10 micron size. Such CCDs are common on
the market today. Slight modifications of the existing EEV CCD42-40
(13.5 micron pixels) meet our requirements and they are proposed for use
on other missions (e.g. STEREO). These CCDs are thinned, back
illuminated low dark current cameras with typical quantum efficiencies
of 85 % (AR-coated) at the wavelength of interest for VIM. Shuttered and
shutterless (frame transfer) modes are available and it has to be
investigated which one suites our application better. As a result
different observing modes may be available (slow but large FOV, fast but
small FOV). Tradeoffs between photon statistics, image stabilization
system bandwidth, rotating frequency of the PMP waveplate and CCD pixel
read rate (1MHz is anticipated) will have to be studied. Marconi (EEV)
is presently developing custom CCDs for the Japanese led Solar-B project
that will be used in both imaging and spectroscopic modes. Their
experience can be of great value for VIM.

**Resources and Requirements**

*Telemetry needs -- Data compression.*

VIM will take advantage of the high rate observing mode by using a high
cadence data acquisition necessary for local helioseismology and
magnetic field dynamics. For example, in this mode one can read only
1024 x 1024 pixels (smaller FOV, but still large enough for magnetic
field studies) or readout the CCD with 2 x 2 on-chip summing (preserving
FOV but with slightly worse spatial resolution that is still adequate
for local helioseismology). In this case, 1 full frame (1024 x 1024, 12
bits/pixel) at a cadence of 1 per minute produces a data rate of 200
kb/s. By using efficient lossy compression algorithms with a reduction
factor of 10, the final rate will be around 20 kb/s. In this mode, we
are considering storing only one physical quantity each time (e.g. the
LOS magnetic field or the LOS velocity). Or one can read all 2048 x 2048
pixels, and send 4 physical magnitudes (for example all 4 Stokes
parameters) at a cadence of 1 frame each 15 minutes using a similar data
rate. In the low rate observing mode, sending 2 (4) physical magnitudes
over the full frame can be done at a cadence of 1 per hour and with a
data rate of 3 (6) kb/s.

*Mass breakdown and power requirements*

A rough breakdown of the estimated mass and power requirements is given
in the two tables below.

  --------------------------------------------------------- -------------
                                                            **Mass (g)**

  **1. High Res Telescope** (incl. ISS, re-imaging, filter, **19500**
  CCD)                                                      

  \- 25 cm telescope (mirrors + structure)                  4000

  \- prefilter (@ primary hole)                             400

  \- rotating quarter wave plate                            1000

  \- linear polarizer                                       200

  \- folding mirror                                         400

  \- tip-tilt mirror + PZT (incl. electronics)              1500

  \- cube beam-splitter                                     200

  \- ISS lens                                               200

  \- ISS detector                                           1000

  \- re-imaging system (4x200)                              800

  \- Fabry-Perots and Oven                                  3000

  \- shutter                                                400

  \- CCD camera (incl. electronics)                         2500

  \- harness (25% of total above)                           3900

                                                            

  **2. Full Disk Telescope**                                **3250**

  \- prefilter                                              200

  \- focal lenses (4\*200)                                  800

  \- rotating quarter wave plate                            1000

  \- linear polarizer                                       200

  \- folding mirror                                         400

  \- harness (25% of total above)                           650

                                                            

  **3. VIM enclosure**                                      **6000**

                                                            

  **4. VIM TOTAL**                                          **28750**
  --------------------------------------------------------- -------------

Table xx: VIM mass break-down

  --------------------------------------------------------- -------------
                                                            **Power (W)**

  CCD                                                       5

  Image Stabilization System                                5

  Fabry Perots & Oven                                       10

  PMP (x2)                                                  4

  **TOTAL**                                                 **28**
  --------------------------------------------------------- -------------

Table xx: VIM power requirements.

**Instrument Summary**

The basic design of the instrument is given in the table below.

+--------------------------+-------------------------------------------+
| Telescope                | HRT: Gregorian reflector                  |
|                          |                                           |
|                          | -   250 mm diameter                       |
|                          |                                           |
|                          | ```{=html}                                |
|                          | <!-- -->                                  |
|                          | ```                                       |
|                          | -   effective focal length 4250 mm        |
|                          |                                           |
|                          | FDT: zoom lens                            |
|                          |                                           |
|                          | \- aperture: 20 -- 50 mm                  |
|                          |                                           |
|                          | \- effective focal length: 233 -- 888 mm  |
+--------------------------+-------------------------------------------+
| Spatial resolving        | 0.28 arcsec pixel size (40 km on Sun at   |
| element                  | 0.2 AU)                                   |
+--------------------------+-------------------------------------------+
| Spectrometer             | -   Fabry-Periot double etalon            |
|                          |                                           |
|                          | -   tuning by rotation                    |
+--------------------------+-------------------------------------------+
| Spectral resolution      | FWHM: \< 50 mÅ                            |
+--------------------------+-------------------------------------------+
| Field of View/Pointing   | 80x80 Mm^2^; pointing to anywhere on Sun  |
|                          | by tilting and decentering secondary      |
|                          | mirror.                                   |
+--------------------------+-------------------------------------------+
| Detector                 | CCD:                                      |
|                          |                                           |
|                          | \- 10 micron pixels                       |
|                          |                                           |
|                          | -   2048x2048 array                       |
|                          |                                           |
|                          | -   back thinned                          |
+--------------------------+-------------------------------------------+
| Stability                | Active image stabilisation system and low |
|                          | order wavefront correction.               |
+--------------------------+-------------------------------------------+
| Telemetry                | 20 kbit/s                                 |
+--------------------------+-------------------------------------------+
| Mass                     | 28 kg                                     |
+--------------------------+-------------------------------------------+
| Power                    | 28 W                                      |
+--------------------------+-------------------------------------------+
| Size                     | Optics box: 160 cm x 40 cm x 30 cm        |
+--------------------------+-------------------------------------------+

**8.3.3 EUV spectrometer (EUS**)

# 

**Scientific Driver**

Observations of the UV/EUV spectral range are critical for the
determination of plasma diagnostics in the solar atmosphere across the
broad temperature range from tens of thousands to several million K.
Analysis of the emission lines, mainly from trace elements in the Sun's
atmosphere, provides information on plasma density, temperature,
element/ion abundances, flow speeds and the structure and evolution of
atmospheric phenomena. Such information provides a foundation for
understanding the physics behind a large range of solar phenomena.

Current spacecraft instrumentation (SOHO) provides EUV spatial and
spectral resolving elements of order 2-3 arcsec and 0.1 Å, respectively,
and UV resolutions of 1 arcsec and 0.02 Å. There are no plans for EUV or
UV spectroscopic measurements from the NASA STEREO mission (2004
launch), the NASA Solar Dynamics Observatory (\~2005 launch) and the
Solar Probe (\~2007 launch). The only planned EUV spectrometer for a
future mission is the EIS instrument on Solar-B with 1 arcsec (750 km on
the Sun) and 0.01 Å resolving elements. However, this instrument is
tuned to active region observations with almost no transition region
capability.

The European solar physics community has a well established expertise in
solar EUV/UV spectroscopy as illustrated by the successful CDS and SUMER
instruments on SOHO.

**Instrument Concept**

Whilst recognising that an EUV spectrometer is an essential component
the Solar Orbiter, we must be aware that it must be compact and
light-weight, must not be too telemetry 'thirsty' and must be able to
cope with the thermal and particle environment of such an orbit.

A Ritchey-Chretien telescope design, feeding a spectrometer, has been
chosen to minimise the size of the EUS instrument. It has the
disadvantage of one extra reflection over, for example, a single-mirror
off axis paraboloid system, but the wavelength range selected is geared
to bright solar EUV lines.

The instrument structure would be made of carbon fibre, with silicon
carbide (SiC) optical components. Multilayers will be considered if the
final wavelength selection requires it.

*Basic Optical Layout*

In the current design, the primary and secondary mirrors are separated
by 900mm, with the focal plane 200 mm beyond the primary, making a total
telescope length of 1.1 m. The telescope diameter is 120 mm, making an
overall diameter of the telescope 'tube' of about 150 mm. The telescope
effective focal length is 3.7 m.

The slit assembly lies at the focal plane and beyond this is the
spectrometer, with a toroidal grating forming a focus at a 2-D detector.
The grating ruling spacing is yet to be decided but we take 4800 l/mm as
a guide for the current calculations. For the wavelength bands selected
(see table), and assuming a spectrometer magnification of 1.5, the
spectrometer will project 1.2 m beyond the focal plane with a 1.76 m
distance to the detector. The detector will lie 0.55 m off-axis,
adjacent to the primary mirror. Thus, without a spectrometer, the
instrument cross-section would be of order 15x15 cm^2^ but the off-axis
path of the spectrometer will ensure that the back-end of the instrument
is wider, making an overall boot-shaped instrument of maximum dimensions
of order 230 cm x 15 cm x 55 cm. The 55 cm dimension is dependent on the
choice of grating and the wavelength selection.

The secondary presents a portion of the Sun at the slit and can be
rotated to allow rastered images (i.e. exposures interlaced with
mechanism movements to build up images simultaneously in selected
wavelengths). Only some 20% of the Sun will be visible to the secondary.
The instrument can be re-pointed to acquire particular targets.

Calculations of the spot-size at the focus from ray-tracing codes have
been performed for various angles of the secondary. This confirms that
the concept of rastering with the secondary is workable for the
scientific requirements over the planned range. We plan a 9 micron
detector pixel size and a field of view of 34x34 arcmin.

The heritage of this instrument comes from the SOHO/CDS, SOHO/SUMER and
Solar-B/EIS projects.

*Resolution/Detector*

We size the instrument to an observing distance of 0.2 AU (perihelion).
A spatial resolving element of 0.1 arcsec represents 75 km on the Sun
from the Earth; the same resolution can be achieved at 0.5 arcsec at 0.2
AU. With regard to spectral resolution, our basic aim is to return
sufficient numbers of well separated spectral lines to allow thorough
spectroscopic studies.

A detector array of 4k x 4k 9 micron pixels is baselined. Thus, the EUS
has a spectral range of 40 Å at 0.01 Å/pixel. The same array will give a
spatial extent (vertical distance on the detector = slit length) of 0.5
arcsec x 4096 = 2048 arcsec = 34 arcmin. The solar diameter at 0.2 AU is
170 arcmin, i.e. we have a slit length of 0.2 of the solar diameter.
Thus, a pointing mechanism is required. For a given pointing location,
rastered imaging will be made up from movement of the secondary mirror
in only one direction.

We baseline an APS (Active Pixel Sensor) detector system. The SOHO CCD
detectors receive many \'cosmic ray\' particle hits, occasionally have
\'bleeding\' of charge for the brightest intensities (e.g. EIT flares
and LASCO planets) and there must be some degradation due to protons
causing damage to the silicon lattice (this creates \'traps\' that can
steal charge which can be transferred to other parts of the image).

In an APS device, the signal charge is sensed by an amplifier within the
pixel (each pixel has its own). Therefore, it avoids the charge transfer
problem. Pixels will still be hit, but the damage will not propagate
because of the transfer process used for CCDs.

This degradation problem has been experienced by Chandra, mainly because
the CCDs have a direct line of sight to space for protons and the damage
was done mainly by relatively low energy protons (\~100 keV) of which
there are many. In SOHO, there is a lot more shielding. However, this
effectively means that the degradation will be at a slower rate, not
that there will be no degradation. [For a long duration mission (7
years) heading into a more intense particle environment, we have to
consider this seriously.]{.underline} We expect a moore intense particle
environment because Solar Orbiter is more likely to encounter
accelerated protons from solar events.

Intense sources display bleeding on CCD devices. This is not the case
for the APS detector system. This alone could be a tremendous advantage.

The APS detectors are still silicon chip devices, it is just that they
incorporate individual pixel amplification and charge extraction. The
EUV sensitivity will be provided in the same way as with CCDs, in this
case with back-thinned devices. There is no mass overhead in using such
a system, relative to the CCD approach.

**Resources and requirements**

*Telemetry*

For EUS we baseline 20 kb/s telemetry. The instrument detector has
4096x4096 pixels. It is clear that data from individual exposures will
have to be selected carefully, as given by pre-defined observing
sequences -- i.e. only selected fractions of the full range would be
returned in any exposure. In addition, a number of compression schemes
should be available with a common ('lossless') compression of order 3-5
but other schemes up to a factor of 10 would be available.

In order to cope with the limited telemetry resources, a combination of
novel compression schemes, large dedicated instrument memory and careful
emission line selection will be applied. In the light of the successful
operation of the CDS instrument on SOHO -- where a similar selection and
compression consideration was applied - and the new opportunities given
by the Orbiter's location, it is anticipated that this instrument can
make significant advances with a telemetry rate of order 20 kbit/s.

*Stability/Pointing*

Given our plan to achieve 0.5 arcsec resolution elements, we must choose
one of the following options:

(i) do not include an image stabilisation system, assuming that the
    variations of the spacecraft stability occur on timescales much less
    than the exposure time of the spectrometer and thus any corrections
    could be done on the ground, or

(ii) include an image stabilisation system possibly making use of VIM
     limb sensor error signals to drive adjustments to the secondary
     mirror.

The EUS instrument would require an independent pointing system to
enable independent observation of any solar location.

*Thermal Environment*

At 1 AU the average solar intensity is 1,371 W/m^2^. During the nominal
phase, in each 149 day period, the spacecraft will encounter a range
from 2142 W/m^2^ (0.8 AU) to 34275 W/m^2^ (0.2 AU -- 25 times the value
at 1 AU). This presents a severe thermal challenge which we tackle in a
number of ways.

The SiC primary mirror 'sees' the full Sun. SiC optical components can
run hotter than traditional components, but the primary receives 310.6 W
at 0.2 AU. Gold-coating is used to reduce the absorption (to 0.2). We
run the primary at 70^o^C, and require a dedicated radiator of 0.11 m^2^
(33 cm x 33 cm).

Some 248.5 W are reflected from the primary. The secondary is sized such
that it catches only 20% of this, i.e. it only sees 20% of the solar
image and much of the image is reflected back out of the aperture. Thus,
49.7 W is received by the secondary mirror which will also be gold
coated. Thus it absorbs 9.9 W which will require dissipation via a small
radiator of area about 11 cm x 11 cm. We run the secondary at 83^o^C.

The reflected component, now 39.8 W continues toward the spectrometer
aperture in the primary. This beam can be reduced considerably by stops;
only about 1/100 of this incoming beam is required for each exposure.

Finally, the APS detector system should run cold and this would require
a cold finger to a further radiator.

For the current design we assume passive radiators to space.

With regard to the remaining structure, the design would include thermal
blankets and the front panel could include a solar shield. Heaters will
be required to maintain the correct temperature gradients throughout the
orbit.

*Instrument Pointing/Spacecraft Shield*

The EUS can view a 34 x 34 arcmin area of the Sun, which is about 170
arcmin across at 0.2 AU. Thus an independent pointing mechanism is
required which needs to cover the range about +/- 3 degrees. The
instrument has an aperture of 120 mm. This must match a viewing port in
the spacecraft +X-plate shield. To minimise the viewing port size, the
instrument will be mounted on six legs (like the SOHO/CDS instrument)
with the actuators at the back and the pivot near the front of the
instrument. Thus, we envisage a port size of approximately 130 mm. The
front shield of the EUS instrument will act as a secondary to the
spacecraft shield when the aperture edge and viewing port edge are not
aligned.

*Mass Breakdown*

A rough breakdown of the estimated mass is given here. We take the view
that this is a new generation instrument made of light-weight materials
which really ought to be significantly lighter than, for example, the
equivalent (100 kg) instruments on SOHO.

+--------------------------------------------------------+-------------+
| ## Component                                           | **Mass**    |
|                                                        |             |
|                                                        | **(kg)**    |
+--------------------------------------------------------+-------------+
|                                                        |             |
+--------------------------------------------------------+-------------+
| Primary mirror                                         | 0.5         |
+--------------------------------------------------------+-------------+
| Mirror support                                         | 0.3         |
+--------------------------------------------------------+-------------+
| Secondary mirror                                       | 0.1         |
+--------------------------------------------------------+-------------+
| Mirror scan mechanism                                  | 0.6         |
+--------------------------------------------------------+-------------+
| Slit assembly                                          | 0.3         |
+--------------------------------------------------------+-------------+
| Grating assembly                                       | 0.6         |
+--------------------------------------------------------+-------------+
| OPS                                                    | 1.5         |
+--------------------------------------------------------+-------------+
| Detector                                               | 1           |
+--------------------------------------------------------+-------------+
| Detector electronics.                                  | 1.5         |
+--------------------------------------------------------+-------------+
| Baffles                                                | 0.5         |
+--------------------------------------------------------+-------------+
| Structure                                              | 5.4         |
+--------------------------------------------------------+-------------+
| Thermal subsystem                                      | 3.5         |
+--------------------------------------------------------+-------------+
| Harness                                                | 1.2         |
+--------------------------------------------------------+-------------+
| Main electronics incl PSU                              | 6           |
+--------------------------------------------------------+-------------+
|                                                        |             |
+--------------------------------------------------------+-------------+
| **[GRAND TOTAL]{.underline}**                          | **[23]{.u   |
|                                                        | nderline}** |
+--------------------------------------------------------+-------------+

**Instrument Summary**

The basic design of the instrument is given in the table.

+--------------------------------+-------------------------------------+
| Telescope                      | Ritchey-Chretien type:              |
|                                |                                     |
|                                | -   120 mm diameter                 |
|                                |                                     |
|                                | -   900 mm primary to secondary     |
|                                |                                     |
|                                | -   1100 mm secondary to focus      |
|                                |                                     |
|                                | ```{=html}                          |
|                                | <!-- -->                            |
|                                | ```                                 |
|                                | -   effective focal length 3.7 m    |
|                                |                                     |
|                                | -   gold-coated SiC mirrors         |
+--------------------------------+-------------------------------------+
| Spectrometer                   | Normal Incidence:                   |
|                                |                                     |
|                                | -   4800 l/mm toroidal grating      |
|                                |                                     |
|                                | -   slit/grating distance 1200 mm   |
|                                |                                     |
|                                | -   grating/detector distance       |
|                                |     1760mm                          |
+--------------------------------+-------------------------------------+
| Spatial resolving element      | 0.5 arcsec (75 km on Sun at 0.2 AU) |
+--------------------------------+-------------------------------------+
| Spectral resolving element     | 0.01 Å/pixel (5 km/s)               |
+--------------------------------+-------------------------------------+
| Raster Mechanism               | Through motion of the secondary     |
+--------------------------------+-------------------------------------+
| Detector                       | Active Pixel Sensor (APS):          |
|                                |                                     |
|                                | -   9 micron pixels                 |
|                                |                                     |
|                                | -   4096x4096 array                 |
|                                |                                     |
|                                | -   back thinned                    |
+--------------------------------+-------------------------------------+
| Preliminary Wavelength         | -   First Order: 580-620 Å          |
| Selection                      |                                     |
|                                | -   Second Order: 290-310 Å         |
|                                |                                     |
|                                | (e.g. Fe XII 291, Si X 293, He I    |
|                                | 584, Si IX 296, O III 599, He II    |
|                                | 304, Mg X 610, Fe XIX 592).         |
+--------------------------------+-------------------------------------+
| Slit Length/Width              | 34 arcmin length; width \<9 micron  |
+--------------------------------+-------------------------------------+
| Field of View/Pointing         | Raster over 34 arcmin; pointing to  |
|                                | anywhere on Sun with independent    |
|                                | pointing mechanism.                 |
+--------------------------------+-------------------------------------+
| Telemetry                      | 20 kbit/s                           |
+--------------------------------+-------------------------------------+
| Mass                           | 25 kg                               |
+--------------------------------+-------------------------------------+
| Power                          | 30 W                                |
+--------------------------------+-------------------------------------+
| Size                           | Instrument = 230 cm x 15 cm x 55 cm |
|                                |                                     |
|                                | Electronics box = 20cm x 20cm x     |
|                                | 20cm                                |
+--------------------------------+-------------------------------------+
| Thermal                        | Mirror operating temperatures 70    |
|                                | and 83^o^C, both with radiators.    |
|                                | APS detector cooled with cold       |
|                                | finger and radiator. Possible       |
|                                | shield at front. Standard thermal   |
|                                | blankets with heaters for thermal   |
|                                | control.                            |
+--------------------------------+-------------------------------------+
| Operation                      | On board deferred command store in  |
|                                | which can be loaded time-tagged     |
|                                | observation sequences designed on   |
|                                | the ground.                         |
+--------------------------------+-------------------------------------+
| Stability                      | Active image stabilisation system   |
|                                | at secondary may be considered if   |
|                                | jitter is significant over common   |
|                                | exposure times.                     |
+--------------------------------+-------------------------------------+

**8.3.4. EUV and X-ray imager (EXI)**

**Scientific Drivers**

The EIT/SOHO and TRACE imagers have recently exemplified the power of
EUV imaging for the understanding of the solar corona at all different
spatial and temporal scales.

The Full-Sun Imager (FSI) EIT not only allowed to survey the general
state of the corona (growing and decaying of active regions, coronal
holes, ..) but also to pinpoint special fine features such as plumes or
bright points and to record impulsive events such as flare, prominence
eruption and the start of CMEs (EIT waves).

With TRACE, the limited field of view (FOV) allowed an increased spatial
resolution which revealed a wealth of new features such as the
filamentary structure of coronal loops and their (apparent) constant
cross-section which puzzles theoreticians. It also appears that the 17.1
nm images are more filamentary than the 19.5 nm images. Unfortunately,
for TRACE (as for EIT), images taken at different wavelengths are also
taken at different times.

In the quiet Sun, a lot of \"events\" of various sizes (bright points,
explosive events, jets, blinkers,... ) all evidence heating probably
related to magnetic reconnection. Their spatio-temporal behavior
(distribution functions exhibiting power laws) is interpreted in terms
of small-scale intermittency due to, e.g., MHD turbulence. As for the
reconnection, it appears to be pervasive whether the magnetic field
lines are closed (quiet Sun) or open (coronal holes). In the last case,
it probably takes place low in the atmosphere (Axford & McKenzie 1991,
Parker 1991).

The results obtained lead to new questions concerning the structuring of
the corona above the poles, the real basal dimension of coronal
structures, the rôle played by nanoflares in the heating of the quiet
solar corona. Moreover, SOHO results point at the rôle of an activated
network from where high frequency waves, emitted to the outward corona,
could resonate with protons and heavy ions and provide the momentum for
the solar wind.

Dramatic new results can be obtained from increasing the spatial
resolution by about one order of magnitude on one hand, from connecting
chromospheric/coronal phenomena with close *in-situ* measurements and
also from exploring the open field regions on the other hand.

This is where an EUV Imager on SO can fully exploit the unique
capabilities of the mission. Such an imager must provide EUV images with
an angular resolution increased by about one order of magnitude in order
to reveal the fine structure of most coronal features. To do so, it
benefits from the closeness of the Sun. It must also provide full disk
EUV images of the Sun in order to reveal the global structure of regions
inaccessible such as the \"far side\" of the Sun and the polar regions.

The European solar physics community has a well established expertise in
solar EUV/UV imaging as illustrated by the successful EIT instrument on
SOHO. Some studies concerning the thermal (radiative) load on the optics
have been shared with Europeans teams involved in the Remote Sensing
Package of the Solar Probe mission.

**Instrument Concept**

A unique imager faces two difficulties : 1/ the best spatial resolution
cannot be obtained with a full disk imager for various reasons, the less
not being the limited number of pixels on current detectors, a
limitation coinciding with the telemetry constraints; 2/ it must cope
with a heavy thermal (radiative) load when at perihelion (25 solar
constants on the first optical element if the field-of-view (fov) is not
reduced). (It also must cope with the particle environment of such an
orbit).

This consideration leads to a separation between a High Resolution Image
(HRI) and a Full Sun Imager (FSI) all being part of EXI. In this way,
one can find specific solutions for the thermal load and have realistic
detectors.

*HRI*

The driver of the design is the need to get rid of the unwanted solar
radiation which goes along the direct light path. A first step consists
in reducing the entrance aperture of the instrument to about 2 cm and to
fully use the available length of the Payload Module. In doing so, one
can cluster a few telescopes together, each one dedicated to one
wavelength, which eliminates any moving part, takes full use of the
aperture and allows for *simultaneous* observations in different
wavelengths.

Essentially one diminishes the FOV through the use of an entrance baffle
and one sets the FOV to 70 000 km (or 500 arcsec when at perihelion) for
a 2\*2 K detector. This sets the equivalent pixel size to 35 km on the
Sun (or 0.25 arcsec). The pixel size of the detector is baselined at 9
microns, so the equivalent focal length is 7.2m.

We baseline a set of three telescopes working at 13.3, 17.4, and 30.4
nm. These three wavelengths cover a very wide range of temperatures
(from 1.6 10^7^ down to 5 10^4^ K) and targets (from flares to the quiet
Sun). Note that the 13.3 band includes a very hot line (Fe XXIII)
visible only during flares and that the 17.4 band only selects the 17.4
Fe X line, excluding the cooler Fe IX line at 17.1.

The basic feature of each telescope relies upon the use of a long baffle
to reduce the FOV and the stray light, an entrance metallic filter and
an off-axis Gregory telescope with multilayers fitted to each
wavelength. The Gregory concept allows to put stops at the focal plane
and at the image of the entrance aperture; it leads to a major reduction
of straylight.

A carbon-carbon composite structure is planned.

Figure X shows the optical scheme of the HRI.

![](media/image9.wmf){width="6.525in" height="3.9631944444444445in"}

Figure X: HRI optical scheme : B = baffle, M1 = primary mirror, M2 =
secondary mirror, F = entrance filter, D = detector. The figure also
shows two possible implementations of the three telescopes (boxes on the
right-hand side).

2kx2k APS detectors (see EUS section in 8.3.X) are baselined. Diamond or
GaN new generation detectors (robust, solar-blind, ..) may be
interesting alternatives.

The electronics (including DPU) is common to the three detectors and to
the FSI telescope. Its volume should fit within the overall volume given
in the Table.

*Counts number*

For the 17.4 and 30.4 channels, the small aperture allows for a few
counts only per pixel per second. One can cope with this low number by
increasing the exposure time up to at least 10 s (as compatible with the
pointing stability) and/or on-board binning (at least a factor 2\*2).
(the counts number for the 13.3 channel can be much higher but only on
very active regions and flares).

*FSI*

The Full Sun Imager needs a wide field (about 5°) when at perihelion in
order to cover the full lower solar corona with about 50 % margin in
each direction. For a 2kx2k detector, this means a pixel size of 9
arcsec (or 1250 km on the Sun at perihelion). In order to solve the
issue of thermal (radiative) load on the entrance filter of the Imager,
different schemes have been envisaged. We found that the only safe
possibility is to move back the entrance filter as much as possible
(1800 mm) and as close as possible to the M1 mirror. Then the mirrors
combination (all off-axis) allows for an image quality of the order of
the pixel size. The entrance filter is crossed twice by the light. This
system can work e.g. in the 17.4 line with the appropriate multilayer
(Mo/Si). Another solution exists if one works above 120 nm (e.g. in the
Lyman alpha line at 121.6 nm). In any case, the dimensions of the
instrument are of the order of 2000 x 250 mm x 400 mm, i.e. shorter than
HRI but wider.

Figure 4 shows the optical scheme of the FSI.

![](media/image10.wmf){width="6.525in" height="3.9631944444444445in"}

Figure 4: FSI optical scheme : B = baffle, M1 = primary mirror, M2 =
secondary mirror, F = entrance filter, D = detector.

*Operations*

Observing sequences should be defined/designed well ahead of time and
include target coordinates, sequence of operating telescopes, along with
the transmission of a few sample data. These sequences would be uplinked
and stored on board in a deferred command store. The operations for a
defined period would be stored in this way and sequences would be run
autonomously. Planning will be done more on a monthly and weekly basis;
however, we think useful to have daily samples of the data (at least a
full-Sun image and an image in each wavelength, i.e. the equivalent of 4
Mbits) in order to check the acquisition of scientific data and possibly
to change the pointing or the sequence of observations.

**Resources and Requirements**

*HRI*

***The average rate of 20 kbit/s can be obtained by observing with the
three telescopes about one day during the \"nominal mission\". This
duration may be increased by on-board binning (e.g. 2\*2) and adopting
an observing scheme different from strictly simultaneous
observations.***

***One can also envisage a set of two basic oberving modes : high time
resolution (short) sequences followed by low time resolution (long)
sequences.***

For a typical exposure time (a few ten seconds), the bit rate for each
telescope is very high and it is assumed that we can compress data by a
factor 20. Actually, with a Rosetta-type compression scheme (Langevin,
2000), tests performed on solar images seem to indicate that rates as
high as 50 could be achieved.

*Mass budget*

The overall mass budget is given in the Table below.

  ------------------------------------------------ ----------------------
  **Component**                                    **Mass (kg)**

  Structure (including the baffle) :               6.3

  Optics (including mounts) :                      0.2

  Detector (including electronics and shielding) : 0.5

  Miscellaneous (harness, ..)                      0.1

  Margin                                           0.8

                                                   

  Total                                            7.9

  Total 3 telescopes                               23.7

  Pointing mechanism                               2\.

  Electronics shared between the 4 telescopes      1.5

                                                   

  Grand total for the HRI                          27.2
  ------------------------------------------------ ----------------------

*Power*

The power requirements concern the detector feeding (it should be
remembered that voltage conversion is not lossless), electronics and the
pointing actuator (common to the three telescopes). For a given
telescope, we anticipate about permanently 4 W for the cycling \[power
on, integration, dump to memory\]. With margin and power for pointing,
we reach no less than 5 W. If the three telescopes are working
simultaneously, they will consume about 15 W for basic functions (above)
to which one should add about 2 W for data compression and dump by the
common DPU.

Non-simultaneous (sequential) observations of course lead to a smaller
power consumption.

*Stability/Pointing*

Given our plan to achieve 0.25 arcsec resolution elements and the fact
that exposure times are of the order of 10 seconds, an image
stabilisation system should be included, possibly making use of the VIM
limb sensor error signals to drive adjustments to the secondary mirror.

The three telescopes cluster has its own pointing mechanism since the
limitation of the FOV prevents any independent internal pointing. The
pointing is performed by a joint displacement of the three telescopes.
We can mention that the S/C having an offset capability of about 10 °,
much larger than what is required (2.5 ° at perihelion), it is well
possible to point a solar target with the whole S/C.

*FSI*

An average rate of 1 kbit/s is sufficient to transmit one full sun image
every 2400 s (about every 40 minutes) with compression.

The mass is of the order of 13 kg (weight of the structure : 8 kg;
optics and mountings : 4 kg; detector and associated electronics : 1
kg).

The power needed is only related to the detector (4 W) since the
electronics is common with HRI.

The FSI also relies on the pointing stability of the PLM but with
constraints much less stringent than for the HRI.

**Instrument Summary**

The characteristics of HRI are summarized in the following table.

+-----------------------+----------------------------------------------+
| 3 Telescopes          | Off-axis Gregory; 20 mm diam.                |
+-----------------------+----------------------------------------------+
| Spatial resolving     | 0.25 arcsec pixel size (35 km on Sun at      |
| element               | perihelion)                                  |
+-----------------------+----------------------------------------------+
| Pointing              | Common pointing mechanism                    |
+-----------------------+----------------------------------------------+
| Detector              | Baselined on a 9 microns, 2000\*2000         |
+-----------------------+----------------------------------------------+
| Field of View         | 70 000 km at perihelion                      |
+-----------------------+----------------------------------------------+
| Peak Telemetry        | 720 kbit/s (with compression 20)             |
|                       |                                              |
| Average Telemetry     | \>20 kbit/s                                  |
+-----------------------+----------------------------------------------+
| Mass                  | 24 kg                                        |
+-----------------------+----------------------------------------------+
| Power                 | 15 W (essentially detector; shared           |
|                       | electronics)                                 |
+-----------------------+----------------------------------------------+
| Size                  | Instrument = 2500 mm x 150 mm x 240 mm       |
+-----------------------+----------------------------------------------+
| Thermal               | Operating temperature 20^o^C with passive    |
|                       | cooling; 140 cm long baffle at the entrance  |
|                       | of each telescope                            |
+-----------------------+----------------------------------------------+
| Stability             | We assume a pointing stability of about 0.2  |
|                       | arcsec during a typical exposure time of 10s |
+-----------------------+----------------------------------------------+

The characteristics of FSI are summarized in the following table.

+------------------------+---------------------------------------------+
| 3 Telescopes           | Off-axis Gregory; 20 mm diam.               |
+------------------------+---------------------------------------------+
| Spatial resolving      | 9 arcsec (1300 km on Sun at perihelion)     |
| element                |                                             |
+------------------------+---------------------------------------------+
| Pointing               | Pointed at Sun center                       |
+------------------------+---------------------------------------------+
| Detector               | Baselined on a 9 microns, 2000\*2000        |
|                        | detector                                    |
+------------------------+---------------------------------------------+
| Field of View          | 4\*4 solar radii at perihelion              |
+------------------------+---------------------------------------------+
| Peak Telemetry         | 240 kbit/s (with compression 20)            |
|                        |                                             |
| Average Telemetry      | 1 kbit/s                                    |
+------------------------+---------------------------------------------+
| Mass                   | 13 kg                                       |
+------------------------+---------------------------------------------+
| Power                  | 4 W (essentially detector; electronics      |
|                        | shared with HRI)                            |
+------------------------+---------------------------------------------+
| Size                   | Instrument = 2000 mm x 250 mm x 400 mm      |
+------------------------+---------------------------------------------+
| Thermal                | Operating temperature 20^o^C with passive   |
|                        | cooling; 600 mm long baffle at the entrance |
|                        | of the telescope                            |
+------------------------+---------------------------------------------+
| Stability              | The constraint is defined by HRI.           |
+------------------------+---------------------------------------------+

Axford W.I. & McKenzie, J.F., 1991, Solar Wind 7, 1

Langevin, Y., 2000, S.P.I.E., submitted

Parker E. N., 1991, ApJ, 372, 719

**8.3.5. Ultraviolet and Visible-light Coronagraph** **(UVC)**

**Scientific Drivers**

The Solar and Heliospheric Observatory (SOHO) has broadened our ability
to investigate the outer solar corona by gaining access to an unexplored
electromagnetic domain, the ultraviolet, and to a much wider field of
view. This was achieved by flying a set of two coronagraphs, the
Ultraviolet Coronagraph Spectrometer (UVCS) and the Large Angle
Spectroscopic Coronagraph (LASCO).

The Ultraviolet and Visible-light Coronagraph (UVC) proposed for the
Solar Orbiter is designed to combine and extend the imaging capabilities
of the SOHO coronagraphs. This instrument can at the same time image the
visible and ultraviolet coronal emissions and diagnose with
unprecedented temporal coverage and space resolution the full solar
corona.

UVC can provide monochromatic images in the neutral hydrogen HI Ly alpha
1216 Å and single-inonized helium HeII Ly alpha 304 Å lines and measure
the polarized brightness (pB) of the K-corona. These observational
quantities are sufficient to diagnose crucial physical parameters, such
as the electron density, the outflow velocity through Doppler dimming
techniques, and the abundance of the major components of the solar
atmosphere.

Corotation, during the helio-synchronous phase of the orbit, will freeze
for many days coronal structures in the plane of the sky. This condition
will be most favourable to investigate the evolution of the magnetic
configuration of streamers in order to test the hypothesis of magnetic
reconnection as one of the main processes leading to the formation of
the slow solar wind.

The out-of-ecliptic vantage point will allow for the first time a unique
view of the plasma distribution and solar wind expansion in the coronal
low-latitude/equatorial belt. Therefore it will be possible to measure
the longitudinal extent of coronal streamers and coronal mass ejections.
These parameters, that at present are unknown, are essential to
determining the magnetic flux carried by plasmoids and coronal mass
ejections in the heliosphere.

UVC will allow the first determination of the absolute abundance (i.e.,
relative to hydrogen) of helium in corona. Helium, the second largest
contributor to the density of coronal plasma, is important for the
dynamics of solar wind, and it may act as a regulator to maintain a
nearly constant solar wind mass flux.

UVC will also determine the differential outflow speed of the major
components (H, He) of the solar wind and discriminate the mechanisms of
solar wind acceleration through different mass-charge ions observations.

The European solar community has a well established expertise in visible
light and UV coronagraphy as illustrated by the successful LASCO and
UVCS instruments on SOHO.

**Instrument** **Concept**

UVC is an externally o**c**culted telescope designed for narrow-band
imaging of the EUV corona in the He II 304 Å and H I 1216 Å lines, and
for broad-band polarization imaging of the visible K-corona, in an
annular field of view between 1.2 and 3.5 solar radii, when the Solar
Orbiter perihelion is 0.21 AU. When it moves to 0.3 AU, at the end of
the nominal mission the UVC field-of-view is within 1.8 and 5.3 solar
radii.

The telescope optical configuration is an off-axis Gregorian. The EUV Lα
lines are separated with multi-layer mirror coatings and EUV
transmission filters. These mirrors with coating optimized for 304 Å
still have a good reflectivity at 1216 Å and in the visible. The visible
light channel includes an achromatic polarimeter, based on
electro-optically modulated liquid crystals.

*Optical Design*

Telescope: The external occulter ensures both thermal protection for the
optics and better stray-light rejection. A Gregorian telescope design
has been chosen because it gives real images of the external occulter
and the edges of the telescope primary mirror. Field and Lyot stops will
be used to reduce stray light and diffracted radiation from the
telescope primary mirror. In addition, use of a annular secondary mirror
and light trap captures and removes diffracted light from the external
occulting disk and its supports. A light trap behind the primary mirror
ensures that only the light reflected by the mirror enters the filters /
detectors assembly structure. The direct light from the solar disk is
mostly dumped into space, though a portion of it may be used for
radiometery and / or coarse imaging of the solar disk in the three
wavelength bandpasses.

![](media/image11.png){width="2.841666666666667in"
height="2.975in"}![](media/image12.png){width="2.8409722222222222in"
height="2.9756944444444446in"}

Fig.xx: Two views of the optical layout of UVC.

UV and Visible-Light Channels: The instrument structure is made of
carbon fiber, with Zerodur optical components. Zerodur is used for the
mirrors\' substrate because of its extremely small thermal expansion
coefficient. Optics made with this material is also easy to polish into
non-spherical figures with sub-angstrom rms surface roughnesses.

The basic specifications and drivers of the instrument are summarized in
the following table (Table I).

*Structure*

The instrument structure is of carbon fiber/cyanate or carbon-carbon
composite, designed for zero coefficient of thermal expansion in the
axial direction. It is in the form of a rigid rectangular baseplate to
which are bonded appropriately placed brackets (also of composite) for
the support of the optical and detector elements. The focus and
orientation of each element is set and maintained by a precision spacer
located between the element and its support bracket. The optical system
is surrounded by a lightweight rectangular composite enclosure and
baffle structure that attaches to and further strengthens the baseplate.
Removable covers on the top of the enclosure allow for installation and
servicing of the optical and detector elements. An external door
mechanism shields the telescope aperture when the instrument is not
observing.

*Thermal*

The thermal control approach is based on shielding the instrument from
most direct solar radiation, rejecting light from the solar disk that
enters the annular entrance aperture, and an opto-mechanical design that
is insensitive to bulk temperature changes. The use of a carbon fiber
composite structural system allows the optical system to be almost
completely athermalized. Moreover, the use of a low secondary
magnification factor in the telescope (i.e.3:1) makes the optical system
less sensitive to dimensional changes in comparison with more
traditional space instruments such as TRACE. Preliminary estimates
suggest that the allowable temperature range of the opto-mechanical
system may be as much as +/- 50 C. Local temperature control will be
needed for the liquid crystal polarization module, which has a smaller
allowable temperature range. Thus we consider it prudent to include a
small power allocation for thermal control heaters (e.g. 1 watt) pending
further analysis.

The occulter is the element of the coronagraph, which reaches the
highest temperature since it is directly exposed to the Sun,
particularly when the Solar Orbiter is near perihelion. The occulter has
an exiguous conductive link (represented by the thin rods of the
support) with the rest of the structure. The occulter and supporting
links are made of titanium and are arranged in a conical geometry such
that thermally induced strain is absorbed without de-centering the
occulter disk.

The temperature achieved by the occulter has been estimated by assuming
a purely radiative heat transfer and an external coating of the same
type as that utilised for the sunshield of the payload module.

The thermal flux incident on the occulter is 31 kW/m^2^ at 0.21 AU and
5.5 kW/m^2^ at 0.5 AU. The temperature difference experienced by the
occulter along the observation phase is at maximum 266.5 C. A small
reduction of the temperature will be caused by the transport of part of
the heat through the occulter supports.

*Detectors*

There are two detectors, optimized for the Visible region (4500-6000 Å)
and for the Extreme UV (304-1216 Å). Both detectors have 15 micron
pixels and an array size of 4096x4096. The baseline visible detector is
an Active Pixel Sensor (APS). This detector system is based on CMOS
technology, which is attractive for high radiation dose environment. In
addition, the APS architecture has the potential to allow on-chip
differencing of the polarimetric signals, yielding substantial
improvements in signal-to-noise performance in comparison with
conventional CCD detectors.

The baseline EUV detector is a Charge Injection Device (CID) array. This
detection system is highly resistant to space radiation. It can easily
withstand more than 100 times the lethal dose of a typical CCD and
radiation-hardened versions can tolerate doses in excess of 1 MegaRad
(Si total dosage). CID detectors have randomly addressable pixels,
permitting the dynamically programmable read-out of individual pixels
and sub-arrays. This feature will be used in UVC for limiting the
readout area to the solar corona (excluding the occulted portion field)
and for defining \"Regions of Interest\" to be observed with higher time
resolution. EUV sensitivity will be provided by a Micro-Channel Plate
image intensifier coupled to the CID array with fiber-optics. This
approach results in a detector that is blind in the visible,
substantially improving the stray light rejection performance of the EUV
channel.

We note that both detector approaches are based to some extent on
anticipated developments in rapidly evolving technologies (e.g. array
sizes, APS noise performance). Heavily shielded CCD detectors are
suitable alternatives in either case.

**Resources and Requirements**

*Data Rates and Volume*

The primary driver of the data rate is UVC efficiency in the three
channels. In the two EUV lines, the coronal signal is weaker and the
instrument efficiency lower than in the visible band. Therefore, longer
exposure times will be required for the EUV coronal observations. On the
basis of the count-rate for H I and He II Lyman-alpha emission estimated
from a coronal hole (i.e., worst case), we assume an average exposure
time of 600 sec. A single image (4096 x 4096) pxl^2^ with 16 bit (b) per
pixel takes about 270 Mbit of memory. Assuming that only 50% of the
image is used (the rest is occulted disk, extreme corners of the square
matrix, etc.), and that a factor of 3 lossless compression can be
achieved, we expect a data rate of 0.07 Mbps per detector. Since the
Visible-light detector works simultaneously with the H I Lyman-alpha
(therefore, about 50% of the time), the data volume accumulated per
orbit is about 260 Gbit. A factor 10 image compression is needed to
reach a data rate of about 5 kbps. This level of compression can be
achieved with an acceptable loss of information with schemes such as,
for instance, the \`\`Adaptive Discrete Cosine Transform\'\' (ADCT). The
ADCT technique is being used in the Large-Angle Spectroscopic
Coronagraph (LASCO) on SOHO.

*Mass*

The UVC mass breakdown is given in the table.

+--------------------------------------------------------+------------+---+
| ## Component                                           | **Mass**   |   |
|                                                        |            |   |
|                                                        | **(kg)**   |   |
+--------------------------------------------------------+------------+---+
| ## Structure                                           | 6          |   |
+--------------------------------------------------------+------------+---+
| ## Optics                                              | 2          |   |
+--------------------------------------------------------+------------+---+
| ## Mechanisms                                          | 1          |   |
+--------------------------------------------------------+------------+---+
| ## Detectors (2)                                       | 2          |   |
+--------------------------------------------------------+------------+---+
| ## CPU & Interface                                     | 2          |   |
+--------------------------------------------------------+------------+---+
| ## DC/DC converter                                     | 1          |   |
+--------------------------------------------------------+------------+---+
| ## ADC                                                 | 1          |   |
+--------------------------------------------------------+------------+---+
| ## Data compressor                                     | 1          |   |
+--------------------------------------------------------+------------+---+
| ## Motor drive                                         | 2          |   |
+--------------------------------------------------------+------------+---+
| ##                                                     |            |   |
+--------------------------------------------------------+------------+---+
| **TOTAL instrument**                                   | **18**     |   |
+--------------------------------------------------------+------------+---+
|                                                        |            |   |
+--------------------------------------------------------+------------+---+

**Instrument Summary**

+-------------------+--------------------------------------------------+
| Telescope         | Off-axis Gregorian                               |
+-------------------+--------------------------------------------------+
| Occultation       | External occulter: 60 mm diam.                   |
+-------------------+--------------------------------------------------+
| Optics            | Primary mirror: 33 mm diam.                      |
|                   |                                                  |
|                   | Secondary mirror: 94 mm diam.                    |
|                   |                                                  |
|                   | (multilayer coated mirrors)                      |
+-------------------+--------------------------------------------------+
| Eff. Focal Length | 490 mm                                           |
+-------------------+--------------------------------------------------+
| Field of View     | Annular, Sun-centered. Coverage:                 |
|                   |                                                  |
|                   | 5.  R at 0.21 A.U.                               |
|                   |                                                  |
|                   | 1.8-5.3 R at 0.30 A.U.                           |
+-------------------+--------------------------------------------------+
| Spatial           | UV (304 Å, 1216Å): 8 arcsec at 1.2 R,            |
| Resolution        |                                                  |
|                   | 40 arcsec at 3.5 R                               |
|                   |                                                  |
|                   | Visible(5000 Å): 40 arcse entire FOV             |
+-------------------+--------------------------------------------------+
| Stray-light       | \< 10^-8^ (visible light); \< 10^-7^ (304 Å,     |
|                   | 1216 Å)                                          |
+-------------------+--------------------------------------------------+
| Wavel. Bandpasses | 1\) He II λ (304 ± 20) Å                         |
|                   |                                                  |
|                   | 2\) H I λ (1216 ± 100) Å                         |
|                   |                                                  |
|                   | 3\) Visible λ (4500 - 6000) Å                    |
+-------------------+--------------------------------------------------+
| Detectors         | Visible-light: Active Pixel Sensor (APS)         |
|                   |                                                  |
|                   | EUV (304 Å, 1216 Å): Charge Injection Device     |
|                   | (CID)                                            |
|                   |                                                  |
|                   | (EUV intensified); 4096 x 4096 array, 15 micron  |
|                   | pixels                                           |
+-------------------+--------------------------------------------------+
| Data Rate         | 10 kb/s                                          |
+-------------------+--------------------------------------------------+
| Phys. Dimensions  | (80 x 30 x 30) cm^3^                             |
+-------------------+--------------------------------------------------+
| Mass (kg)         | 22 Total                                         |
+-------------------+--------------------------------------------------+
| Power (W)         | 35 (Operational mode, peak); 25 (Stand-by mode)  |
+-------------------+--------------------------------------------------+

**8.3.6 Neutron and γ-ray detector (NGD)**

The Solar Orbiter provides a unique opportunity to measure many, and
also low-energy (few MeV), neutrons in the perihelion passages. Although
the simultaneously produced γ rays have been measured frequently on
other spacecraft, co-ordinated measurements are possible in
scintillation detectors via pulse-height and shape analysis of the
signals. Proton-recoil spectrometers or inorganic scintillators allow
neutrons to be detected in the energy range between 0.5 and 100 MeV and
γ rays between 0.1 and 10 MeV, with a large dynamic range, energy
resolution of a few percent, and time resolution in the s range. A
serious problem is the large mass of several kg of conventional
detectors. Mass reduction to less than about 1 kg is technologically
feasible (e.g. with a fission detector as used on Pioneers 10/11). For
example, McKibben et al. (1997) have suggested a low-mass design based
on detection of neutron-induced fission of heavy nuclei. Their
instrument will allow measurements down to very low neutron fluxes for a
mass of 1-2 kg and power less than 2 W.

**8.3.7** **Radiometer** **(RAD)**

**Scientific Drivers**

The irradiance of the Sun (i.e. its brightness as measured above the
Earth's atmosphere) is known to vary by 0.1% over the solar cycle. There
is also some evidence for a longer term, secular variation. In spite of
its small magnitude, the irradiance variation is a potential cause for
climate change. Two basic questions need to be answered before we can
reach an understanding of the causes of this variability:

-   How does the solar luminosity (i.e. the radiation escaping in all
    directions) vary? Does it change or is a brightening at the equator
    compensated by a darkening over the poles?

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-   Why is the irradiance variability of the Sun a factor of three
    smaller than that of Sun-like stars?

By carrying out irradiance studies from out of the ecliptic and on the
side away from the Earth, the Solar Orbiter will provide a unique
opportunity to answer these questions. For example, one major difference
between the Sun and Sun-like stars is that whereas we observe the Sun
nearly equator on, stars are on average seen from a latitude of
approximately 30^o^. The Solar Orbiter will reach this latitude in the
course of the extended mission and test whether this inclination effect
really is responsible for the Sun\'s anomalous behaviour.

The mission profile of the Solar Orbiter is ideally suited for studies
to improve our understanding of the angular redistribution of the
radiance in active regions of the photosphere. This can be done by
observing e.g. the influence of faculae on TSI from different angles:
from near the equator (from the Earth-bound mission) and from higher
latitudes from the Solar Orbiter. Similar studies can also be done and
have been done from one platform (e.g. SOHO) using the fact that the Sun
rotates, which allows to view the same feature from different angles at
different times; the uncertainty, however, is limited by the evolution
of the active regions themselves which are known to change on time
scales of days at the level of up to 0.1%. From our present
understanding, an uncertainty of about 1 part 10000 is needed for such
comparative studies. Thus, we have to ensure that IPSIR can be traced
back to 1 AU radiometers for all its operational modes and at all times
with such an uncertainty. This seems feasible, as shown below, but has
to be studied in detail and confirmed during a Phase A.

**Instrument Concept**

IPSIR is a new generation, room temperature radiometer with four
channels, which can be operated independently or in groups of two for
internal comparisons. Each channel has its shutter and an independent
cover for protection. IPSIR has electrically calibrated cavity
receivers. Two cavities are operated together with one closed and the
other with cyclic shutter phases of 50s open and 50s closed durations.
Phase sensitive detection at the fundamental shutter frequency is used
to analyze the signal. This operational mode has many advantages over
the normal "active cavity" operation in e.g. SOHO/VIRGO, which relies on
reaching a steady state before measurements can be taken (for a
description of this novel way of operating radiometers see e.g. George
M. Laurence *et al.* , 2000, "The Total Solar Irradiance Monitor: TIM",
*Metrologia,* **37**, in press).

In order to cope with the large change in solar irradiance from 1 to 0.2
AU three different apertures with 8, 4 and 2 mm diameter will be used at
different phases of the mission. A set of four of these apertures is
mounted on an aperture wheel which allows the selection of the same
aperture size for each of the four cavities. These determine the
receiver area (the m^2^ of the Wm^-2^ ) and about 80mm back from these
apertures are the view-limiting apertures of 10mm diameter in front of
each conical-cylindrical cavity (12mm diameter and about 38mm long)
yielding a full-angle field of view of 7^0^ and a slope angle of 0.7^0^,
2.1^0^ and 2.9^0^ for 8, 4 and 2mm apertures. This reversed arrangement
of the field-of-view apertures has the advantage of less energy in the
radiometer and a simpler treatment of the diffraction correction. At 1
AU with the 8mm aperture an area of 0.5cm^2^ is illuminated with about
70mW with a radial fall-off to zero within 0.4mm, whereas at 0.2 AU (2mm
aperture) a central area of 0.03cm^2^ is illuminated with about 110mW
and the radial fall-off to zero is within 1.7mm. Thus, the spatial
irradiance distribution within the cavity changes substantially for the
two cases and thus the non-equivalence between the homogenous electrical
heating and the inhomogeneous radiative heating will probably be the
most important limiting factor of the uncertainty of the radiometer.

The operation will be mostly digital and done by software (S/W) in the
instrument\'s computer. There are essentially only two analog circuits:
(i) the amplifier of the thermistor bridge measuring the temperature
difference along the thermal impedance and between two cavities (one
measuring irradiance the other as thermal reference) and (ii) the highly
stable voltage source for the electrical power fed to the cavity. The
bridge is powered by a sine wave at about 100Hz produced by S/W and its
output is A-D converted and synchronously demodulated by S/W. The
demodulated signal is used to PID control the heater power of the
shuttered cavity to compensate for the extra radiant power during the
open periods. The heater power is deduced from the highly stable voltage
source which is pulse-width modulated by the output of the PID
controller. This power signal is frequency analyzed and the amplitudes
and phases of the shutter fundamental frequency and the first few
harmonics of it (up to about the 8th) are transmitted to ground. Only
the in-phase signal at the shutter fundamental is needed to calculate
the irradiance, the harmonics are for the in-flight characterization of
the temporal behavior of the detectors. Some of the correction factors,
needed for the final evaluation of the irradiance can be determined
and/or checked in flight.

# **Instrument Summary**

-   *4-Channel Radiometer* with an uncertainty of \<500ppm (goal)

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-   *Three precision aperture* diameters: 8, 4 and 2mm corresponding to
    areas of 0.5, 0.12 and 0.03cm^2^; the full-angle field of view is
    7^0^ and the slope angles are 0.7^0^, 2.1^0^ and 2.9^0^ for 8, 4 and
    2mm apertures

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-   *Volume:* 110x110x220mm with the 110x110mm face looking to the sun;
    in front there will be some (probably 2) radiation shields which
    will extend the length of the box by 40mm and the area of the front
    to 140x140mm.

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-   *Mass:* 4.0 kg including a contingency of 10%

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-   *Power:* total power consumption is estimated to 6.6W including a
    contingency of 10%. This estimate is still uncertain, as the
    computer type has not yet been selected.

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-   Normal *data rate* (measuring with one radiometer) is 750 b/s.
    During a backup operation (about once a week) 280 b/s have to added.
    During check-out a total of 3500 b/s will be needed; this mode is
    during commissioning and if problems have to be analyzed.

    **8.3.8 Other potential solar instruments *Still old prop.
    Version!***

A valid option to be considered in future studies of the Solar Orbiter
is the addition of three corona-related instruments:

-   A wide-angle white-light instrument with a FOV from 1.5 to 6 R~s~,
    equivalent to the LASCO-C2 on SOHO, would allow one to see extended
    coronal structures and evolving CMEs. It would cover the cusp
    regions of helmet streamers, where the slow solar wind is released
    and begins to accelerate. The highly successful LASCO-C3 instrument
    (FOV from 3 to 32 R~s~ at 1 AU) could serve as a design baseline.
    The required mass should not exceed 4 kg (assuming common electronic
    systems with the other imagers).

-   A heliospheric imager might extend the view of the heliosphere to
    almost a full hemisphere, from about 3 R~s~ (at 0.2 AU) to as much
    as the anti-solar direction. Thus, it would have almost the full
    Earth-Sun line within its FOV for major parts of the observation
    time, depending on the spacecraft attitude. This instrument would
    allow for the first time to observe the propagation of transient
    disturbances towards the Earth from various distances and aspect
    angles, in particular from outside the ecliptic plane. Preliminary
    design studies have demonstrated that such an instrument can
    probably be built for less than 2 kg. A crucial requirement would be
    that the hemisphere to be watched must be kept free of
    Sun-illuminated S/C appendages.

-   An X-ray instrument dedicated to the high-temperature portion of the
    coronal emission could excellently complement the Solar Orbiter
    payload. A novel configuration for an imaging spectroscopic
    instrument would operate in the wavelength range from 40 Å to 6 Å
    and detect emissions from plasma at temperatures between a few 10^6^
    K and a few 10^7^ K. Such an instrument working in extremely-grazing
    incidence will obtain very high spectral and spatial resolution. It
    would also be able to measure densities equal to or greater than
    10^9^ cm^-3^ (through density-sensitive lines in this
    high-temperature range).

    **Table 8.4: Solar Instrumentation (\*depth x width x length)**

+---------+------+----------+---------+------+----------+-----+-------+
| Name    | Ab   | Mea      | Specifi | Mass | Size∗    | Po  | Tele  |
|         | bre- | surement | cations |      |          | wer | metry |
|         |      |          |         | kg   | cm x cm  |     |       |
|         | via  |          |         |      | x cm     | W   | kb/s  |
|         | tion |          |         |      |          |     |       |
+---------+------+----------+---------+------+----------+-----+-------+
| Visibl  | VIM  | H        | Na D1   | 35   | 30 x 40  | 30  | 20    |
| e-light |      | igh-res. |         |      | x 120    |     |       |
| Imager  |      | disk     | Na D2   |      |          |     |       |
| and     |      |          |         |      |          |     |       |
|         |      | imaging  |         |      |          |     |       |
| Magne   |      | and      |         |      |          |     |       |
| tograph |      |          |         |      |          |     |       |
|         |      | pol      |         |      |          |     |       |
|         |      | arimetry |         |      |          |     |       |
+---------+------+----------+---------+------+----------+-----+-------+
| EUV     | EUS  | Imaging  | EUV     | 35   | 30 x 15  | 30  | 20    |
| Imager  |      | and      | e       |      | x 140    |     |       |
|         |      | dia      | mission |      |          |     |       |
| and     |      | gnostics | lines   |      |          |     |       |
|         |      | of       |         |      |          |     |       |
| Spect   |      |          |         |      |          |     |       |
| rometer |      | TR and   |         |      |          |     |       |
|         |      | corona   |         |      |          |     |       |
+---------+------+----------+---------+------+----------+-----+-------+
| X-ray / | EXI  | Corona   | He and  | 15   | 10 x 10  | 5   | 5     |
| EUV     |      | imaging  | Fe      |      | x 30     |     |       |
|         |      |          |         |      |          |     |       |
| Imager  |      |          | Ion     |      |          |     |       |
|         |      |          | lines   |      |          |     |       |
+---------+------+----------+---------+------+----------+-----+-------+
| Ultr    | UVC  | Imaging  | Coated  | 25   | 20 x 20  | 10  | 10    |
| aviolet |      | and      | mirror  |      | x 50     |     |       |
| and     |      | dia      | coro    |      |          |     |       |
| Visible |      | gnostics | nagraph |      |          |     |       |
| Light   |      | of the   |         |      |          |     |       |
| Coro    |      | corona   | CCD     |      |          |     |       |
| nagraph |      |          | d       |      |          |     |       |
|         |      |          | etector |      |          |     |       |
+---------+------+----------+---------+------+----------+-----+-------+
| Neutron | NGD  | Solar    | 0.5-100 | 1    | 10 x 10  | 2   | 0.5   |
| and     |      | neutrons | MeV     |      | x 10     |     |       |
|         |      |          |         |      |          |     |       |
| γ-ray   |      | gamma    | 0.1-10  |      |          |     |       |
| D       |      | rays     | MeV     |      |          |     |       |
| etector |      |          |         |      |          |     |       |
+---------+------+----------+---------+------+----------+-----+-------+
| Rad     | RAD  | Solar    | Visible | 4    | 15 x 15  | 5   | 0.5   |
| iometer |      | constant | light   |      | x 25     |     |       |
+---------+------+----------+---------+------+----------+-----+-------+