**8. Scientific Payload**

**8.1 Measurement requirements**

The Helios, Ulysses, Yohkoh, SOHO and TRACE missions may be used as a
baseline for designing the payload of the Solar Orbiter. However, this
mission will take a scientific spaceprobe into closer vicinity of the
sun than has ever before been reached by a man-made object. In
particular, the thermal load of almost 25 solar constants means an
unprecedented technical challenge. The instruments have to be designed
accordingly. They may dwell on developments implemented in many space
missions in the past, but the extreme thermal environment requires
substantial redesign. In order to make sure that the ambitious mission
goals of the Solar Orbiter can really be achieved, all individual
instruments were given extraordinary attention. Some of them were
designed to such detail as is usually not the case at this phase of
mission preparation. Also, we made sure that the required technology is
available within Europe.

The strawman payload discussed in this chapter contains sufficient
detail for system design to demonstrate its feasibility. However, it is
important to note that no pre-selection of certain instruments or
instrument principles is implied.

The payload 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 8.6 below): Solar wind
    plasma analyser measuring electrons and ions, magnetometer for DC
    and AC fields, particle detectors 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): X-ray/EUV
    full-Sun 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 -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)**

##### 

##### Scientific drivers

Three principal science goals require the inclusion of a solar wind
plasma analyser in the payload:

-   to provide observational constraints on kinetic plasma properties
    for a fundamental and detailed theoretical treatment of all aspects
    of coronal heating,

-   to investigate charge- and mass-dependent fractionation processes of
    the solar wind acceleration process in the inner corona,

-   to correlate comprehensive in-situ plasma analysis and compositional
    tracer diagnostics with space-based and ground-based optical
    observations of individual stream elements.

Furthermore, the SWA will

-   investigate in detail ^3^He and "unusual" charge states in
    CME-related flows and to correlate the observations with other
    in-situ particle and field data and with optical observations,

-   investigate the "recycling" of solar wind ions on dust grains in the
    distance range which has been located as the "inner source". Freshly
    produced pick-up ions from this inner source are specially suited as
    test particles for studying the dynamics of incorporation of these
    particles into the solar wind.

The SWA will measure separately the three-dimensional distribution
functions of the major solar wind constituents: protons, 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.

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).

**Instrument concept**

In view of the limited resources of mass, volume and telemetry allocated
to SWA, a compromise between sensitivity, mass/charge- and mass- and
time resolution has to be found. The SWA has to cover a large dynamic
range in ion fluxes. Since there is an enormous difference between the
proton fluxes at perihelion (typically 10^14^ m^-2^ s^-1^) and the
fluxes of relevant minor ion tracers at 1 AU (e.g. Fe^10+^ at typically
10^8^ m^-2^ s^-1^ etc.) it is suggested to implement three different
sensors:

1.  A Proton/Alpha-particle Sensor (PAS) with the principal aim to
    investigate the velocity distribution of the major ionic species at
    a time resolution equivalent to the ambient proton cyclotron
    frequency.

2.  An Electron Analysers System (EAS) consisting of two or three
    combined sensor heads covering a solid angle of almost 4and
    allowing the determination of the principal moments of the electron
    velocity distribution at a time resolution at least equivalent to
    the proton cyclotron period near perihelion.

3.  A Minor Ion Sensor (MIS) which allows the independent determination
    of the major charge states of oxygen and iron and a coarse mapping
    of the three-dimensional velocity distribution of some prominent
    minor species. Also, pick-up ions of various origins, such as
    weakly-ionised species (C^+^, N^+^, Ne^+^, Mg^+^, Si^+^, etc.), will
    be measured.

The PAS is designed to resolve the relevant ion-kinetic scales. It will
achieve a time resolution which is compatible with the characteristic
ion cyclotron periods near perihelion. The characteristic time scale is
about 100 ms. For the minor species (heavier than He) a time resolution
of 1 s is needed.

To fulfil the requirements of compositional diagnostics at least the
charge state distributions of oxygen and iron should be measured
routinely and with a temporal resolution of better than 10 s.

The instrument specifications are outlined in Table 8.1:

+-------+------------+------------+----------+------------+----------+
| **Sen | **P        | **Field of | **Energy | **Angular  | **Time   |
| sor** | articles** | View**     | reso     | re         | reso     |
|       |            |            | lution** | solution** | lution** |
+-------+------------+------------+----------+------------+----------+
| ###   | Protons    | Conical    | E/E = 10 | Azimuth    | 10 ms    |
| # PAS | and        | ±45^o^    |          | 30^o^      |          |
|       |            |            |          |            |          |
|       | particles  |            |          | Zonal      |          |
|       |            |            |          | ±5^o^     |          |
+-------+------------+------------+----------+------------+----------+
| # **  | Electrons  | 4          | E/E = 10 | Azimuth    | 10 ms    |
| EAS** |            |            |          | 30^o^      |          |
|       |            |            |          |            |          |
|       |            |            |          | Zonal      |          |
|       |            |            |          | ±5^o^     |          |
+-------+------------+------------+----------+------------+----------+
| #     | O, (^3^He, | Conical    |          |            | 1 s      |
|       | C, Ne, Mg, | ±45^o^    |          |            |          |
| # **  | Si), Fe    |            |          |            |          |
| MIS** |            |            |          |            |          |
+-------+------------+------------+----------+------------+----------+

Table 8.1 Specifications of the SWA sensors

The instrument package proposed for SWA is a further development of
instruments successfully flown on several missions in the past, e.g.,
SWICS on Ulysses, CELIAS/CTOF on SOHO, TRIPLET on INTERBALL. A possible
scheme for PAS based upon the design of TRIPLET is shown in Fig. 8.1.

######## 

######## 

######## Fig. 8.1 Basic schematic of the TRIPLET instrument as example for a possible PAS design. The points in the collimator layout correspond to pinhole channels. Each pinhole aperture gives one energy strip on the detector surface.

**Resources and requirements**

*Telemetry*

The transmission of velocity distributions of the major species
(protons, electrons and particles) drives the telemetry requirements of
SWA. In the normal mode of operation this information should be
compressed to transmit the principal moments of velocity distributions.
Periodically however, the transmission of the complete velocity
distributions of protons and particles with a (velocity) resolution of
10x10x10 pixels at a resolution limit of 5 b is needed. This requires 5
kb per spectrum corresponding, at a time resolution of 10 ms, to a
telemetry rate of 500 kb/s. This rate should be maintained during 1s,
i.e., typically 100 ion cyclotron periods to cover a growth time of
plasma instabilities. With an average telemetry rate of 5 kb/s a duty
cycle of the order of 1% can be achieved.

*Mass allocation and volume*

A mass allocation for MIS is 4 kg, 1 kg for PAS and 1 kg for EAS
(excluding the electronics) is foreseen. The volume allocation for MIS
is 20x20x20 cm^3^, and 10x10x10 cm^3^ for PAS and for each of the two
EAS sensors.

*Power*

The total power allocated to SWA is 5 W.

*Thermal and radiation requirements*

The instrument should be able to operate under the specified spacecraft
conditions. Special attention has to be devoted to the extreme solar UV
radiation flux. The design of the particle instruments should reject
this radiation to a sufficiently low level in order not to compromise
the detection of rare particles.

*Electrostatic cleanliness*

Special care has to be taken to ensure near-equipotential conditions in
the vicinity of the EAS entrance apertures.

**Instrument summary**

+------------------+-----------+-----------+-------------+------------+
| **Instrument**   | **Mass**  | **Power** | **Volume    | **Data     |
|                  |           |           | (cm^3^)**   | rate       |
|                  | **(kg)**  | **(W)**   |             | (bps)**    |
+------------------+-----------+-----------+-------------+------------+
| **PAS**          | 1         | 1.2       | 1000        | 2000       |
+------------------+-----------+-----------+-------------+------------+
| **EAS**          | 1         | 0.8       | 1000        | 2000       |
+------------------+-----------+-----------+-------------+------------+
| **MIS**          | 4         | 3         | 8000        | 1000       |
+------------------+-----------+-----------+-------------+------------+
| **Total SWA**    | 6         | 5         | 10000       | 5000       |
+------------------+-----------+-----------+-------------+------------+

**8.2.2 Radio and Plasma Waves Analyser (RPW)**

**Science drivers**

*In-situ* wave investigations on many spacecraft (e.g., Ulysses, Wind)
have amply demonstrated the crucial role of wave observations in the
study of a broad range of solar terrestrial phenomena. Wave observations
have been particularly useful in correlation with SOHO optical and EUV
observations. Typical waves instrumentation provides measurements of
both the electric field and magnetic field in a broad frequency band,
usually from a fraction of a Hertz up to several 10s of MHz, covering
characteristic frequencies in the solar corona and interplanetary
medium. Both electrostatic waves and electromagnetic waves can be
measured, leading to different diagnostics: electrostatic waves mostly
provide *in situ* information, in the close vicinity of the spacecraft,
while electromagnetic waves can provide an extensive remote sensing of
energetic phenomena in the solar corona and interplanetary medium.

Solar radio astronomy in a broad range of frequencies provides a unique
means to monitor, track and analyse energetic phenomena taking place in
the solar corona and interplanetary medium, in association with solar
activity. Radio emission is essentially generated through the
interaction of energetic electrons (produced in solar flares and shock
waves) with the ambient solar wind plasma. Most radiation mechanisms
produce radiation on a frequency close to the local plasma frequency
which is directly proportional to N~e~ , leading to the strong
frequency-distance dependence shown in Table 8.2.

  --------------------- -------------------------- -----------------------
                        electron density N~e~      radio frequency
                        (cm^-3^)                   

  low corona            \~ 10^8^ -- 10^10^         100 MHz -- 1 G Hz

  medium corona         \~ 10^6^ -- 10^8^          10 MHz -- 100 MHz

  \~ 10 R~o~            \~ 10^4^                   \~1 MHz

  \~ 40 R~o~            \~ 10^3^                   \~ 300 kHz

  \~ 1 AU               \~ 10                      \~ 30 kHz
  --------------------- -------------------------- -----------------------

Table 8.2. Radio frequency as function of electron density

Using proven techniques, the Solar Orbiter mission profile will allow a
radio and plasma wave package to face quite different and new
situations: access to waves and turbulence that occur much closer to the
Sun; access to viewing angles from well out of the ecliptic plane,
permitting for the first time studies of the north-south beaming of the
radio radiation in the solar corona.

**Instrument concept**

The Radio and Plasma Waves Analyser (RPW) comprises two main
sub-systems: the Plasma Waves Analyser (PWA) and the Radio Spectrometer
(RAS) covering respectively *in situ* measurements and remote sensing.
The two sub-systems share some of the sensors (essentially the low
frequency boom antenna) and have a common digital signal processing unit
and interface with the spacecraft.

The same receivers can be used to analyse the different types of waves
detected by different sensors, for instance electric antennas and
magnetic coils. There is a lot of experience in the design of the
sensors and receivers, as well as in the analysis techniques. Also all
aspects of electromagnetic cleanliness, are well understood and are no
longer a problem for modern spacecraft designs (e.g., ULYSSES,
WIND/POLAR, CASSINI, IMAGE, CLUSTER).

*Concept of the PWA*

The PWA will identify the various plasma waves and kinetic modes
comprising the high-frequency part of the fluctuation and turbulence
spectra. We provide below a brief description of the characteristics of
the plasma waves expected to be present at 0.2 AU.

The expected field strength may range between a few V/m and mV/m, and up
to about 1 V/m for the convection electric field. The magnetic field
strength is expected to vary between a few nT and T, 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).

From Helios observations one can estimate that a sensitivity equal to
10^‑6^nT/(Hz)^1/2^ will be required at 0.2 AU to identify without
ambiguity whether the observed waves are electromagnetic or
electrostatic.

At and above the electron plasma frequency intense electron plasma
oscillations and solar radio waves are expected. Helios observations
show that wave intensities associated with type III burst increase very
rapidly with decreasing radial distance from the Sun. Intense emission
could extend up to 10^-4^ nT/(Hz)^1/2^.

The measurement requirements necessary for the identification and
analysis of the various electromagnetic phenomena expected to be
measured at 0.2 AU are given in Table 8.3.

+-------------------+--------------+----------------+----------------+
| Phenomena         | Frequency    | minimum        | maximum        |
|                   | band         | amplitude      | amplitude      |
|                   |              |                |                |
|                   |              | (nT/Hz)        | (nT/Hz)        |
+-------------------+--------------+----------------+----------------+
| Dopplerised       | 1 Hz -- 10   | 10 ^--5^ at 10 | 10 ^--1^ at 10 |
| ELF/VLF waves and | kHz          | Hz             | Hz             |
| turbulence        |              |                |                |
+-------------------+--------------+----------------+----------------+
| Ion acoustic      | 10 kHz -- 10 | A few10 ^--7^  | 10 ^--4^ at    |
| waves and Radio   | MHz          | at 1 MHz       | 1MHz           |
| waves             |              |                |                |
+-------------------+--------------+----------------+----------------+

# 

# Table 8.3: Measurement requirements for the fluctuating magnetic field 

The PWA will cover a broad band in frequencies, extending from about 1
Hz into the MHz range. Resolving the vector components of the electric
and magnetic fields is scientifically highly desirable and required to
determine wave modes unambiguously. However, due to different
constraints of the mission, we propose to measure one electric field
component only, with a boom antenna mounted in the anti-sun direction,
in the shadow of the spacecraft, while the three components of the
fluctuating magnetic field can be easily measured with a 3-axial search
coil magnetometer arranged in a compact configuration and mounted on a
short boom, required for magnetic cleanliness reason, on the anti-sun
direction. This boom could be shared with the MAG as long as a minimum
distance between the two sensors is respected. There is a strong
heritage for those sensors mounted on three-axis stabilised spacecraft,
e.g., on Galileo, Cassini, STEREO. The PWA will also benefit from
studies and developments under progress for the Solar Probe.

The PWA will perform on-board processing of the data and deliver
selectable wave forms (Time Domain Sampler), to identify non-linear
coherent fluctuations and correlate them with measurements from the EPD.

*Concept of the RAS*

The RAS will measure the solar and interplanetary radio waves in the
frequency range from 100 kHz to 1 GHz, with a sweep period between 0.1 s
and 10 s and a high spectral resolution (f/f 0.07). The RAS will observe
plasma processes associated with energetic electrons from the low corona
up to about 0.5 AU. 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. Since radio radiation is generally beamed (beam widths
sometimes down to a few ten degrees) more or less along a radial
direction from the sun, this technique is particularly relevant for
different vantage points, for instance when the Solar Orbiter will
observe the far side of the sun. The time history provided by the
regular acquisition of the radio dynamic spectrum will help to trace the
development of the active region in a synthetic manner. Activity indices
could thus be produced.

The time resolution required to detect the rapidly varying solar bursts
vary with the radio frequency. Typically, the duration of a type III
burst (stream of energetic electrons) is t(s) = 220 / f(MHz). This
points to time resolutions of the order of 0.1 s or better for the high
frequencies and of 10 s for the low frequencies.

A single boom antenna can be used for the low frequency part of the band
(below, e.g., 20 MHz), where the sensitivity is essentially determined
by the sky background radiation. Above 20 MHz and particularly above
200 MHz, the background is determined by the receiver temperature and
gain is needed in order to detect solar radio bursts: a simple boom
cannot be used anymore. A broad-band antenna with significant gain is
thus required.

The RAS instrument could be divided into three sub-spectrometers, for
instance:

-   RAD1: 100 kHz - 16 MHz

-   RAD2: 16 MHz - 200 MHz

-   RAD3: 200 MHz - 1 GHz

Different designs are used on many space mission and have a lot of
heritage since the birth of space radio astronomy at long wavelengths.

Each one of the receivers will be fed by different antenna or search
coil systems.

*RPW instrument description*

An overall block diagram of the Solar Orbiter RPW instrument is shown on
Figure 8.2.

![](media/image2.png){width="6.522222222222222in"
height="4.049305555555556in"} Figure 8.2.: Block diagram of the Solar
Orbiter RPW instrument

Plasma waves and low frequency activity will be monitored by a receiver
covering the band from about 1 kHz up to 100 kHz. In order to allow for
a high time resolution, while minimising resources, it can be of the
type flown on WIND and projected for STEREO, using a digital filtering
technique.

The wave form analyser or Time Domain Sampler (TDS) is a broad band
waveform sampler with 4 channels (1 electric field component and 3
magnetic field components). The signal from the boom antenna and search
coil are sampled at a rate up to 10^6^ samples per second. Events are
detected on board and transmitted according to the available bitrate.
This is done using 'intelligent' algorithms which can return a quality
value for the events. A number of random events can be sent to the
ground without regard of their quality in order to provide unbiased
statistics. Real time triggers can be sent to the EPD for special
studies of wave-particle phenomena. The bitrate can be quite variable. A
compromise can be found in selecting the number of events that will be
transmitted per day.

The radio receivers can be of the classical super-heterodyne type.
Frequency synthesisers will allow for a maximum flexibility in the
choice of the observing frequencies (for instance to avoid polluted
frequencies on the spacecraft). Each sub-receiver can consist of up to
256 selectable channels. Only a selection of these channels could be
transmitted to the telemetry stream.

The DPU will command the different subsystems and interface with the
spacecraft. The variety of instruments call for a packetised telemetry
system, that can be managed by the DPU itself, as is done for similar
packages on the WIND spacecraft (fixed format telemetry system) and
STEREO. The DPU will interface with other instruments on board the
spacecraft, particularly the Energetic Particle Detector. The DPU could
be shared with some other instruments of the of the payload, e.g., with
MAG.

Tables 8.3a, 8.3b, 8.3c provide the main characteristics of the sensors
and analysers comprised in the RPW investigation. Numbers are indicative
and can vary according to design options

+------+--------+---------+--------+---------+----------+----------+
|      | Ma     | fr      | sensi  | size    | mass (g) | power    |
|      | gnetic | equency | tivity |         |          |          |
|      | s      | band    |        |         |          |          |
|      | ensors |         | (      |         |          |          |
|      |        |         | nT/Hz) |         |          |          |
+------+--------+---------+--------+---------+----------+----------+
| B1   | Search | 1Hz--10 | 10     | 3 rods, | 500      | 200 mW   |
|      | coils  | kHz     | ^--5^  |         |          |          |
|      |        |         | at 10  | 18 cm   |          |          |
|      |        |         | Hz     | each    |          |          |
+------+--------+---------+--------+---------+----------+----------+
| B2   | Ma     | 10      | 1      | 20 cm   | 600      | 200 mW   |
|      | gnetic | kHz--10 | 0^--6^ | d       |          |          |
|      | loop   | MHz     | at     | iameter |          |          |
|      |        |         | 1MHz   |         |          |          |
|      | (opt   |         |        |         |          |          |
|      | ional) |         |        |         |          |          |
+------+--------+---------+--------+---------+----------+----------+
| *To  |        |         |        |         | 1100     |          |
| tal* |        |         |        |         |          |          |
+------+--------+---------+--------+---------+----------+----------+

Table 8.3.a: Main characteristics of the RPW magnetic sensors

  --------- -------------- ------------- ------------- ----------- -----------
            Electric       frequency     size          mass (g)    power
            sensors        band                                    

  E1        radio boom     1 Hz -- 20    5 m boom      1300        N/A
                           MHz                                     

  E2        active dipole  20 MHz -- 200 \~ 1m x 2cm x 800 (TBC)   300 mW
                           MHz           2 cm                      

  E3        Log-periodic   200 MHz--1    \~ 1m x 1m x  1500 (TBC)  N/A
            antenna        GHz           1cm                       

  *Total*                                              2600        
  --------- -------------- ------------- ------------- ----------- -----------

Table 8.3.b: Main characteristics of the RPW electric sensors

+--------+-----------------+--------+---------------+----+-----+-----+
| an     | frequency band  | Sen    | sensitivity   | ma | B   | Po  |
| alyser |                 | sor(s) |               | ss | itr | wer |
|        |                 |        |               | (  | ate | (W) |
|        |                 |        |               | g) | (b  |     |
|        |                 |        |               |    | /s) |     |
+--------+-----------------+--------+---------------+----+-----+-----+
| TDS    | wave form       | E1, B1 | 30mV/m, 0.2   | 8  | 1   | 1.0 |
|        |                 |        | pT            | 00 | 700 |     |
|        | 1,000,000       |        |               |    |     |     |
|        | samples/s       |        |               |    |     |     |
+--------+-----------------+--------+---------------+----+-----+-----+
| PWR    | 2 kHz -- 128    | E1, B1 | 10 nV/Hz,     | 8  | 1   | 0.7 |
|        | kHz             |        |               | 00 | 250 |     |
|        |                 |        | 10^--5^ nT/Hz |    |     |     |
+--------+-----------------+--------+---------------+----+-----+-----+
| RAD1   | 100 kHz -- 16   | E1,    | 10 nV/Hz,     | 8  | 800 | 0.8 |
|        | MHz             | (B2)   |               | 00 |     |     |
|        |                 |        | 10^--6^ nT/Hz |    |     |     |
+--------+-----------------+--------+---------------+----+-----+-----+
| RAD2   | 16 MHz --200    | E2     | \~ 10^-20^ W  | 7  | 600 | 0.7 |
|        | MHz             |        | m^2^ Hz^-1^   | 00 |     |     |
+--------+-----------------+--------+---------------+----+-----+-----+
| RAD3   | 200 MHz -- 1    | E3     | \~ 10^-20^ W  | 7  | 600 | 0.7 |
|        | GHz             |        | m^2^ Hz^-1^   | 00 |     |     |
+--------+-----------------+--------+---------------+----+-----+-----+
| D      |                 |        |               | 12 |     | 1.5 |
| SP/DPU |                 |        |               | 00 |     |     |
+--------+-----------------+--------+---------------+----+-----+-----+
| Con    |                 |        |               | 8  |     | 1.2 |
| verter |                 |        |               | 00 |     |     |
+--------+-----------------+--------+---------------+----+-----+-----+
| H      |                 |        |               | 5  | 50  |     |
| arness |                 |        |               | 00 | (H  |     |
| etc.   |                 |        |               |    | /K) |     |
+--------+-----------------+--------+---------------+----+-----+-----+
| *      |                 |        |               | 63 | 5   | 6.8 |
| Total* |                 |        |               | 00 | 000 |     |
+--------+-----------------+--------+---------------+----+-----+-----+

Table 8.3.c: Main characteristics of the RPW analysers and subsystems

**Resources and requirements**

*Antenna systems*

Antenna systems for the range from 20 MHz to 1 GHz are readily available
off the shelf for ground based systems, including 'portable' ones.
Different types and techniques are available and the theory of the
measurement is well known. They would have to be developed specifically
for the Solar Orbiter mission. The RAD2 band (16 MHz - 200 MHz) could be
covered by an active dipole antenna and the RAD3 band (200 MHz - 1 GHz)
by a log-periodic antenna. These two antenna systems can be grouped
together and mounted on the same boom in order to have a common
mechanical (deployment mechanism) and electrical interface with the
spacecraft. They would have to be pointed to the sun (accuracy of a few
degrees) and thus protected from the direct sunlight by a shield not
affecting the radio wave propagation (radome type). The whole system can
be kept to a reasonable size (longest element of the order of 1 m) and
provide the required gain. The beam shape of the antennas mounted on the
spacecraft will have to be modelled and measured using existing
techniques.

*Electro-Magnetic Cleanliness (EMC)*

Many of the RPW scientific objectives are based on the measurement of
low-level signals for which maximum sensitivity is required. Several
inexpensive measures can be taken at the spacecraft and project level to
ensure that the Solar Orbiter spacecraft is clean from the point of view
of both conducted and radiated electromagnetic interference. The
sensitivity of the Solar Orbiter/RPW instrument will be approximately
10^-8^ V/m /Hz in the frequency range of 1 kHz to 20 MHz and 310^-7^
V/m/Hz below 100 Hz. In addition, the wave form analyser is sensitive to
impulsive interference of duration as short as a fraction of a
microsecond. Although these sensitivities may appear to require an
excessively \"clean\" spacecraft, they are not difficult to achieve if
good EMC practices are incorporated in the spacecraft design.

It is important to realise that an EMC-clean spacecraft is beneficial
for every instrument, as well as for spacecraft systems. Given the
progress made in understanding the principles, these EMC controls are
easy to implement and do not require extensive or expensive testing. The
Project could set up a consultant EMC board to study the general
grounding and shielding philosophies, instrument-particular
exceptions/waivers, and test procedures for the mission. This committee
would define an inexpensive, coherent scheme for minimising
interference. For example, a single point or \"star\" ground has often
been suggested as the path to a clean spacecraft. This can be a
burdensome requirement on many types of measurement and has been shown
not to be very effective at high frequencies. Instead of such arbitrary
rules, the EMC committee would work to find acceptable compromise using
a \"best engineering practices\" point of view. Many of the interests of
an EMC committee will overlap with other spacecraft cleanliness issues
such as magnetic cleanliness and electrostatic cleanliness.

Currents induced in the spacecraft structure are a major source of EMC
disturbances because any voltage on the spacecraft surface induced by
such currents will couple readily into the antennas. Good grounding and
harness shielding practices will control these currents and should be
the basis of an EMC \"philosophy\". Power converters also represent a
serious source of interference if left uncontrolled. The EMC committee
would develop a plan for the use of crystal controlled DC-DC conversion
in all spacecraft systems (including instruments) which can limit the
interference to narrow frequency bands. The committee would review all
subsystem test results and, in the case of test results in excess of
defined limits, determine what remedial action (if any) is required.

**Instrument summary**

RPW Sensors:

-   triaxial search coil mounted on the magnetometer boom

-   magnetic loop antenna

-   boom antenna (5 to 6 m on side opposed to sun)

-   VHF/UHF system comprising an active dipole and a log-periodic
    antenna

Role of main subsystems:

-   RPW/PWA: identify plasma waves and kinetic modes (high frequency
    part of the fluctuation and turbulence spectra)

-   RPW/RAS: measure and track solar and interplanetary radio bursts;
    connect low-altitude coronal regions observed by the optical
    instruments with the *in situ*, local solar wind conditions at the
    spacecraft, study directivity of radio emissions

+------------------+-----------+-----------+-------------+------------+
| **Instrument**   | **Mass**  | **Power** | **Volume    | **Data     |
|                  |           |           | (cm^3^)**   | rate       |
|                  | **(kg)**  | **(W)**   |             | (bps)**    |
+------------------+-----------+-----------+-------------+------------+
| **Total RPW**    | 10        | 7.5       | 7500        | 5000       |
+------------------+-----------+-----------+-------------+------------+

**8.2.3 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 Solar Orbiter
orbit and the proximity of the spacecraft to the sun will create unique
situations for this kind of analysis. The CRS investigation 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.

It is presently unclear whether or not an ultra-stable oscillator would
be necessary to guarantee the frequency stability of the downlinks. 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). The
practicality of using linear polarised signals for Faraday rotation
measurements is another trade-off to be explored. The final decision
depends critically on the geometry and mission operation plan.
Considerable interaction between the radio scientists and the
spacecraft's radio subsystem contractor is a prerequisite for conducting
a successful coronal sounding experiment. The occultation geometry
during the solar conjunctions is another aspect that needs to be worked
in detail.

**8.2.4 Magnetometer (MAG)**

**Scientific drivers**

Knowledge of the magnetic field is of crucial importance for
characterising the state of the local interplanetary medium and
indispensable for the interpretation of measurements of energetic
particles and thermal plasma particles.

The magnetometer will measure the three orthogonal components of the
interplanetary magnetic field vector over a broad dynamic range from a
few nT up to the T range, and for frequencies well below the local
heavy-ion or proton gyrofrequencies.

**Instrument concept**

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. 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. The MAG
sensor and electronics for the Solar Orbiter require 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 \< 1m) 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 shadow. 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 of MAG could
be shared with either the plasma instruments or the wave instruments.

##### 

##### 

##### Resources and requirements

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.

**Instrument summary**

Instrument triaxial fluxgate magnetometer

Telemetry \< 200 bps

Mass \< 1 kg

Power \< 1 Watt

Special requirements mounting on short boom (\<1 m), in shadow of
spacecraft, basic magnetic cleanliness of spacecraft.

5.  **Energetic Particle Detector (EPD)**

    **Scientific drivers**

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 plasma measurements. It is essential that all
measurements are done fast, typically at 1 s, with complete angular
coverage to resolve the pitch-angle distributions. The EPD will

-   determine *in-situ* the generation, storage, release and propagation
    of different species of solar energetic particles in the inner
    heliosphere,

-   identify the links between magnetic activity and acceleration on the
    Sun for energetic particles, by virtue of combined remote-sensing of
    their source regions and *in-situ* measurements of their properties,
    while staying magnetically connected with the acceleration sites
    during solar corotation passes,

-   characterise gradual (CMEs) and impulsive (flares) particle events
    and trace their spatial and temporal evolution near the Sun

**Instrument concept**

The EPD will determine chemical and charge composition and energy
spectra 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 and is also a
possible design option for the Solar Orbiter.

The specific design suggested here resembles closely the one recently
selected for the IMPACT instrument on the STEREO mission. It is a
multi-head sensor system using solid state technology. Ongoing research
in this field is promising and will certainly lead to further
improvements concerning energy thresholds and noise level as well as to
size reduction.

**Resources and requirements**

The measurement 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. The whole EPD package will have a total mass not exceeding 4
kg and consume an average power of 3 Watts. The required telemetry rate
is 2 kb/s.

The characteristics of the EPD could, e.g., resemble the 4 sensors of
the IMPACT instrument on STEREO, as given in Table 8.5.

+-------+--------------+-----------+--------+------------------------+
| **Sen | **Energy     | **GF**    | *      | **Name**               |
| sor** | (Me          |           | *FOV** |                        |
|       | V/nucleon)** | **(cm^2^  |        |                        |
|       |              | sr )**    | **(° x |                        |
|       |              |           | °)**   |                        |
+-------+--------------+-----------+--------+------------------------+
| **    | e: 0.02 -    | 4 x 0.30, |        | Electron proton        |
| EPT** | 0.40         |           |        | telescope              |
|       |              | 4 x 0.24, |        |                        |
|       | p: 0.02 -    |           |        |                        |
|       | 7.00         | 4 view    |        |                        |
|       |              | cones     |        |                        |
+-------+--------------+-----------+--------+------------------------+
| **    | 0.03 - 2     | 0.3,      | 17 x   | Supra-thermal Ion      |
| SIT** |              |           | 44     | Telescope              |
|       |              | 1 view    |        |                        |
|       |              | cone      |        | He - Fe                |
+-------+--------------+-----------+--------+------------------------+
| **    | 1.5 - 40     | 4.5,      | 130 x  | Low Energy Telescope   |
| LET** |              |           | 30     |                        |
|       |              | 10 view   |        | He - Ni                |
|       |              | cones     |        |                        |
+-------+--------------+-----------+--------+------------------------+
| **    | 13 - 100     | 2.4,      | 48 x   | High Energy Telescope  |
| HET** |              |           | 48     |                        |
|       |              | 2 view    |        | He - Fe                |
|       |              | cones     |        |                        |
+-------+--------------+-----------+--------+------------------------+

Table 8.5: EPD instrument characteristics

**Instrument summary**

+------------------+-----------+-----------+-------------+------------+
| **Instrument**   | **Mass**  | **Power** | **Volume    | **Data     |
|                  |           |           | (cm^3^)**   | rate       |
|                  | **(kg)**  | **(W)**   |             | (b/s)**    |
+------------------+-----------+-----------+-------------+------------+
| **EPT**          | 0.94      | 0.60      | 300         | 300        |
+------------------+-----------+-----------+-------------+------------+
| **SIT**          | 0.78      | 0.55      | 620         | 600        |
+------------------+-----------+-----------+-------------+------------+
| **LET**          | 0.44      | 0.12      | 430         | 800        |
+------------------+-----------+-----------+-------------+------------+
| **HET**          | 0.27      | 0.04      | 210         | 300        |
+------------------+-----------+-----------+-------------+------------+
| **DPU/H-L-VPS**  | 1.37      | 1.17      | 3000        |            |
+------------------+-----------+-----------+-------------+------------+
| **Total EPD**    | **3.80**  | **2.48**  | **4560**    | **2000**   |
+------------------+-----------+-----------+-------------+------------+

**8.2.6 Dust Detector (DUD)**

**Science drivers**

The DUD will analyse interplanetary dust particles with respect to their
composition and orbital characteristics and determine *in-situ* the
distribution, composition and dynamics of dust particles in the near-Sun
heliosphere in and out of the ecliptic. The spatial distribution of
particles with masses ranging between 10^-16^ g and 10^-6^ g will be
determined. DUD will

-   help outlining the extent of the dust-free zone around the Sun,

-   find out the sources of dust, e.g. from Sun-grazing comets,

-   unravel the role played by near-Sun dust for pick-up ions,

-   deliver data relevant for an understanding of proto-planetary discs.

**Instrument concept**

The science objectives can be met with an instrument of, e.g., Giotto
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 look as
sketched in Fig 8.3. The dust particles enter the detector through a
thin aluminium foil and evaporate during impact on a gold target. A set
of wires in front of the target plate collects the produced ions. The
signal is analysed to deduce the mass and possibly speed of the dust
particle.

![](media/image3.png){width="6.072916666666667in" height="2.4375in"}

Fig.8.3: Instrument design of DUD **to be improved by BF**

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 but may require a
more sophisticated design.

###### 

###### Resources and requirements

Two identical detector units should be located on the sides of the
spacecraft body, one looking 90 degrees off the spacecraft-sun line in
the orbital plane to measure the near-ecliptic dust particles, the other
one in the direction perpendicular to the orbital plane in order to
measure dust particles in orbits highly inclined to the ecliptic. The
FOV of each sensor is ± 80°.

The mass of each sensor is 0.5 kg, consumes a power of 0.3 W and has a
volume of 1200 cm^3^. The average bitrate will not exceed 50 b/s. The
DPU should

be shared with other particle instruments.

####### 

####### 

####### Instrument summary

+---------------+------------+------------+-------------+------------+
| *             | **Mass**   | **Power**  | **Volume    | **Data     |
| *Instrument** |            |            | (cm^3^)**   | rate       |
|               | **(kg)**   | **(W)**    |             | (b/s)**    |
+---------------+------------+------------+-------------+------------+
| DUD           | 1          | 0.6        | 2400        | 50         |
+---------------+------------+------------+-------------+------------+

**8.2.7 Neutral Particle Detector (NPD)**

## Scientific drivers

The L 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 atoms and thus render possible the detection of the
interplanetary \"neutral solar wind\". 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^-1^ at 0.21 AU, but it could be up to 10^6^ atoms cm^-2^ s^-1^ in a
CME. 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.

**Instrument concept**

Neutral atom detectors have been flown on SOHO, CASSINI and IMAGE.
Time-of-flight instruments are for example flown on ULYSSES, SOHO and
ACE.

As a baseline, the instrument parameters of the NPD on the Solar Orbiter
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 entrance system must block radiation emerging from the
solar disk. The UV scattered in the solar corona, i.e. L, 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.

## A feasible neutral atom detector (see Fig. 8.4) may consist of two sections: the first compartment accommodates the baffle plates keeping out solar UV and the deflection unit for keeping off unwanted charged particles. The second section contains a time-of-flight (TOF) double-coincidence detector. Between 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 sawtooth-like structures. The collimator consists of,
alternatively positive and negative, high-voltage biased plates and
therefore acts as a charged-particle deflector. The collimator plate
length, distance and potential determine the cut-off of the
charged-particle transmission (about 100 keV/q). L light 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 matching the
L peak. 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
g/cm^2^ carbon foil. As they pass the carbon foil, they release
secondary electrons which are accelerated towards a START MCP
(multi-channel plate). The atoms travel through a 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/image4.png){width="4.825in" height="3.1006944444444446in"}

Figure 8.4: The instrument schematics of the Neutral Particle Detector.

## 

## Resources and requirements

The NPD can be built for a mass of 1 kg. The power consumption is 2 W.
The required telemetry rate is less than 0.5 kb/s. It is envisaged to
use a common DPU and to share the mechanical structure with the SWA. The
FOV has to be 17° off the S/C - Sun line because of the typical 10°
aberration of a 1 keV hydrogen atom travelling at solar wind speed.

**Instrument summary**

+---------------+------------+------------+-------------+------------+
| *             | **Mass**   | **Power**  | **Volume    | **Data     |
| *Instrument** |            |            | (cm^3^)**   | rate       |
|               | **(kg)**   | **(W)**    |             | (bps)**    |
+---------------+------------+------------+-------------+------------+
| NPD           | 1.00       | 2.00       | 2000        | 500        |
+---------------+------------+------------+-------------+------------+

**8.2.8 Neutron Detector (NED)**

**Scientific drivers**

Solar neutrons (and rays) are a by-product of solar flares due to
spallation reactions between energetic protons and particles. The
overwhelming majority of these neutrons are thermalised within minutes
and later absorbed by hydrogen to form deuterium and the well known 2.2
MeV line. Some neutrons, however, escape into space and can eventually
be detected (before they decay with a half-life of 619.9 s) by a sensor
located in close proximity to the Sun. The temporal correlation of
neutron fluxes with solar energetic particle fluxes and with optical
observations in the EUV provides crucial information on the solar
particle acceleration process, which is inaccessible by other means.

The advantage of neutrons over charged particles not being delayed by
travel along complicated paths through magnetic fields will allow an
unambiguous temporal correlation with optical observations of features
of flares in different energy ranges.

Incorporation of NED on the Solar Orbiter which offers routine
monitoring of neutron fluxes and, optionally, rays will be a "first" in
solar science. A major contribution to the solar ray flux stems from the
2.2 MeV line, which is due to the reaction of thermal neutrons with
protons to form deuterium. Hence, the ray detector registering "extinct"
solar neutrons would ideally complement the neutron detector, which is
to record "life" neutrons in situ.

**Instrument concept**

Following the design of the instrument proposed for the Solar Probe the
heart of the suggested NED instrument consists of a cell with thin (\<
10 m) thorium foils. With a cross section of typically 100 mbarn (at
energies above 1 MeV) an efficiency of 10^-5^ for producing fission
fragments - to be detected in a pulse ionisation chamber - can be
achieved. The anti-coincidence chamber surrounding the sensitive volume
contains a counting gas, e.g. methane or argon, and energetic particles
which might produce false counts in the ionisation chamber will be
detected with an efficiency of close to 100%. The angular resolution of
the NED will not be sufficient to map individual features on the solar
surface.

Optionally, incorporation of a facility to monitor solar rays into the
experiment would greatly enhance the scientific return.

Planetary flybys could be used to block the solar neutron flux and to
determine the intrinsic background rate of the spacecraft. Further
information, also on the average energy of solar neutrons will be gained
by comparison of long time averages of the neutron flux at different
solar distances.

**Resources and requirements**

The main concern with NED will be to minimise the energetic particle
background in the sensor cell. On one hand the anti-coincidence cell
should be as efficient as possible to prevent false coincidences, on the
other hand, the location of the sensor within the payload has to be
chosen such that as to keep the particle radiation level at a minimum.

Fig. 8.5 Schematics of the NED sensor

**Instrument summary**

+------------------+-----------+-----------+-------------+------------+
| **Instrument**   | **Mass**  | **Power** | **Volume    | **Data     |
|                  |           |           | \[cm^3^\]** | rate       |
|                  | *         | **\[W\]** |             | \[b/s\]**  |
|                  | *\[kg\]** |           |             |            |
+------------------+-----------+-----------+-------------+------------+
| **NED**          | 1         | 1         | 1000        | 100        |
+------------------+-----------+-----------+-------------+------------+
| **NED+ ^\*)^**   | 2         | 1.5       | 2000        | 150        |
+------------------+-----------+-----------+-------------+------------+
| **^\*)^**        |           |           |             |            |
| Version          |           |           |             |            |
| including ray    |           |           |             |            |
| detection        |           |           |             |            |
| facility         |           |           |             |            |
+------------------+-----------+-----------+-------------+------------+

**8.2.8 Heliospheric instrument summary**

+-----------+----+--------+--------+-----+-----------+-----+--------+
| Name      | A  | Measur | Spe    | M   | Size      | Po  | Tel    |
|           | cr | e-ment | cifica | ass |           | wer | emetry |
|           | o- |        | -tions |     | cm x cm x |     |        |
|           |    |        |        | kg  | cm        | W   | kb/s   |
|           | n  |        |        |     |           |     |        |
|           | ym |        |        |     |           |     |        |
+-----------+----+--------+--------+-----+-----------+-----+--------+
| Solar     | S  | T      | 0-30   | 6   | 20 x 20 x | 5   | 5      |
| Wind      | WA | hermal | keV/Q; |     | 20        |     |        |
|           |    |        |        |     |           |     |        |
| Plasma    |    | ions   | 0-10   |     |           |     |        |
|           |    |        | keV    |     |           |     |        |
| Analyser  |    | and    |        |     |           |     |        |
|           |    | ele    |        |     |           |     |        |
|           |    | ctrons |        |     |           |     |        |
+-----------+----+--------+--------+-----+-----------+-----+--------+
| Radio and | R  | AC     | V/m -  | 10  | 15 x 20 x | 7.5 | 5      |
|           | PW | el     | V/m    |     | 30        |     |        |
| Plasma    |    | ectric |        |     |           |     |        |
|           |    |        | 0      |     | (el       |     |        |
| Wave      |    | and    | .1nT - |     | ectronics |     |        |
| Analyser  |    | ma     | T      |     | box)      |     |        |
|           |    | gnetic |        |     |           |     |        |
|           |    | fields |        |     |           |     |        |
+-----------+----+--------+--------+-----+-----------+-----+--------+
| Magn      | M  | DC     | to 500 | 1   | 10 x 10 x | 1   | 0.5    |
| eto-meter | AG | ma     | Hz     |     | 10        |     |        |
|           |    | gnetic |        |     |           |     |        |
|           |    | field  |        |     |           |     |        |
+-----------+----+--------+--------+-----+-----------+-----+--------+
| Energetic | E  | Solar  | Ions   | 4   | 10 x 20 x | 3   | 2      |
| Particle  | PD | and    | and    |     | 20        |     |        |
| Detector  |    | cosm   | ele    |     |           |     |        |
|           |    | ic-ray | ctrons |     |           |     |        |
|           |    | par    | .01-10 |     |           |     |        |
|           |    | ticles | MeV    |     |           |     |        |
+-----------+----+--------+--------+-----+-----------+-----+--------+
| Neutral   | N  | N      | 0.6 -  | 1   | 10 x 10 x | 2   | 0.5    |
| Particle  | PD | eutral | 100    |     | 20        |     |        |
| Detector  |    | hy     | keV    |     |           |     |        |
|           |    | drogen |        |     |           |     |        |
|           |    | and    |        |     |           |     |        |
|           |    |        |        |     |           |     |        |
|           |    | atoms  |        |     |           |     |        |
+-----------+----+--------+--------+-----+-----------+-----+--------+
| Dust      | D  | In     | Mass   | 2   | 10 x 10 x | 1   | 0.5    |
| Detector  | UD | terpla | (g):   |     | 10        |     |        |
|           |    | netary |        |     |           |     |        |
|           |    |        | 10     |     |           |     |        |
|           |    | dust   | ^-16^- |     |           |     |        |
|           |    | par    | 10^-6^ |     |           |     |        |
|           |    | ticles |        |     |           |     |        |
+-----------+----+--------+--------+-----+-----------+-----+--------+
| Radio     | C  | Wind   | X-band | (   | 5 x 5 x 5 | 3   | 0      |
|           | RS | d      |        | USO |           |     |        |
| Sounding  |    | ensity | K      | 20  |           |     |        |
|           |    |        | a-band | 0g) |           |     |        |
|           |    | and    |        |     |           |     |        |
|           |    | ve     |        |     |           |     |        |
|           |    | locity |        |     |           |     |        |
+-----------+----+--------+--------+-----+-----------+-----+--------+
| Neutron   | N  | Solar  | e      | 1   | 10 x 10 x | 1   | 0.2    |
| detector  | ED |        | \>1MeV |     | 10        |     |        |
|           |    | ne     |        |     |           |     |        |
|           |    | utrons |        |     |           |     |        |
+-----------+----+--------+--------+-----+-----------+-----+--------+

**Table 8.6: *In-situ* heliospheric instrumentation**

**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 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/image8.png){width="4.35in"
height="2.1743055555555557in"}![](media/image9.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/image11.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, **19000**
  CCD)                                                      

  \- 25 cm telescope (mirrors + structure)                  4000

  \- prefilter (@ primary hole)                             400

  \- rotating quarter wave plate                            800

  \- linear polarizer                                       200

  \- folding mirror                                         400

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

  \- cube beam-splitter                                     200

  \- ISS lens                                               200

  \- ISS detector                                           800

  \- re-imaging system (4x200)                              800

  \- Fabry-Perots and Oven                                  3000

  \- shutter                                                300

  \- CCD camera (incl. electronics)                         2500

  \- harness (25% of total above)                           3900

                                                            

  **2. Full Disk Telescope**                                **3000**

  \- prefilter                                              200

  \- focal lenses (4\*200)                                  800

  \- rotating quarter wave plate                            800

  \- linear polarizer                                       200

  \- folding mirror                                         400

  \- harness (25% of total above)                           600

                                                            

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

                                                            

  **4. VIM TOTAL**                                          **28000**
  --------------------------------------------------------- -------------

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 8500 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. We expect a more 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                           | 23 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 the Solar Orbiter 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/image12.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/image13.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     | 25 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 L
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/image14.png){width="2.841666666666667in"
height="2.975in"}![](media/image15.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 (304 20) Å                             |
|                   |                                                  |
|                   | 2\) H I (1216 100) Å                             |
|                   |                                                  |
|                   | 3\) Visible (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         | 5 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)  |
+-------------------+--------------------------------------------------+

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-   **8.3.6** **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 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        | Fe      | 28   | 30 x 40  | 28  | 20    |
| e-light |      | igh-res. | 630nm   |      | x 120    |     |       |
| Imager  |      | disk     |         |      |          |     |       |
| and     |      |          |         |      |          |     |       |
|         |      | imaging  |         |      |          |     |       |
| Magne   |      | and      |         |      |          |     |       |
| tograph |      |          |         |      |          |     |       |
|         |      | pol      |         |      |          |     |       |
|         |      | arimetry |         |      |          |     |       |
+---------+------+----------+---------+------+----------+-----+-------+
| EUV     | EUS  | Imaging  | EUV     | 23   | 30 x 15  | 30  | 20    |
| Imager  |      | and      | e       |      | x 140    |     |       |
|         |      | dia      | mission |      |          |     |       |
| and     |      | gnostics | lines   |      |          |     |       |
|         |      | of       |         |      |          |     |       |
| Spect   |      |          |         |      |          |     |       |
| rometer |      | TR and   |         |      |          |     |       |
|         |      | corona   |         |      |          |     |       |
+---------+------+----------+---------+------+----------+-----+-------+
| EUV /   | EXI  | Corona   | He and  | 37   | 10 x 10  | 19  | 20    |
| X-ray   |      | imaging  | Fe      |      | x 30     |     |       |
|         |      |          |         |      |          |     |       |
| Imager  |      |          | Ion     |      |          |     |       |
|         |      |          | lines   |      |          |     |       |
+---------+------+----------+---------+------+----------+-----+-------+
| Ultr    | UVC  | Imaging  | Coated  | 18   | 20 x 20  | 35  | 5     |
| aviolet |      | and      | mirror  |      | x 50     |     |       |
| and     |      | dia      | coro    |      |          |     |       |
| Visible |      | gnostics | nagraph |      |          |     |       |
| Light   |      | of the   |         |      |          |     |       |
| Coro    |      | corona   | CCD     |      |          |     |       |
| nagraph |      |          | d       |      |          |     |       |
|         |      |          | etector |      |          |     |       |
+---------+------+----------+---------+------+----------+-----+-------+
| Rad     | RAD  | Solar    | Visible | 4    | 15 x 15  | 7   | 0.75  |
| iometer |      | constant | light   |      | x 25     |     |       |
+---------+------+----------+---------+------+----------+-----+-------+