ESA UNCLASSIFIED – For Official Use
Prepared by Yannis Zouganelis, Anik De Groof, Andrew Walsh, David Williams, Daniel Müller
Reference SOL-EST-PL-8539
Issue 0
Revision 1
Date of Issue 10 July 2017
Status Draft
Document Type
Distribution
Solar Orbiter Science Activity Plan
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Title
Issue 0 Revision
Author Yannis Zouganelis Date
Approved by Date
Daniel Müller (ESA Project Scientist)
Chris St. Cyr (NASA Project Scientist)
Javier Rodríguez-Pacheco (EPD PI)
Pierre Rochus (EUI PI)
Timothy Horbury (MAG PI)
Ester Antonucci (METIS PI)
Sami Solanki (PHI PI)
Milan Maksimovic (RPW PI)
Russ Howard (SoloHI PI)
Frédéric Auchère (SPICE Operations PI)
Sam Krucker (STIX PI)
Christopher J. Owen (SWA PI)
Reason for change Issue Revision Date
Issue 0 Revision
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Table of Contents:
1 INTRODUCTION ....................................................................................................................... 7 1.1 Purpose and Scope ........................................................................................................................ 7 1.2 Structure of the Document ........................................................................................................... 7 1.3 Assumptions and Definitions ....................................................................................................... 7 1.4 Applicable Documents ................................................................................................................. 8 1.5 Reference Documents ................................................................................................................... 8 2 MISSION PLANNING OVERVIEW ........................................................................................ 9 2.1 Mission Planning Cycle ................................................................................................................ 9 2.1.1 Long-Term Planning (LTP) ...................................................................................................... 10 2.1.2 Medium-Term Planning (MTP)................................................................................................ 10 2.1.3 Short-Term Planning (STP) ...................................................................................................... 10 2.1.4 Very-Short-Term planning (VSTP) .......................................................................................... 10 3 SCIENCE OBJECTIVES ......................................................................................................... 12 3.1 Solar Orbiter’s Top-Level Science Objectives ........................................................................... 12 3.1.1 What drives the solar wind and where does the coronal magnetic field originate? ................. 12 3.1.2 How do solar transients drive heliospheric variability? ........................................................... 13 3.1.3 How do solar eruptions produce energetic particle radiation that fills the heliosphere? .......... 13 3.1.4 How does the solar dynamo work and drive connections between the Sun and the heliosphere?
14 3.2 Lower-Level Objectives and Required Observations................................................................. 15 3.2.1 Objective 1: What drives the solar wind and where does the coronal magnetic field originate?
15 3.2.1.1 1.1 What are the source regions of the solar wind and heliospheric magnetic field? ............ 16 3.2.1.2 1.2 What mechanisms heat the corona and heat and accelerate the solar wind? ................... 32 3.2.1.3 1.3 What are the sources of solar wind turbulence and how does it evolve? ......................... 39 3.2.2 Objective 2: How do solar transients drive heliospheric variability? ....................................... 45 3.2.2.1 2.1 How do CMEs evolve through the corona and inner heliosphere? .................................. 45 3.2.2.2 2.2 How do CMEs contribute to solar magnetic flux and helicity balance? .......................... 49 3.2.2.3 2.3 How and where do shocks form in the corona? ............................................................... 51 3.2.3 Objective 3: How do solar eruptions produce energetic particle radiation that fills the
heliosphere? ....................................................................................................................................... 55 3.2.3.1 3.1 How and where are energetic particles accelerated at the Sun? ....................................... 55 3.2.3.2 3.2 How are energetic particles released from their sources and distributed in space and
time? 68 3.2.3.3 3.3 What are the seed populations for energetic particles? .................................................... 73 3.2.4 Objective 4: How does the solar dynamo work and drive connections between the Sun and the
heliosphere? ....................................................................................................................................... 76 3.2.4.1 4.0 Overall remarks and feasibility concerning Objective 4 observations with Solar Orbiter
77 3.2.4.2 4.1 How is magnetic flux transported to and re-processed at high solar latitudes? ............... 80 3.2.4.3 4.2 What are the properties of the magnetic field at high solar latitudes? ............................. 83 3.2.4.4 4.3 What is the nature of magnetoconvection? ...................................................................... 85 3.2.4.5 4.4 Are there separate dynamo processes acting in the Sun? ................................................. 86
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3.2.4.6 4.5 How are coronal and heliospheric phenomena related to the solar dynamo? .................. 87 3.2.5 5. Additional science objectives ............................................................................................... 88 4 INSTRUMENT DESCRIPTIONS AND THEIR OPERATIONAL CONSTRAINTS ....... 93 4.1 EPD ............................................................................................................................................ 93 4.1.1 EPD Observables ..................................................................................................................... 93 4.1.1.1 Energy range .......................................................................................................................... 94 4.1.1.2 Mass resolution ...................................................................................................................... 95 4.1.1.3 Fields of view and angular resolution .................................................................................... 95 4.1.1.4 Geometric factor..................................................................................................................... 96 4.1.2 EPD modes ............................................................................................................................... 97 4.1.2.1 EPD Normal mode ................................................................................................................. 97 4.1.2.2 EPD Burst mode ..................................................................................................................... 97 4.2 EUI ............................................................................................................................................. 98 4.2.1 EUI observables ........................................................................................................................ 98 4.2.2 EUI modes and telemetry ......................................................................................................... 98 4.2.2.1 FSI Beacon mode (B) ........................................................................................................... 100 4.2.2.2 FSI Synoptic mode (S) ......................................................................................................... 100 4.2.2.3 FSI Reference Synoptic mode (R) ....................................................................................... 101 4.2.2.4 FSI Global eruptive event mode (G) .................................................................................... 102 4.2.2.5 FSI Find Event mode (FE) ................................................................................................... 103 4.2.2.6 FSI Faint High Corona mode (FHC) .................................................................................... 103 4.2.2.7 EUV & LYA Beacon modes (HB)....................................................................................... 104 4.2.2.8 EUV & LYA Coronal hole modes (C) ................................................................................. 104 4.2.2.9 EUV & LYA Quiet Sun modes (Q) ..................................................................................... 105 4.2.2.10 EUV & LYA Active Region modes (A) ......................................................................... 106 4.2.2.11 EUV & LYA Eruptive Event modes (E) ........................................................................ 106 4.2.2.12 EUV & LYA Discovery modes (D) ................................................................................ 107 4.3 MAG ......................................................................................................................................... 108 4.3.1 MAG observables ................................................................................................................... 109 4.3.2 MAG modes ........................................................................................................................... 109 4.3.2.1 MAG Normal mode ............................................................................................................. 109 4.3.2.2 MAG burst mode.................................................................................................................. 109 4.4 METIS ...................................................................................................................................... 110 4.4.1 METIS observables ................................................................................................................ 110 4.4.2 METIS modes ......................................................................................................................... 111 4.5 PHI ............................................................................................................................................ 119 4.5.1 PHI observables ...................................................................................................................... 119 4.5.2 PHI modes .............................................................................................................................. 119 4.6 RPW ......................................................................................................................................... 127 4.6.1 RPW observables .................................................................................................................... 129 4.6.2 RPW modes ............................................................................................................................ 129 4.7 SoloHI ...................................................................................................................................... 131 4.7.1 SoloHI observables ................................................................................................................. 131 4.7.2 SoloHI modes ......................................................................................................................... 131 4.8 SPICE ....................................................................................................................................... 134
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4.8.1 SPICE observables ................................................................................................................. 134 4.8.2 SPICE modes .......................................................................................................................... 134 4.9 STIX ......................................................................................................................................... 145 4.9.1 STIX observables ................................................................................................................... 146 4.9.2 STIX modes ............................................................................................................................ 146 4.10 SWA ......................................................................................................................................... 148 4.10.1 SWA observables ............................................................................................................... 149 4.10.2 SWA modes ....................................................................................................................... 149 5 SCIENCE ACTIVITIES ......................................................................................................... 153 5.1 Introduction and SOOPs ........................................................................................................... 153 5.2 List of SOOPs ........................................................................................................................... 153 5.2.1 I_DEFAULT ........................................................................................................................... 154 5.2.2 L_IS_STIX ............................................................................................................................... 161 5.2.3 L_IS_SoloHI_STIX ................................................................................................................ 163 5.2.4 L_FULL_LRES_MCAD_Coronal_Synoptic ......................................................................... 166 5.2.5 L_FULL_LRES_MCAD_ProbeQuadrature ........................................................................... 169 5.2.6 L_FULL_MRES_MCAD_CME_SEPs .................................................................................. 171 5.2.7 L_FULL_HRES_LCAD_MagnFieldConfig .......................................................................... 174 5.2.8 L_FULL_HRES_MCAD_Coronal_He_Abundance .............................................................. 176 5.2.9 L_FULL_HRES_HCAD_Eruption_Watch ............................................................................ 179 5.2.10 L_FULL_HRES_HCAD_Coronal_Dynamics................................................................... 182 5.2.11 L_SMALL_MRES_MCAD_Ballistic-connection............................................................. 186 5.2.12 L_SMALL_MRES_MCAD_Connection_Mosaic............................................................. 188 5.2.13 L_SMALL_HRES_HCAD_Fast_Wind............................................................................. 192 5.2.14 L_SMALL_HRES_HCAD_SlowWindConnection........................................................... 197 5.2.15 L_BOTH_LRES_MCAD_Pole-to-Pole............................................................................. 204 5.2.16 L_BOTH_MRES_MCAD_Farside_Connection ............................................................... 206 5.2.17 L_BOTH_MRES_MCAD_Flare_SEPs ............................................................................. 209 5.2.18 L_BOTH_HRES_LCAD_CH_Boundary_Expansion ....................................................... 215 5.2.19 R_FULL_LRES_HCAD_GlobalHelioseismology ............................................................ 219 6 PLANNING STRATEGY ....................................................................................................... 252 7 OCTOBER 2018 OPTION E TRAJECTORY AND MEDIUM TERM PLANNING...... 267 7.1 Option E trajectory ................................................................................................................... 267 7.2 Planning periods for Option E (MTPs) .................................................................................... 267 7.2.1 MTP05 - 2021/01/01 - 2021/07/01 ......................................................................................... 268 7.2.1.1 RS window (default) placement ........................................................................................... 268 7.2.1.2 Science planning .................................................................................................................. 268 7.2.2 MTP06 - 2021/07/01 - 2022/01/01 ......................................................................................... 270 7.2.2.1 RS windows (default placement): ........................................................................................ 270 7.2.2.2 New placement for the RS window: .................................................................................... 270 7.2.2.3 SOOP planning..................................................................................................................... 271 7.2.3 MTP07 - 2022/01/01 - 2022/07/01 ......................................................................................... 273 7.2.3.1 RS window (default) placement ........................................................................................... 273 7.2.3.2 Current scenario ................................................................................................................... 273 7.2.4 MTP08 - 2022/07/01 - 2023/01/01 ......................................................................................... 276
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7.2.4.1 RS window (default) placement ........................................................................................... 276 7.2.4.2 Current scenario ................................................................................................................... 276 7.2.5 MTP09 - 2023/01/01 - 2023/07/01 ......................................................................................... 279 7.2.6 MTP10 - 2023/07/01 - 2024/01/01 ......................................................................................... 281 7.2.7 MTP11 - 2024/01/01 - 2024/07/01 ......................................................................................... 284 7.2.8 MTP12 - 2024/07/01 - 2025/01/01 ......................................................................................... 287 7.2.9 MTP13 - 2025/01/01 - 2025/07/01 (EMP) ............................................................................. 289 7.2.10 MTP14 - 2025/07/01 - 2026/01/01 .................................................................................... 291 7.2.11 MTP15 - 2026/01/01 - 2026/07/01 .................................................................................... 293 7.2.12 MTP16 - 2026/07/01 - 2027/01/01 .................................................................................... 295 7.2.13 MTP17 - 2027/01/01 - 2027/07/01 .................................................................................... 297 7.2.14 MTP18 - 2027/07/01 - 2028/01/01 .................................................................................... 299 7.2.15 MTP19 - 2028/01/01 - 2028/07/01 .................................................................................... 301 7.2.16 MTP20 - 2028/07/01 - 2029/01/01 .................................................................................... 303 7.2.17 MTP21 - 2029/01/01 - 2029/07/01 .................................................................................... 304 8 SIMULATIONS ....................................................................................................................... 306 9 APPENDIX .............................................................................................................................. 307 9.1 Bibliography ............................................................................................................................. 307
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1 INTRODUCTION
1.1 Purpose and Scope
The Science Activity Plan (SAP) for Solar Orbiter describes in a structured way all scientific
activities to be carried out by the instruments throughout all science mission phases in order to fulfil
the Science Requirements of the mission. It tracks how high-level science objectives are mapped to
more specific scientific objectives and these, in turn, to scientific activities that will be scheduled at
specific times during the mission. As context information, it also contains instrument operations
scenarios and modelling, and a description of all mission phases.
This first draft, which is not yet approved by the SWT, is issued as Reference Document for the
Solar Orbiter OGS-SGS Joint Ground Segment Implementation Review. It has been built by the
joint efforts of the Science Working Team (SWT), the Science Operations Working Group
(SOWG), instrument teams, PIs and Co-Is, and external experts, under the coordination and
guidance of the Science Operations Centre (SOC) team and the Project Scientists.
In general, it is foreseen to make use of the successful Joint Operations Plan (JOP) concept of the
SOHO mission, which is also being used by the Hinode mission, adapted to the specific needs of
Solar Orbiter (hence to be called SOOP). The general characteristics of the scientific mission
planning approach for Solar Orbiter are described in Section 2.
1.2 Structure of the Document
Section 2 provides an overview of the mission planning process. Section 3 describes the mission’s
science objectives, on the basis of [AD01, AD02]. Section 4 contains brief descriptions of each
instrument along with its operational modes and constraints. Science activities, defined as sets of
observations (SOOPs) which address each of the scientific objectives, are described in Section 5.
Section 6 describes the planning strategy that is used for filling the mission timeline with the various
science activities. Section 7 provides mission profiles (orbits etc.) and the description of all
Medium-Term Planning periods together with the activities planned in each one. Section 8 will
contain in the future the results of the simulations performed regarding the SSMM fill state, the
instrument stores fill state, power consumption etc. Any additional information, as well as a
bibliography of scientific papers is provided in Section 9.
1.3 Assumptions and Definitions
This first draft (SAP v0) is sufficiently detailed for its purpose as a reference document for the
Ground Segment Implementation Review. However, it still lacks information and details before the
first version can be released. The following assumptions and limitations will be relaxed in future
versions:
• The Cruise Phase is not yet described.
• An October 2018 launch date has been assumed with the Option E trajectory. This is only
affecting the Sections 7-8 and part of Section 6. Sections 3-5 are trajectory independent.
• Even though the Objective 4 sub-objectives have been well detailed and the corresponding
SOOPs have been described, these have not been planned in the mission timeline. This is
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due to the fact of the recent project delays and postponement of the launch date, which
resulted in equivalent delays of the SAP activities. These will be resumed once the new
launch date and information on trajectory get stabilized.
• For the same reasons, the preliminary simulations (Section 8) performed for the current
version of the SAP are not included.
1.4 Applicable Documents
[AD01] Solar Orbiter Science Management Plan (SMP, SOL-EST-PL-00880)
[AD02] Solar Orbiter Science Requirements Document (SciRD, SOL-EST-RS-1858)
1.5 Reference Documents
[RD01] Technical Note: Solar Orbiter Pointing Strategy during Remote-Sensing Windows: Science
Requirements (SOL-EST-TN-4020)
[RD02] Solar Orbiter – Exploring the Sun-Heliosphere Connection (Solar Physics, Vol. 285, Issue
1-2, pp. 25-70, 2013)
[RD03] Solar Orbiter Science Implementation Requirements Document (SIRD, SOL-EST-RS-
4514)
[RD04] Solar Orbiter Science Operations Centre Science Implementation Plan (SIP, SO-SGS-PL-
001)
[RD05] Solar Orbiter Science Operations Concept Document (SOL-SGS-PL-0001)
[RD06] Solar Orbiter Instrument Operation Request Interface Control Document (IOR ICD, SOL-
SGS-ICD-0003)
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2 MISSION PLANNING OVERVIEW
While the general mission planning approach for all routine science operations of Solar Orbiter will
be built on the experience of ESA’s precursor solar system missions Mars Express, Venus Express,
and Rosetta, a fundamental difference with respect to planetary missions is the highly dynamic
nature of the Sun. Given the short time-scales on which the targets of remote-sensing observations
(e.g. solar active regions) change, together with the narrow fields-of-view of the high-resolution
imaging telescopes (which cover less than 3% of the solar disk at perihelion), turn-around times
between defining the pointing and executing the observations of at most three days is required
during the Remote-Sensing Windows, as described in [RD01].
The resulting requirements on the science operations planning for Remote-Sensing Windows are
summarized below, on the basis of [RD03]. The mission planning cycle described is the result of a
collaborative effort of OGS, SOC and PS.
2.1 Mission Planning Cycle
For each mission phase (Cruise Phase (CP), Nominal Mission Phase (NMP), Extended Mission
Phase (EMP)), a baseline science plan will be established and documented in the SAP before each
mission phase commences. This plan will take into account the general characteristics and major
constraints of each orbit. The SAP will be a ‘live’ document in the sense that it will be frequently
updated as more and more science activities get worked out in detail, feedback from earlier
observations (and their planning) is injected into the planning of later orbits, etc.
For each orbit, a typical mission planning cycle starts with the SWT deciding upon the top-level
science objectives for this orbit, based on the general goals previously agreed upon by the SWT and
defined in the SAP. Given this input, the SOWG defines a coherent mission-level observing plan. In
this task, they will be assisted by the SOC, which will provide detailed information on the resources
available (e.g., on-board memory management, telemetry downlink). In turn, the PI teams provide,
at fixed deadlines and with a fixed periodicity, inputs to the SOC for the requested science
operations, which implement the observing plan defined by the SOWG. The SOC passes a
consolidated request to the MOC which checks the requests against mission, environmental and
resource constraints.
It is important to stress that top-level science operations planning for the mission needs to be done
well in advance due to the severe constraints on Solar Orbiter’s data downlink volume and on-board
storage in the Solid State Mass Memory (SSMM). In particular, this entails feasibility studies of
planned operations taking into account the expected time-dependent telemetry downlink profile as
well as SSMM load levels.
The mission planning cycle for the routine science operations phase will therefore be divided into
different levels, as described in RD05:
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2.1.1 Long-Term Planning (LTP)
LTP will fix the ground station schedule. This process typically will take place once for each
mission phase (covering the entire mission phase) until the start of the NMP, and once for every 6-
months period, approximately covering one orbit around the Sun, after the start of NMP. In each
orbit, this process will in general commence after the end of the last Remote-Sensing Window,
counted from aphelion. At LTP level, event skeleton fields are provided by the OGS to the SGS,
which will contain all the information required by the SGS to conduct science planning.
2.1.2 Medium-Term Planning (MTP)
MTP will fix the usage of spacecraft resources. This process will typically take place at several
intervals during Cruise Phase, and once per orbit during NMP, covering the entire orbit. At MTP
level, the top-level science operations plan for the entire orbit is defined and a default pointing
profile for Remote-Sensing Windows is defined (with possible updates at VSTP level, see below).
2.1.3 Short-Term Planning (STP)
STP will generate detailed schedules of commands for the spacecraft and for the ground stations.
This process will take place typically every week covering one week as described in detail in RD06.
At STP level, the instrument activities can be modified, provided they fit into the resource envelope
defined at MTP level.
2.1.4 Very-Short-Term planning (VSTP)
In the case of Remote Sensing Windows (RSWs), VSTP may be required. This planning level, with
turn-around times of at most three days between observations and execution of the new Pointing
Request (PTR), is required for RSWs in which features on the solar disk, e.g. active regions, shall be
tracked over time. This is due to the short lifetimes and non-deterministic motion of targets on the
Sun [RD01] and the relatively large (compared to the fields-of-view of the high-resolution imaging
telescopes) and temperature-dependent absolute pointing error (APE) of the spacecraft. This VSTP
consists of (i) initial target selection and (ii) updates to the pointing.
Prior to the start of an RSW, a limited set of precursor observations with the full-disk imaging
telescopes of the EUI and PHI instruments is performed and downlinked with high priority. Based
on the returned data, the target for the start of the RSW will be defined. This step is required to
make a decision on the pointing of the spacecraft and, in turn, the high-resolution imaging
telescopes. In case the orbital constellation permits making this decision by means of other
observations (e.g. using ground-based telescopes), this step can be omitted.
During the course of an RSW, a limited set of daily low-latency data, consisting of full-disk and
high-resolution images, will be downlinked with high priority. Based on the evaluation of these
images, the pointing will be updated by means of uploading a PTR. The final science target
selection (and retargeting) shall be responsibility of a person nominated by the SOWG, who shall
make him/herself available according to the schedule required by the planning process. Refinement
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of the tracking of this target shall be done by SOC, according to a set of rules established by the
SOWG.
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3 SCIENCE OBJECTIVES
Solar Orbiter’s mission is to address the central question of heliophysics: How does the Sun create
and control the heliosphere? This, in turn, is a fundamental part of the second science question of
ESA’s Cosmic Vision programme: “How does the solar system work?” Solar Orbiter is specifically
designed to identify the origins and causes of the solar wind, the heliospheric magnetic field, solar
energetic particles, transient interplanetary disturbances, and the Sun's magnetic field itself.
The supersonic solar wind, driven by dynamic plasma and magnetic processes at the Sun’s surface,
expands to surround the solar system’s planets and the space far beyond. Below the surface, the
solar dynamo drives magnetic fields whose buoyancy brings them to the surface where they form
huge arcades of loops, which contain enormous amounts of stored energy. These magnetic loops are
stretched and sheared by the Sun’s differential rotation and unknown surface processes, eventually
erupting in explosions, which eject magnetic structures that fly into the solar system, occasionally
impacting the Earth and its magnetic shield with disruptive effects on space and terrestrial systems.
Understanding the complex physical processes at work in this system is the central goal of
heliophysics. Since the Sun and presumably the heliosphere are typical of many small stars and their
stellar spheres, these studies are relevant to astrophysics, but are unique since the Sun alone is close
enough for detailed study.
Over the past ~20 years, an international effort to understand the Sun and heliosphere has been
undertaken with an array of spacecraft carrying out both remote observations at visible, UV, and X-
ray wavelengths, as well as in-situ observations of interplanetary plasmas, particles, and fields.
Combined and coordinated observations from missions such as Ulysses, Yohkoh, SOHO, TRACE,
RHESSI, Hinode and STEREO have resulted in an enormous advance in our understanding of the
Sun and heliosphere, and have proven that critical progress in understanding the physics requires
both remote and in-situ observations working together.
Although our vantage point at 1 AU is close by astrophysical measures, it has been long known that
much of the crucial physics in the formation and activity of the heliosphere takes place much closer
to the Sun, and that by the time magnetic structures, shocks, energetic particles and solar wind pass
by Earth they have already evolved and in many cases mixed so as to blur the signatures of their
origin. With the proven effectiveness of combined remote and in-situ studies on the missions cited
above, it is expected that critical new advances will be achieved by flying a spacecraft combining
remote and in-situ observations into the inner solar system. From this inner-heliospheric vantage
point, solar sources can be identified and studied accurately and combined with in-situ observations
of solar wind, shocks, energetic particles, etc., before they evolve significantly.
3.1 Solar Orbiter’s Top-Level Science Objectives
The four top-level scientific questions to be addressed by Solar Orbiter are:
3.1.1 What drives the solar wind and where does the coronal magnetic field originate?
The solar corona continuously expands and develops into a supersonic wind that extends outward,
interacting with itself and with the Earth and other planets, to the heliopause boundary with
interstellar space, far beyond Pluto’s orbit. The solar wind has profound effects on planetary
environments and on the planets, themselves – for example, it is responsible for many of the
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phenomena in Earth’s magnetosphere and is thought to have played a role in the evolution of Venus
and Mars through the erosion of their upper atmospheres.
Two classes of solar wind – ‘fast’ and ‘slow’ – fill the heliosphere, and the balance between them is
modulated by the 11-year solar cycle. The fast solar wind (~700km/s and comparatively steady) is
known to arise from coronal holes. The slow solar wind (~400 km/s) permeates the plane of the
ecliptic during most of the solar cycle so it is important to Earth's space environment. The slow solar
wind differs from the fast wind in mass flux and composition, which is consistent with confined
plasma in the solar corona. The specific escape mechanism through the largely closed magnetic field
is not known since candidate sites and mechanisms cannot be resolved from 1 AU. Fast and slow
wind carry embedded turbulent fluctuations, and these also display different properties compatible
with different solar origins. It is thought that such fluctuations may be responsible for the difference
in heating and acceleration between different solar wind streams.
Understanding the physics relating the plasma at the solar surface and the heating and acceleration
of the escaping solar wind is crucial to understanding both the effects of the Sun on the heliosphere
and how stars in general lose mass and angular momentum to stellar winds.
3.1.2 How do solar transients drive heliospheric variability?
The largest transient events from the Sun are coronal mass ejections (CMEs), large structures of
magnetic field and material that are ejected from the Sun at speeds up to 3000 km/s. CMEs are also
of astrophysical interest since they are the dominant way that stars shed both magnetic flux and
magnetic helicity that build up as a result of the stellar dynamo. Interplanetary CMEs (ICMEs) are
the major cause of interplanetary shocks, but the locations and mechanisms by which shocks form
around them are not known since they occur in the inner solar system. Similarly, the longitudinal
structure of ICMEs is not directly observable from the ecliptic, while their extent has a large impact
on the acceleration of energetic particles. ICMEs are a major cause of geomagnetic storms but their
effectiveness at disrupting the magnetosphere is only loosely related to the parent CME, because the
evolution of the propagating cloud with the surrounding heliosphere is complex and has not been
well studied. These unknowns have direct impact on our ability to predict transient (“space
weather”) events that affect Earth.
3.1.3 How do solar eruptions produce energetic particle radiation that fills the
heliosphere?
Like many astrophysical systems, the Sun is an effective particle accelerator. Large solar energetic
particle (SEP) events produce highly energetic particles that fill the solar system with ionizing
radiation. CME-driven shocks can produce relativistic particles on time scales of minutes, and many
CMEs convert ~10% of their kinetic energy into energetic particles. Other processes produce high-
energy particles on magnetic loops without involving shocks. The multiple processes operating in
SEP events are not well understood or distinguishable from observations at 1 AU. In particular,
particles accelerated in the corona and inner heliosphere are scattered by inhomogeneities in the
interplanetary magnetic field (IMF) before they arrive at Earth, destroying much of the information
they carry about the processes that accelerated them. Particle transport and scattering in the inner
solar system are poorly understood since the turbulence properties cannot be determined from 1 AU.
The actual seed population of particles energized by CME-driven shocks in the inner solar system is
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unexplored, and needs to be understood to construct a complete picture of particle acceleration in
shock-related events.
3.1.4 How does the solar dynamo work and drive connections between the Sun and the
heliosphere?
The Sun’s magnetic field connects the interior of the star to interplanetary space and is dominated
by a quasi-periodic 11-year sunspot cycle that modulates the form of the heliosphere and strongly
affects the space environment throughout the solar system. The large-scale solar field is generated in
the Sun’s interior, within the convection zone, by a dynamo driven by complex three-dimensional
mass flows that transport and process magnetic flux. Despite notable advances in our knowledge
and understanding of solar magnetism made possible by Ulysses, SOHO, and Hinode observations
as well as by recent theoretical models and numerical simulations, fundamental questions remain
about the operation of the solar dynamo and the cyclic nature of solar magnetic activity. Of
paramount importance to answering these questions is detailed knowledge of the transport of flux at
high latitudes and the properties of the polar magnetic field. To date, however, the solar high
latitudes remain poorly known owing to our dependence on observations made from the ecliptic. In
addition to questions about the global dynamo and the generation of the large-scale field, there are
unanswered questions about the origin of the small-scale internetwork field observed in the quiet
photosphere. Is this weak field produced by turbulent local dynamo action near the solar surface?
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3.2 Lower-Level Objectives and Required Observations
In this section, we describe detailed lower-level objectives of each of the above four top science
objectives.
3.2.1 Objective 1: What drives the solar wind and where does the coronal magnetic field
originate?
Hot plasma in the Sun’s atmosphere flows radially outward into interplanetary space to form the
solar wind, filling the solar system and blowing a cavity in the interstellar medium known as
the heliosphere. During solar minimum, large-scale regions of a single magnetic polarity in the
Sun’s atmosphere – polar coronal holes – open into space and are the source of high speed (~700
km/s), rather steady solar wind flows. There is also a slow wind (300-500 km/s) that emanates from
magnetically complex regions at low latitudes and the periphery of coronal holes. It is highly
variable in speed, composition, and charge state. The origin of the slow wind is not known. At solar
maximum, this stable bimodal configuration gives way to a more complex mixture of slow and fast
streams emitted at all latitudes, depending on the distribution of open and closed magnetic regions
and the highly tilted magnetic polarity inversion line.
The fast wind from the polar coronal holes carries magnetic fields of opposite polarity into the
heliosphere, which are then separated by the heliospheric current sheet (HCS) embedded in the slow
wind. Measurements over a range of latitudes far from the Sun show that this boundary is not
symmetric around the Sun’s equator, but is on average displaced southward. This offset must reflect
an asymmetry on the Sun; but since there cannot be a mismatch between the inward and outward
magnetic flux on the Sun, its origin is unclear. In situ, the HCS is warped and deformed by the
combined effects of solar rotation and inclination of the Sun’s magnetic axis, effects that are even
more prominent at solar maximum.
The energy that heats the corona and drives the wind comes from the mechanical energy of
convective photospheric motions, which is converted into magnetic and/or wave energy. In
particular, both turbulence and magnetic reconnection are implicated theoretically and
observationally in coronal heating and acceleration. However, existing observations cannot
adequately constrain these theories, and the identity of the mechanisms that heat the corona and
accelerate the solar wind remains one of the unsolved mysteries of solar and heliospheric physics.
How the coronal plasma is generated, energized, and the way in which it breaks loose from the
confining coronal magnetic field are fundamental physical questions with crucial implications
for predicting our own space environment, as well as for the understanding of the natural plasma
physics of other astrophysical objects, from other stars, to accretion disks and their coronae, to
energetic phenomena such as jets, X- and gamma-ray bursts, and cosmic-ray acceleration.
The solar wind contains waves and turbulence on scales from millions of kilometers to below the
electron gyroradius. The turbulence scatters energetic particles, affecting the flux of particles that
arrives at the Earth; local kinetic processes dissipate the turbulent fluctuations and heat the plasma.
Properties of the turbulence vary with solar wind stream structure, reflecting its origins near the
Sun, but the turbulence also evolves as it is carried into space with the solar wind, blurring the
imprint of coronal conditions and making it difficult to determine its physical origin. The inner
heliosphere, where Solar Orbiter will conduct its combination of remote-sensing and in-situ
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observations, provides the ideal laboratory for understanding the magnetohydrodynamic turbulence
of natural plasmas expected to be ubiquitous in astrophysical environments.
In the following sections, we discuss in more detail three interrelated questions which flow down
from this top-level question: What are the source regions of the solar wind and the heliospheric
magnetic field? What mechanisms heat and accelerate the solar wind? What are the sources of
turbulence in the solar wind and how does it evolve?
3.2.1.1 1.1 What are the source regions of the solar wind and heliospheric magnetic field?
1.1.1 Source regions of the fast solar wind
1.1.1.1 Low FIP fast wind origins
Description of the objective:
The fast wind, which does not exhibit strong first ionization potential (FIP) enhancements, could
come directly from the photosphere, from small cool coronal loops and open magnetic funnels (Tu
et al., 2005; Schwadron and McComas, 2003) at the base of coronal holes or spicules, which also
exhibit small FIP enhancements. Remote observations have revealed many cases of macrospicules
undergoing reconnection and erupting within coronal holes (Yamauchi et al., 2005). Do they
contribute to the fast solar wind streams? Polar plumes have also long been suspected to be a
significant source of fast solar wind (Deforest et al., 1997), as well as polar regions within plumes
("interplume lanes") (Giordano et al., 2000). Micro-streams of plasma originating in the coronal
holes may be related to polar plumes (Neugebauer et al., 1995), though evidence for this is
controversial (e.g., McComas et al., 1996). However, the relation could be difficult to observe with
Solar Orbiter since large amplitude Alfvénic fluctuations generate micro-streams signals in the fast
stream (Matteini et al., 2013).
Remarks:
In order to address this objective, we need to observe if the fast wind comes from the different
above sources. This objective can be split into two different observational strategies:
• A pure coronal characterization by pointing to the center of a coronal hole and its
boundaries, mostly focused on remote sensing observations (closest possible at high
latitude).
• An attempt to observe fast wind in situ and link it to its origin back to the coronal hole. In
order to be connected, we have to choose a coronal hole at the west limb, and observe it
from a location close to the Sun (at least one perihelion window would be good for the high-
resolution observations of plume-interplume FIP variation) and preferably during high-
latitude windows. Another possibility would be to observe during a window before
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perihelion in order to obtain the overall context and then observe during the perihelion for
connectivity. Polar coronal holes could also be observed with METIS from the equatorial
plane (no need of high-latitude windows, since it can observe the sky above the hole). In this
case, however, there is less chance of doing linkage science. Low-latitude coronal holes
should also be observed for intermediate speed outflow. Even if this has a lower priority, it
should not be neglected since this is the typical solar wind observed at Earth.
• The minimum of the solar activity cycle is preferable and/or the declining phase for the polar
coronal holes. Low-latitude coronal holes can be observed at any phase of the cycle, but the
probability to be at the right longitude at perihelion is low.
The needed observations include:
• EUI/HRI in Coronal Hole mode with 1 min of cadence, during 1-2 hours and 12 hours at
lower cadence.
• PHI at high resolution at 1 min cadence.
• SPICE: FIP and velocity maps. A combination of composition, dynamics and a 30’’-wide
movie. 3.2 hours for SPICE, 3 times per day.
• In situ instruments: normal mode for the connection and burst modes for more details. MAG
Burst mode is required for ion cyclotron wave identification in order to distinguish from
different acceleration and heating mechanisms as well as small scale changes. Recent Helios
data analysis shows velocity changes with jet-like features probably well below the 40s
cadence. We would, therefore, need 3D distributions at 1s scales in order to determine the
properties inside and outside these features.
Other remarks:
• Additional observations from near-Earth assets are desirable but not required.
• EMC Quiet is required for linkage science. Noisy periods can be tolerated during the RS
observations, but EMC quiet is required approximately 12 hours later (depending on which
distance Solar Orbiter is).
• EPD, STIX are not required for this objective.
The possible remote sensing targets should include wide regions of well extended coronal holes, as
well as smaller regions for focusing on the following:
• Small cool coronal loops.
• Open magnetic funnels at the base of coronal holes (Tu et al., 2005; Schwadron and
McComas, 2003).
• Spicules with short lifetimes, fast motions, and hot plasma components. Macrospicules
reconnection and eruption within coronal holes (Yamauchi et al., 2005).
• Polar plumes (Deforest et al., 1997).
• Polar regions within plumes (“interplume planes”) (Giordano et al. 2000).
• Diverging polar regions of the extended corona.
• Coronal hole boundaries.
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The SOOP that addresses this objective is L_SMALL_HRES_HCAD_Fast_Wind, which also
addresses objective 1.1.1.2.
1.1.1.2 Origin of the small-scale X-ray and UV jets in polar coronal holes
Description of the objective:
One of the biggest discoveries of the Hinode satellite is the frequent occurrence of small-scale X-ray
and UV jets in polar coronal holes (Cirtain et al. 2007; Nistico et al., 2009; Krucker et al., 2011).
Such jets have widths between 2 × 103 and 2 × 104 km. Their origin is thought to be magnetic
reconnection of coronal field lines. At the reconnection site, Alfvén waves develop and produce
outflow velocities up to 800 kms-1, while the energy released by the reconnection heats the plasma
locally, generating mass motions with sonic speeds of ∼200 kms-1.
Given the high velocities and frequency of these events, it has been suggested that they contribute to
the fast solar wind. Their relations to the photospheric magnetic field, or the relevance of
photospheric processes for triggering them, have not been established yet. The high latitudes at
which they are observed hamper the accurate determination of their photospheric footpoints either
from the ground or NEO. Observations with remote-sensing instruments aboard Solar Orbiter will
allow us to understand the relation of these jets from a unique vantage point, providing comparable
high-angular resolution data simultaneously in the corona, chromosphere, and photosphere.
Remarks:
This objective could also be considered as part of the previous one (1.1.1.1). It only appears in a
separate section because it mostly requires remote sensing observations including STIX. From the
operational point of view, it can be jointly addressed with 1.1.1.1 with the
SOOP L_SMALL_HRES_HCAD_Fast_Wind. Observation of a sufficiently extended coronal hole
is required in order to catch multiple jets.
1.1.2 Source regions of the slow solar wind
1.1.2.1 Does slow solar wind originate from the over-expanded edges of coronal holes
(Antonucci et al., 2006)?
Description of the objective:
Study of the anti-correlation between the speed of the solar wind and the expansion rate of the
magnetic field. Wang and Sheeley (1990), Arge and Pizzo (2000), and others observe an anti-
correlation between solar wind speed and the so-called expansion factor of a flux tube near the Sun
calculated with a Potential-Field-Source-Surface (PFSS) model. Near the edges of coronal holes, the
flux-tube expansion factor is larger and generates a slower solar wind (Antonucci et al., 2006). The
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physical connection between flux-tube expansion and solar wind speed is unclear. Larger flux-tube
expansion increases the downward electron and proton heat flux (e.g., Cranmer et al., 2007) and
creates an energy sink that slows the solar wind (Schwadron and McComas, 2003). However,
detailed solar models (e.g., Lie-Svendsen et al., 2002) have difficulty accounting for the solar wind
properties under these circumstances. If slow solar wind has a coronal hole origin, we expect a
smooth transition across the coronal hole boundary. Boundary layers outside the coronal hole
(Schwadron et al., 2005) need to be correlated with their remotely observed source
structures. Furthermore, an “anomalous” slow solar wind is observed to be originating from the
boundaries of coronal holes with the same level of Alfvénic fluctuations as in the nearby fast wind,
but of considerable smaller amplitude (Antonucci et al., 2005; D’Amicis et al., 2015).
Remarks:
• This objective can be addressed with two combined observing strategies:
o At first, several (2-3) days of disc center pointing for overall coronal hole
configuration during a perihelion window that would scan over many latitudes. PHI
and EUI/FSI will give us a ~3D view of the coronal hole edges and METIS can
observe close to the Sun, while SPICE will do composition mapping scanning.
o At a second phase, the coronal hole boundary will be mapped with SPICE mosaics
during 1 day (with METIS being off). The SPICE raster area should be optimized to
make sure that both open and closed field boundary is captured. The choice of the
lines has to be optimized depending on the type of target.
• An alternative strategy would be to observe with SoloHI and METIS during a high-latitude
window, before trying to do the linkage science at the following perihelion window.
• The connectivity being of main importance for the solar wind origin objectives, we will need
modeling in order to choose the right pointing.
• Observations during two successive windows (high-latitude – especially for high-latitude
coronal holes - and the following perihelion) are preferred.
• It is easier to perform these observations during the minimum or declining phase of solar
activity, when we have a strong magnetic dipole. Even though it will not be easy to choose
where to point during the maximum, it is worth planning it as well at least once during the
maximum.
• Earth context observations are important for better constraining the models and improving
pointing decisions.
• Parker Solar Probe joint observations at a radial alignment would be beneficial, especially if
PSP happens to be at its perihelion.
The SOOP that addresses this objective is L_BOTH_HRES_LCAD_CH_Boundary_Expansion,
which was specifically created for this objective and does not currently cover any other sub-
objectives. All other slow wind origin objectives are addressed by the similar SOOP
L_SMALL_HRES_HCAD_SlowWindConnection.
1.1.2.2 Does slow and intermediate solar wind originate from coronal loops outside of coronal
holes?
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Description of the objective:
Solar observations show that the elemental composition, temperatures and thus the charge states of
large coronal loops outside of coronal holes are similar to the composition of the slow solar wind
(Feldman et al. 2005; Baker et al., 2013; Brooks et al., 2015). Loop source models require that the
foot-points of open field lines are interspersed with large loops outside coronal holes (Fisk et al.,
1999; Fisk and Schwadron, 2001). Interchange reconnection between open field lines and loops
(along topological boundaries: quasi-separatrix layers, including coronal nulls, e.g. van Driel-
Gesztelyi et al., 2012) releases the material stored on the loops and generates the slow wind.
Observations indicate the legs or tips of the helmet streamer, loops near active regions, and 80-300
Mm loops may all contribute to the slow wind. The challenge is to associate definitively the features
and structures in the slow wind with the morphology of the coronal complex. A loop origin for the
slow wind would show sharp interfaces and characteristic variations that can be correlated to the
remotely observed sources. Other interesting aspects to be explored include the existence of narrow
regions of open field lines between multiple closed field configurations within the streamer belts
(Noci et al., 1997; Wang et al., 2000; Noci and Gavryuseva, 2007), slow wind coming from small
coronal holes near or inside active regions (Wang, 2017, 10.3847/1538-4357/aa706e) vs from open
magnetic flux rooted in active regions (Liewer et al., 2004, 10.1007/s11207-004-1105-z), the
differences with the solar activity maximum when the slow wind appears to emanate rather from
small coronal holes and active regions (e.g. Neugebauer et al., 2002), the possibility of turbulent
reconnection (Rapazzo et al., 2012), the quasi-steady flow from streamer legs (Suess and Nerney,
2006), the transient contribution to the slow wind from parcels of plasma escaping from the
streamer core (Suess et al., 2009), explore the sharp interfaces in the slow wind indicative of loop
origin, observe the fans at the edge of active regions.
Remarks:
• This objective corresponds to different kinds of structures (helmet streamer, loops near
active regions, streamer core, belts, edge of active regions...). These will be different targets
observed by the same SOOP.
• As for all solar wind origin objectives, we need to compare the elemental composition,
temperatures and charge states of large coronal loops to the in situ slow wind ones (Feldman
et al., 2005; Baker et al., 2013; Brooks et al., 2015).
• Several days of observations of a large region surrounding an emerging active region are
required, it is probably interesting to observe during a full window for studying the
abundance of active regions.
• Observations near the perihelion are preferred for an increased resolution and a higher
probability for connectivity.
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• For the high-resolution observations, EUI needs to support SPICE with at least one or a few
high-resolution frames for each SPICE raster.
• As for all solar wind connectivity objectives, an EMC Quiet period is needed some hours
after the remote sensing observations. Note, however, that this is not easy to effectively plan
in practice since this would decrease the duration of the remote sensing observation of the
active region.
• Parker Solar Probe joint observations at a radial alignment would be beneficial, especially if
PSP happens to be at its perihelion.
• Reconnection events at the boundaries of a coronal hole caused by the emergence of a new
active region occur on timescales of the order of 1 hour (Rappazzo et al. 2012). Edmondson
et al. (2010) also found timescales of the order of 1-2 hours for interchange reconnection
processes between an active region loop and an adjacent coronal hole. The typical network
supergranule reconfiguration time is of the order of 1-2 days. This timescale can also be
adopted as the reconfiguration time of the global magnetic field lines. The magnetic dipole
emerging rate in coronal holes is 1-2/day (Abramenko et al. 2006). Schrijver at al. (1998)
found that flux concentrations are enhanced and disappear with a characteristic timescale of
about 1.5 days. Fisk and Schwadron (2001) state that the characteristic time for a change in
the open flux, due to reconnection with loops at low latitudes, is about 36-38 hr. On the other
hand, Antiochos et al. (2011) describe the large-scale field evolution as approximated to a
sequence of topologically smooth quasi-steady states (quasi-steady models). Therefore, we
are facing with phenomena occurring on timescales ranging from about 1 hour, if we
consider a single active region, to about 36 hr, when we consider the global coronal
magnetic field. On this basis, a PHI cadence of 6 hr is enough for investigating interchange
reconnection phenomena at a global scale (larger than the active region FoV). A shorter
cadence (about 1 hr) might be required to investigate at high spatial resolution interchange
reconnection phenomena between an active region and an adjacent coronal hole.
This objective is addressed by the SOOP L_SMALL_HRES_HCAD_SlowWindConnection. The
similar SOOP L_SMALL_MRES_MCAD_Connection_Mosaic can be used if we want to cover a
wider area around the active region with lower resolution, in particular for the SPICE FoV.
1.1.2.3 Abundance of minor ions as a function of height in the corona as an indicator of slow
or fast wind.
Description of the objective:
Abundance of minor ions as a function of height in the corona as indicator of slow or fast wind
(Antonucci et al., 2006).
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Remarks:
• This is a SPICE-led observation. METIS cannot contribute to the study of the abundance,
but it can give the context of the coronal configuration before the limb off-pointing or farther
out.
• After METIS observations, we need to off-point to the limb in order to observe either a limb
active region and/or along the boundary of a streamer.
• SPICE will observe with a narrow slit in order to enhance the resolution for properly
measuring the abundance. A typical duration would be of 3.2 h.
• In the case of streamers, a bigger height can be studied from farther out, it is better to
observe from a distance >0.55 AU (for METIS support). In the case of active regions, the
perihelion is preferred for higher resolution.
• It might be difficult to address this objective at the maximum of solar activity.
• This observation should be tested at least once to see if it works and depending on the
success, re-plan it in the future.
The SOOP R_SMALL_HRES_LCAD_Composition_vs_Height is specifically designed for (and
currently only serves) this objective.
1.1.2.4 Study of density fluctuations in the extended corona as a function of the outflow
velocity of the solar wind while evolving in the heliosphere
Description of the objective:
Study of density fluctuations in the extended corona as a function of the outflow velocity of the
solar wind while evolving in the heliosphere (Telloni et al., 2009).
Remarks:
• A complete and detailed study of coronal density fluctuations requires a long temporal
baseline (more than one day).
• EUI/FSI 174 would need deep exposures in order to observe up to METIS FoV with good
SNR.
• METIS should observe with the FLUCTS mode for about 1 hour and then turn to MAGTOP
(with a cadence of 5-20 minutes). The current SOOP is defined with cycles of 8 hours, but
multiple cycles should be planned in order to observe for at least one day. Another
possibility would be to observe 8 hours per window for 3 consecutive windows, but this
might not meet the minimum observation duration for the lower frequencies.
• This objective should better be addressed at perihelion for a better METIS SNR for the
FLUCTS mode combined to a lower cadence during the rest of the window.
• If SPICE participates, we will need to do limb pointing for a certain period of time and then
return to disc center for METIS.
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• PSP in quadrature would be an asset for observing the fluctuations in situ.
The SOOP R_FULL_HRES_HCAD_Density_Fluctuations is specifically designed for (and
currently only serves) this objective.
1.1.2.5 Structure and evolution of streamers
Description of the objective:
The extended observation at the limb of the same portion of the corona is essential to address the
study of streamer structure, evolution, and dynamics. This will focus either on the quiescent, slowly
evolving streamer belt or on the rapidly varying active streamers. The investigation of the
interrelation of solar wind acceleration region and magnetic topology of the flux tube guiding the
expansion necessitates as well long-term observations. The opportunity to freeze a streamer section
at the limb offers, in addition, the possibility of increasing dramatically the statistics, being, in this
case, only limited by the intrinsic evolution of the structure and not by solar rotation. A significant
increase in statistics is then coupled with the possibility of observing at high spatial resolution; this
allows us to resolve the fine structure and relevant dynamics in the slow wind coronal source
regions. This kind of studies requires the simultaneous knowledge of the electron density and
morphology of the corona, by observations in visible light, and of the radial flow velocities,
obtained with better accuracy and detail from images of the ultraviolet H I and He II emission, by
Doppler dimming techniques.
High-latitude observations will greatly impact on the study of large-scale structures of the solar
corona. If this phase occurs during solar minimum, from high latitudes the Orbiter views the
streamer belt running at low-latitude, or close to the equator, as an approximately continuous
annular structure around the solar disk. Furthermore, the out-of-ecliptic vantage point is indeed
providing another means to observe the corona not affected by solar rotation. When the morphology
is relatively simple and the relevant coronal features are at low-latitude, or close to the equator, their
intrinsic evolution can indeed be easily separated from rotational effects if viewed from high-
latitudes.
Another interesting aspect concerns the possibility to study the dynamics of coronal expansion all
around the streamer belt. During solar minimum, slow solar wind studies would be privileged, since
their supposed low-latitude and equatorial sources are predominant on the plane of the sky. It would
then be possible to (i) assess the contribution to the slow wind of sporadic reconnection events, such
as the coronal blowouts, and (ii) evaluate the total mass and magnetic flux injection into the
heliosphere all along the streamer belt.
Therefore, we have to:
• search for evidence of pseudo-streamers in the solar wind,
• study the detailed structure of the pseudo-streamers and
• study if slow wind is originating from pseudo-streamers.
Remarks:
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• Noci & Gavryuseva (2007) observed a quiescent streamer at solar minimum with UVCS for
4 consecutive days and determined the velocity pattern in it. Several days should be enough,
but several hours of observations could also give us useful insight.
• We should not observe at perihelion since we need METIS and a rather big FoV.
• A declining phase of the solar activity is preferred for a fully developed pseudo-streamer.
• Quadrature with Earth for combining with Solar Orbiter RS observations with L1 in situ, or
Earth coronagraphs observations with Solar Orbiter in situ.
This objective is addressed by the SOOP L_FULL_HRES_HCAD_Coronal_Dynamics, which is
aimed at observing structures in the outer corona and linking them to the heliosphere observed in
situ. METIS and SPICE are leading this SOOP, while in situ payload provides continuous
observations. Synoptic support from other remote sensing instruments is provided. Disk center
pointing is preferred.
1.1.2.6 Disentangle the spatial and temporal variability of the solar wind
Disentangle the spatial and temporal variability of the slow wind: comparison of the variability of
the slow wind on a number of orbits, with varying relative spacecraft-Sun longitudinal speeds at the
same distance. This objective is part of the general characterization of the solar wind and can be
addressed by the in-situ instruments. No specific operational planning is needed. However, it will
benefit from joint observations with PSP and L1, especially during radial alignments.
This objective is covered by the general in situ SOOP I_DEFAULT.
(To be moved to an overarching section (to be created) of characterization of the solar wind
together with other general in situ objectives).
1.1.2.7 Trace streamer blobs and other structures through the outer corona and the
heliosphere
Description of the objective:
• Trace streamer blobs ("helmet streamer plasmoids") and other structures through the outer
corona and the heliosphere (Sheeley and Rouillard, 2010).
• Study periodic density structures (Viall and Vourlidas, 2015) in the low corona and for
different times in the solar cycle.
The helmet streamer plasmoids are:
• observed in white light images as the continual, episodic releases of plasma from the tips of
helmet streamers
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• also, observed in situ at 1 AU and tracked from the upper corona to 1 AU (swept up and
compressed by the fast solar wind from low latitude coronal holes) (Sheeley and Rouillard,
2010).
• small interplanetary transients (size of 0.05-0.12 AU, magnetic field rotation between 2 and
13 hours) can be traced back to streamer events, but also to CMEs (Rouillard et al., 2011)
• flux ropes (Sheeley et al., 2009; Rouillard et al., 2011)
• exclusively plasma sheet phenomena (Wang et al., 2000)
• rate of ~4/day or approximately every 6 hours
• large in scale, initial sizes about 1 Rs in the radial direction and 0.1 Rs in the transverse
direction (Sheeley et al., 1997)
• originate at about 3-4 Rs (Sheeley et al., 1997)
• speeds: 150 km/s at 5 Rs, 300 km/s at 25 Rs
• thought to be released through either interchange reconnection, and/or complete
disconnection, and in either case, the reconnection takes place at high altitudes (Wang et al.
2000; Zurbuchen 2001; Crooker et al. 2004; Suess et al. 2009).
Other periodic density structures are:
• often not flux ropes (Viall et al., 2009)
• observed in 70%-80% of the slow solar wind and in much of the ecliptic fast wind (Viall et
al., 2008)
• rates of several minutes, characteristic timescale of 90 minutes, with a range of 65-100
minutes (Viall and Vourlidas, 2015)
• small scales (70-3000Mm)
• largest scale structures often contain smaller ones embedded within them (Viall et al., 2010)
• always associated with streamers and the HCS (Viall and Vourlidas, 2015)
• formed at or below 2.5 Rs, no evidence of formation above these heights (Viall and
Vourlidas, 2015).
Remarks:
• In general, not easy to distinguish radial vs transverse scales of structures. Joint observations
with PSP will be very important.
• At least one perihelion would be required for this objective for looking at the birth of the
structures. METIS would better observe from 0.3-0.35 during a zooming-out (North)
window.
• Low-latitude preferred.
• Minimum activity preferred, but the maximum could also be interesting if we could manage
to disentangle other effects.
• For linkage science, we would need Earth in quadrature and either observe in situ on Solar
Orbiter what Earth sees or observe in situ at 1 AU what SoloHI sees. For SoloHI to be
useful, we need a GSE –Y quadrature.
• The SoloHI turbulence mode should be used at perihelion for looking at the formation,
acceleration, shape, and evolution of the fastest periodic structures. For making sure that we
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will see at least one blob at its birth (frequency of 1 every 6 hours), the turbulence mode (20
sec) should be used for approximately 12 hours (or 1 day).
This objective is addressed by the SOOP L_FULL_HRES_HCAD_Coronal_Dynamics, which is
aimed at observing structures in the outer corona and linking them to the heliosphere observed in
situ. METIS and SPICE are leading this SOOP, while in situ payload provides continuous
observations. Synoptic support from other remote sensing instruments is provided. Disk center
pointing is preferred.
Other SOOPs are partially addressing the connectivity aspects of the objective (but not the
formation and evolution of the fastest structures). These are
L_SMALL_MRES_MCAD_Connection_Mosaic,
L_SMALL_HRES_HCAD_SlowWindConnection, L_BOTH_MRES_MCAD_Farside_Connection.
1.1.2.8 Determine the velocity, acceleration profile and the mass of the transient slow wind
flows
This objective is part of the general characterization of the solar wind and can be addressed by the
in situ instruments. No specific operational planning is needed. However, it will benefit from joint
observations with PSP and L1, especially during radial alignments.
This objective is covered by the general in situ SOOP I_DEFAULT.
(To be moved to an overarching section (to be created) of characterization of the solar wind
together with other general in situ objectives).
1.1.3 Source regions of the heliospheric magnetic field
1.1.3.1 Full characterization of the photospheric magnetic fields and fine structures
Description of the objective:
• Full characterization of photospheric magnetic fields: the magnetic field at the photosphere
can be determined quantitatively by recording the full polarization state of light in
appropriate spectral lines. Such measurements can be inverted to provide the full magnetic
vector at the photosphere and the LOS component of the plasma flow velocity. These
measurements will allow the emergence of magnetic flux to be determined, as well as its
redistribution through its interaction with convection. From surface maps of the magnetic
vector, it is possible to extrapolate the field into the Sun’s upper atmosphere, where its
evolution gives rise to numerous dynamic and energetic phenomena. Time series of the
velocity maps of the photosphere also allow a reconstruction of the subsurface structure
using the techniques of helioseismology. Observations of the photospheric fields are, thus,
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essential for studying both the generation and atmospheric evolution of solar magnetic
fields.
• Probe fine-scale structure of the solar photosphere, including waves and the emergence,
evolution, and dynamics of magnetic flux. This needs PHI observations at high spatial
resolution.
• Uncover the effects of waves and magnetic field changes on the upper atmospheric layers
(e.g. production of transient events (Galsgaard et al., 2005)).
• Study of the magnetic fields at the poles of the Sun.
o The polar fields are responsible for the polar coronal holes and largely drive the fast
solar wind, but are poorly known.
o Polar plumes are bright structures reaching far into the corona, which appear to
harbor gas moving slowly, compared to the fast solar wind in the interplume regions.
Observations of the magnetic field at their footpoints from high latitude will be
crucial for an understanding of their origin. The aim is to provide sufficiently
accurate and detailed magnetic maps, unhampered by the massive foreshortening that
current observations suffer from, to allow high-quality extrapolations of the field.
• Origin of the spicules and other chromospheric features (De Pontieu et al., 2004). Polar
spicules (Johannesson and Zirin, 1996). Spicules are a prominent chromospheric
phenomenon, cool and dense fibrils that intermittently connect the photosphere with the hot
and rarefied corona. They are short-lived (5-10 min), narrow (diameters less than 500 km)
and display upward motions with speeds up to 20 km/s. If this is really mass motion, then
the mass flux in spicules is 100 times larger than that of the solar wind. It has been shown
that photospheric p-modes, which are evanescent in the field-free photosphere and
chromosphere, can indeed propagate into and through the chromosphere if they are guided
by inclined magnetic field lines (De Pontieu et al. 2004). Due to the steep vertical density
gradient, the oscillations develop into shocks which may result in significant excursions of
the top of the chromosphere, i.e. cause spicules. According to De Pontieu et al. (2004), the
crucial ingredients for spicule formation are the photospheric velocities, the temperature
stratification and the inclination of the magnetic field lines. However, it is well known that
polar spicules are larger than ordinary spicules (Johannesson & Zirin 1996). They may,
therefore, differ in cause. Viewing them from out of the ecliptic will advantageously reveal
the underlying differences.
Remarks:
• Except for PHI/HRT (1 min cadence), EUI/HRI should observe at very high cadence (1-30s,
possible for approximately up to 20 minutes), maybe interleaved with burst modes with high
cadence (0.1 s, during 1 to a few seconds). The FoVs can be released if need be for telemetry
reasons (for instance down to ¼ of the PHI/HRT FoV).
• SPICE could participate for a dynamics study with many rasters and a similar FoV.
Alternatively, it could operate at a sit-and-stare mode.
• No other instruments are required.
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• Due to its demanding nature in terms of telemetry, the relevant SOOP cannot be planned for
long durations. It should be repeated whenever the telemetry budget configuration is
favorable.
This objective is covered by the SOOP
R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure run for different targets at the quiet
Sun (observed from the perihelion) and the coronal holes (observed from high-latitude windows).
The SOOP L_SMALL_HRES_HCAD_SlowWindConnection will also provide valuable
information on photospheric structures but at lower spatial resolution.
1.1.3.2 How does the Sun’s magnetic field link into space?
For any plasma to leave the Sun and reach space, it must travel along the magnetic field. Any study
of how the Sun exerts its influence in space must, therefore, address the connection of the solar
magnetic field far into the solar system.
The large scale structure of the interplanetary magnetic field (IMF) is well known (e.g. Mariani and
Neubauer, 1990; Smith, 2008): the Sun's rotation winds the field into the Parker spiral; compression
and rarefaction of the plasma in co-rotating interaction regions (CIRs) produces increases and
decreases of the field strength; the polarity of the solar source field is reflected in that of the IMF;
and the field is pervaded by waves and turbulence over a wide range of scales. Over the Sun's 22
years magnetic cycle, the IMF reflects the changing character of the solar field, from approximate
dipole to a much more complex, multi-pole structure.
However, the mapping between solar and interplanetary fields is only known on relatively large
scales and in a crude manner. Observations of the Sun's surface photospheric magnetic field
combined with coronal observations or MHD models of the corona, make it possible to estimate the
mapping between the lower corona and the “source surface” at several solar radii, but many simple
questions remain about how the Sun's magnetic field opens into space (e.g. Antiochos et al., 2007),
particularly with regard to the emergence of new coronal holes and the long-range connectivity of
active regions, as well as how the IMF disconnects from the Sun. Distant observations by Ulysses
over the Sun's poles have helped to constrain such mappings (e.g. Hoeksema, 1995; Forsyth et al.,
1997) but Solar Orbiter, being much closer to the Sun and hence eliminating many of the
uncertainties caused by local stream-stream interactions, will dramatically improve the precision
with which this can be constrained.
Beyond the source surface, MHD models must be used. These models will be greatly constrained by
Solar Orbiter magnetic field data, with the important consequence of improving the systematic
prediction ability of such models throughout the heliosphere. This connection is also essential for
many elements of the Solar Orbiter science objectives of linking solar and interplanetary
phenomena.
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1.1.3.2.1 How does the Sun’s magnetic field change over time?
Description of the objective:
The process by which the Sun's field reverses at solar maximum is highly complex: lack of
observations of the polar solar field greatly hampers our understanding of this process. Ulysses in
situ observations contributed to measurements of the previous cycle's reversal (e.g. Jones et al.,
2003). Solar Orbiter will image the polar field with PHI while simultaneously measuring the field in
space at a range of locations, making more precise measurements of the solar reversal and its effects
on the heliosphere. While the nature of the solar field reversal is part of Objective 4 (How does the
solar dynamo work and drive connections between the Sun and the heliosphere?), here we are
interested in how this affects the coronal and heliospheric magnetic fields.
We will measure how the polarity and large scale structure of the Sun's magnetic field in
interplanetary space varies close to the Sun as it moves from solar minimum towards maximum and
the global field reverses.
Even though part of this objective can be addressed with in situ only measurements, it would make
more sense to have access to full disk remote sensing observations in order to properly understand
how the in-situ changes are linked to the overall changes in the solar field. For this, we need long-
term observations at synoptic modes (for telemetry reasons), by targeting the full disk for both
photospheric and coronal fields. Since such (low resolution) observations already exist from Earth,
it would mostly be interesting to observe when Solar Orbiter is at the far side of the Sun or for
intermediate regions between an alignment with the Earth and the far side (regular spacing in
longitude). This should also be repeated for different latitudes. Since in situ MAG measurements are
key for this objective, good statistics during EMC Quiet periods are required.
This objective can be addressed by SOOP L_FULL_HRES_LCAD_MagnFieldConfig, which has
been mainly created for this goal.
1.1.3.2.2 How is the heliospheric current sheet (HCS) related to coronal structure?
Description of the objective:
The Heliospheric Current Sheet (HCS) is the interplanetary extension of the neutral line between the
two magnetic polarities in the corona and is a topological magnetic boundary. The HCS is also vital
in the motion of cosmic rays throughout the heliosphere: depending on the polarity of the solar
cycle, ions or electrons tend to migrate to low latitudes and along the HCS as they enter the solar
system.
The HCS is remarkably thin: just a few thousand km across (Zhou et al., 2005), but is surrounded by
the much thicker, denser Heliospheric Plasma Sheet (HPS) (Vourlidas and Riley, 2007). It is not
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clear how thin the HCS and HPS are close to the Sun, which could provide clues to their origin – the
HCS seems actually to thin with distance, for example. Both the HCS and HPS are also highly
variable: what is the origin of this variability? Reconnection appears to occur here (Gosling et al.,
2006): how frequent is this close to the Sun?
Data from a number of sources (cosmic rays: Simpson et al., 1996; the IMF: Luhmann et al., 1988;
geomagnetic activity: Mursala and Zieger, 2001; source surface models: Zhao et al., 2005) suggest
that the Sun's neutral line, and also the HCS, are persistently displaced Southward during solar
minimum. This is consistent with a solar dipolar field offset from the Sun's center, but why (or even
if) this should occur is unknown.
We, therefore, have to:
• Determine the local tilt and latitudinal extent of the HCS as a function of solar distance and
relation to the time-varying source surface neutral line. Asymmetry of the HCS. Determine
whether the HCS is offset from the equator in near-Sun space.
• Study the relation of the Heliospheric Plasma Sheet (HPS) to the HCS. (Vourlidas and Riley,
2007).
• Measure the variability of the HCS and HPS in time and space, with the goal of determining
the solar or local origin of the variations. Reconnection (Gosling et al., 2006)?
• Determine the link between heliospheric plasma sheet and coronal streamers (Bavassano et
al., 1997).
Remarks:
• As for objective 1.1.3.2.1 How does the Sun's magnetic field change over time?, in situ
observations are key during EMC Quiet periods, but full disk synoptic remote sensing
observations will also be important for a better understanding.
• This objective can be addressed by SOOP L_FULL_HRES_LCAD_MagnFieldConfig.
• It would be interesting to study the HCS crossings at the perihelia during the Cruise Phase.
• Observations during one solar rotation will not be enough since we will have gaps due to the
RS windows durations. High-latitude observations are also required.
1.1.3.2.3 How does the heliospheric magnetic field disconnect from the Sun?
Description of the objective:
At any given time, magnetic field lines stretch out from the Sun deep into interplanetary space,
carried by the solar wind. However, these fields must eventually disconnect from the Sun, resulting
in a complex, tangled magnetic topology in the heliosphere, with important consequences for
energetic particle propagation. It is essential to quantify the large-scale connectivity, but this is
remarkably difficult. The magnetic field polarity is an important indicator of the solar region of
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origin of a packet of solar wind, but cannot determine its connectivity alone. Signatures of
connectivity such as suprathermal electron streaming (e.g. Zurbuchen and Richardson, 2006) are
difficult to interpret. Indeed, (Owens and Crooker, 2007) suggests that electron scattering and fading
can result in signatures which look like disconnected field lines, but are actually due to interchange
reconnection. Only by simultaneously measuring the streaming electrons and magnetic field polarity
over a wide range of distances, and particularly close to the Sun, can we distinguish dis- and re-
connections from electron fade-out.
Evidence of reconnection at solar wind magnetic discontinuities (e.g. Gosling et al., 2007)
demonstrates that the connectivity of the IMF can change in space. How often does this occur close
to the Sun, where the solar wind is much more dynamic than at 1 AU? There is also evidence for
folds in the magnetic field, from cross-helicity (Balogh et al., 1999) and proton-alpha streaming data
(Yamauchi et al., 2004): what is their origin? Are they related to chromospheric reconnection
features such as jets (Shibata et al., 2007) or velocity shears (Landi et al., 2006)?
Remarks:
We have to search for local reconnection events using MAG and SWA in order to determine their
radial distribution and significance for the connectivity of the solar wind.
We will use signatures of the varying connectivity of the solar wind (e.g. suprathermal electrons),
combined with the magnetic field orientation and other measures of source polarity (e.g.
alpha/proton streaming and the normalized cross helicity) to search for bends, folds and small scale
polarity reversals. In this way, we will determine the small-scale polarity structure within coronal
hole flows and its relation to the global field.
This is an in situ objective and can be addressed with the I_DEFAULT SOOP as well as during the
connectivity operating plans for the fast and slow winds,
e.g. L_SMALL_HRES_HCAD_Fast_Wind, L_SMALL_HRES_HCAD_SlowWindConnection, L_
SMALL_MRES_MCAD_Connection_Mosaic.
1.1.3.3 What is the distribution of the open magnetic flux?
Description of the objective:
• Understand what determines the amount of open flux from the Sun, how open field lines are
distributed at the solar surface at any given time, and how these open field lines reconnect
and change their connection across the solar surface in time.
• Origin of the open magnetic flux from:
▪
▪ Coronal holes.
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▪ Active regions.
▪ Quiet Sun.
▪ Loops of varying heights.
▪ Tips and legs of the helmet streamer.
▪ Polar plumes.
▪ Other solar structures.
Remarks:
• We need MAG measurements for the magnetic field polarity and strength, over a wide range
of distances and latitudes, combined with signatures of connectivity such as suprathermal
electrons (SWA/EAS) to determine the amount of open magnetic flux in the heliosphere
from solar minimum to maximum.
• METIS can provide electron density and outflow velocity maps with a 20-30 min cadence.
• EUI can support with FSI for complementing observations in the lower lying corona, at low
cadence as METIS.
• PHI can provide low-resolution synoptic context.
• As for all goals of the category 1.1.3 Source regions of the heliospheric magnetic field, far
side observations will be most innovative, as well as high-latitude. EMC Quiet periods are
required for the MAG measurements.
Part of this objective can be addressed with the I_DEFAULT SOOP as well as during the
connectivity operating plans for the fast and slow winds,
e.g. L_SMALL_HRES_HCAD_Fast_Wind, L_SMALL_HRES_HCAD_SlowWindConnection, L_
SMALL_MRES_MCAD_Connection_Mosaic.
However up to now, we have only included it in R_SMALL_MRES_MCAD_AR_LongTerm, but
this would only partially address it.
1.1.4 Transverse themes (reconnection)
This section will be modified or completely removed by integrating the different parts in other
objectives. For this document, I keep it empty for now.
3.2.1.2 1.2 What mechanisms heat the corona and heat and accelerate the solar wind?
Present state of knowledge:
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Despite more than a half-century of study, the basic physical processes responsible for heating the
million-degree corona and accelerating the solar wind are still not known. Identification of these
processes is important for understanding the origins and impacts of space weather and to make
progress in fundamental stellar astrophysics.
Ultimately, the problem of solar wind acceleration is a question of the transfer, storage, and
dissipation of the abundant energy present in the solar convective flows. The key question is to
establish how a small fraction of that energy is transformed into magnetic and thermal energy above
the photosphere. Both emerging magnetic flux and the constant convective shaking and tangling of
magnetic field lines already threading the corona contribute to the processing of the energy in what
is an extremely structured, highly dynamic region of the solar atmosphere, the route to dissipation
involving cascading turbulence, current sheet collapse and reconnection, shocks, high-frequency
waves, and wave-particle interactions. The advent of high-cadence high-resolution observations has
demonstrated the extremely complex phenomenology of the energy flux in the lower atmosphere,
including many types of transient events discovered and classified by Yohkoh, SOHO, TRACE,
RHESSI, Hinode and, most recently, SDO.
Energy deposited in the corona is lost in the form of conduction, radiation (negligible in coronal
holes), gravitational enthalpy, and kinetic energy fluxes into the accelerating solar wind plasma.
Transition region pressures, coronal densities and temperatures, and the asymptotic solar wind speed
are sensitive functions of the mode and location of energy deposition. The mass flux is not,
however, as it depends only on the amplitude of the energy flux (Hansteen and Leer 1995). A
relatively constant coronal energy flux therefore explains the small variations in mass flux between
slow and fast solar wind found by Ulysses during its first two orbits, although the dramatic decrease
in mass flux over the last cycle points also to a decreased efficiency of coronal heating and therefore
to its dependence on the solar magnetic field (McComas et al. 2008; Schwadron and McComas
2008).
One of the fundamental experimental facts that has been difficult to account for theoretically is that
the fast solar wind originates in regions where the electron temperature and densities are low, while
the slow solar wind comes from hotter regions of the corona. The anticorrelation of solar wind
speed with electron temperature is confirmed by the anti-correlation between wind speed and
‘freezing in’ temperature of the different ionization states of heavy ions in the solar wind (Geiss et
al. 1995) and implies that the electron pressure gradient does not play a major role in the
acceleration of the fast wind. On the other hand, the speed of the solar wind is positively correlated
with the in-situ proton temperature, and the fastest and least collisionally coupled wind streams also
contain the largest distribution function anisotropies. Observations of the very high temperatures
and anisotropies of coronal heavy ions suggest that other processes such as magnetic mirror and
wave-particle interactions should also contribute strongly to the expansion of the fast wind (Li et al.
1998; Kohl et al. 1997, 1998, 2006; Dodero et al. 1998). In particular, either the direct generation of
high-frequency waves close to the cyclotron resonance of ions or the turbulent cascade of energy to
those frequencies should play an important role.
Theoretical attempts to develop self-consistent models of fast solar wind acceleration have followed
two somewhat different paths. First, there are models in which the convection-driven jostling of
magnetic flux tubes in the photosphere drives wave-like fluctuations that propagate up into the
extended corona. The waves partially reflect back toward the Sun, develop into strong turbulence,
and/or dissipate over a range of heights. These models also tend to attribute the differences between
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the fast and slow solar wind not to any major differences in the lower boundary conditions, but to
the varying expansion factor of magnetic field lines in different areas of coronal holes (Cranmer et
al. 2007).
In the second class of models, the interchange reconnection models, the energy flux usually results
from magnetic reconnection between closed, loop-like magnetic flux systems (which are in the
process of emerging, fragmenting, and being otherwise jostled by convection) and the open flux
tubes that connect to the solar wind. Here the differences between fast and slow solar wind result
from qualitatively different rates of flux emergence, reconnection, and coronal heating in different
regions on the Sun (Axford and McKenzie 1992; Fisk et al. 1999; Schwadron and McComas, 2003).
It has been difficult to evaluate competing models of fast wind acceleration and to assess
observationally the relative contributions of locally emerging magnetic fields and waves to the heat
input and pressure required to accelerate the wind largely because of the absence of measurements
of the solar wind close to the Sun where they can be mapped with sufficient precision to a solar
source region.
How Solar Orbiter will address this question:
Solar Orbiter’s combination of high-resolution measurements of the photospheric magnetic field
together with images and spectra at unprecedented spatial resolution will make it possible to identify
plasma processes such as reconnection/shock formation and wave dissipation in rapidly varying
surface features, observe Doppler shifts of the generated upflows, and determine
compositional signatures. Whatever the scale, magnetic reconnection leads to particle dissipative
heating and acceleration and wave generation, which have the net effect of a local kinetic energy
increase in the lower solar atmosphere that can be revealed through high-resolution extreme
ultraviolet (EUV) imaging and spectroscopy. Wave propagation will be traced from the source site
to the region of dissipation through observations of EUV-line broadening and Doppler shifts.
Global maps of the H outflow velocity, obtained by applying the Doppler dimming technique to the
resonantly scattered component of the most intense emission line of the outer corona (H I 121.6),
will provide the contours of the maximum coronal expansion velocity gradient for the major
component of the solar wind, and the role of high-frequency cyclotron waves will be
comprehensively assessed by measuring spectroscopically the particle velocity distribution across
the field and determining the height where the maximum gradient of outflow velocity occurs
(Telloni et al. 2007).
Solar Orbiter’s heliospheric imager will measure the velocity, acceleration, and mass density of
structures in the accelerating wind, allowing precise comparison with the different acceleration
profiles of turbulence- driven and interchange reconnection-driven solar wind models.
As it is performing imaging and spectroscopic observations of the corona and photosphere, Solar
Orbiter will simultaneously measure in situ the properties of the solar wind emanating from the
source regions. The in-situ instrumentation will determine all of the properties predicted by solar
wind acceleration models: speed, mass flux, composition, charge states, and wave amplitudes.
Moving relatively slowly over the solar surface near perihelion, Solar Orbiter will measure how
properties of the solar wind vary depending on the changing properties of its source region, as a
function of both space and time, distinguishing between competing models of solar wind
generation.
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1.2.1 What mechanisms heat the corona?
The detailed objectives for understanding the heating of the corona include, but are not limited to,
the following:
1.2.1.1 Energy flux in the lower atmosphere
Description of the objective:
• Understand dissipation involving cascading turbulence, current sheet collapse and
reconnection, shocks, high-frequency waves and wave-particle interactions. Difference
between resonant and stochastic?
• Understand oscillations in small-scale flux tubes (Jess et al., 2009; Martínez González, 2011;
Stangalini et al., 2013; Requerey et al., 2015).
The scientific aim is to characterize the properties of waves in the photosphere and their coupling
with the upper atmosphere, chromosphere, and corona.
Waves are one clear mechanism for transferring energy from the photosphere to the chromosphere
and corona. Measuring the properties of the waves requires, in part, a determination of the velocity
field. The line-of-sight velocity component can be determined at different heights in the atmosphere
by observing Doppler shifts in different spectral lines. From the Earth’s vantage point we have high-
resolution ground-based, balloon-borne, and satellite instruments. Determining the horizontal
velocity has previously relied on using correlation tracking of intensity variations and rely on the
questionable assumption that the changes in the location of the brightness fluctuations reflect the
actual velocity. The orbit and capability to measure Doppler velocities, in conjunction with existing
and upcoming ground-based or near-earth observatories, offers the unique chance to directly
measure two components of the velocity field using the Doppler effect.
Remarks:
High-resolution co-temporal measurements including Doppler velocity maps from Solar Orbiter as
well as ground and NEOs are required. In particular, the ground-based and NEOs should include
high-resolution Doppler images in the same line (with a higher cadence than that of SO), as well as
lines sampling different heights of the atmosphere. Co-observation with IRIS would be desirable.
During the observing period, the Earth-Sun-SO angle should be between 30◦ and 60◦ – a range
which represents a compromise between determining the two components of the velocity field and
allowing magnetic features which can act as wave guides to be partially resolved.
For ease of understanding the connection between the different heights, the observations would best
be performed at the center of the disk as observed from the Earth (where observations over different
wavelengths are possible). Because also the achievable cadence will be higher on ground than
with PHI, it is preferable to select targets which are closer to the disk center as seen from Earth and
at higher heliocentric angles as seen from Solar Orbiter. The highest possible cadence is desirable,
and a shorter time series (of down to 30 minutes of Solar Orbiter observations) would still allow the
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scientific objectives to be met. (The ground-based and NEO should be made for a period of 90
minutes centered on the 30 minute SO observations). However, in order to guarantee reliable
conditions (seeing) at the coordinating ground-based observing facility (e.g. DKIST) a continuous
high-cadence observation period of several hours is required.
High-resolution context magnetic maps from SO immediately before and after the 30 minute
observing window are required to provide context and aid co-alignment.
A second observational campaign of an area 45◦ from disk center, with an Earth-Sun-SO angle of
90◦, would be desirable.
This objective can be partly covered by the high-resolution high-cadence SOOP that studies the fine
scale of the photosphere R_SMALL_HRES_HCAD_RSburst and the
SOOP R_SMALL_HRES_HCAD_WaveStereoscopy, which has been specifically designed for
studying the properties of the waves in the photosphere.
The SOOPs should be run for a bright source (e.g. active region but also quiet sun) at the perihelion
for a duration from 20 minutes to several hours. It would be preferable to address this objective
earlier in the mission, because of the Lyα degradation with time.
1.2.1.2 Energy and mass flux in the corona. Loss in form of conduction, radiation, gravitational
enthalpy, and kinetic energy fluxes into the accelerating solar wind plasma.
Description of the objective:
Energy and mass flux in the corona. Loss in form of conduction, radiation, gravitational energy,
enthalpy, and kinetic energy fluxes into the accelerating solar wind plasma.
This goal needs better definition and detailed description.
1.2.1.3 Contribution of flare-like events on all scales
Description of the objective:
• Statistical study of more than 70k flare-like events on all scales that are expected to be
seen during the NMP for better defining the flares’ distribution that possibly contains
enough energy to heat the corona (Parker 1983; Hudson, 1991).
• Reduce the lower energy limit of detectable flares.
• Understand the mechanisms responsible for producing nanoflares (field line shear, braiding,
coronal tectonics, (Priest, 2003)).
• Search for small 3He-rich events close to the perihelion.
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1.2.1.4 Observe and explore flare-like ‘heating events’ from the quiet corona
• Observe and explore flare-like ‘heating events’ from the quiet corona (Krucker et al., 1997).
• Possibility to observe or not hard X-ray counterparts not yet detected (Hannah et al., 2007).
1.2.1.5 Determine whether coronal heating is spatially localized or uniform, and time steady or
transient or impulsive for a wide range of magnetic loops with different spatial scales. Observe
coronal nanoflares in active region moss (Testa et al., 2013).
1.2.1.6 Resolve the geometry of fine elemental loop strands (Aschwanden et al., 2002; Reale et al.,
2007) and define the possible multi-temperature nature of such structures and their transversal
expansion in the corona, giving insight to the level of the tangling of the magnetic flux tubes (Lopez
Fuentes et al., 2006).
1.2.1.7 Detect and characterize waves in closed and open structures (e.g. De Pontieu et al., 2007),
looking for signatures of their dissipation in the transition region and low corona, with a
corresponding evaluation of the energy released to heat the solar plasma.
1.2.1.8 Investigate the role of small scale magnetic flux emergence in energizing the above laying
layers (e.g. Galsgaard et al., 2007).
1.2.1.9 Multiple-temperature diagnostics of flaring coronal loops (Battaglia and Kontar, 2013).
1.2.1.10 Heating in flaring loops vs in active regions.
1.2.2 What mechanisms heat and accelerate the solar wind?
1.2.2.1 Determine where energy is deposited in the solar wind (Cranmer et al. 1997; Telloni et al.
2007).
1.2.2.2 What drives the evolution of the solar wind distribution?
Description of the objective:
• High time resolution variability in the magnetic field associated with particle distributions in
different plasma parameter regimes (low and high plasma β, fast and slow wind) in order to
quantify the links between particle distribution evolution and wave-particle interactions
(Kasper et al., 2002; Matteini et al., 2007; Maksimovic et al., 2005).
• Role of the electron heat flux. Characterize the non-thermal character of the electron
distributions at perihelion and their evolution with heliocentric distance (Stverak et al., 2009;
Maksimovic et al., 2005).
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• Generation of non-thermal ion distributions (beams, drifting heavy ions, hot heavy ions,
proton ‘strahl’).
• Generation of non-thermal electron distributions (beams, drifting ‘strahl’).
• What role do local evolution (turbulence, shears), global evolution (expansion), and
collisions have in determining the properties of the proton distributions?
• Sub-Debye length electric fields (Randol & Christian, 2014), measure electric fields and
suprathermal proton tails.
1.2.2.3 What are the origins of waves, turbulence and small scale structures?
• Quantify the undisturbed waves and relate the wave power and other characteristics to the
source regions (by measuring photospheric motion in the regions from which the plasma
originated) (Matthaeus and Goldstein, 1986; Bruno and Carbone, 2005; Matthaeus et al.,
2007; Bello González et al., 2010).
• Identify and characterize the waves associated with the plasma instabilities that isotropize
and heat the solar wind (Hellinger et al., 2006; Matteini et al., 2010).
• Resonant absorption and emission by thermal particle distributions: role of the high-
frequency cyclotron waves.
• Ion energization processes in the solar wind (study of the electric fluctuations near the ion
cyclotron frequencies).
• Ion cyclotron resonance damping of the high-frequency part of the Alfvén spectrum (e.g.
Cranmer, 2002).
• Solve the problem of the mode conversion from Langmuir to electromagnetic waves (Bale et
al., 1998; Kellogg et al. 1999; Farrell et al., 2004; Ergun et al., 2008). Characterize the
energy balance between electron beams, Langmuir waves and electromagnetic radio waves
at several heliocentric distances.
• How do variations and structure in the solar wind affect low-frequency radio wave
propagation?
• Small scale structures such as solitons (Rees et al., 2006), mirror modes (Stevens and
Kasper, 2007), and draped fields and - in the case of dust trail signatures (Jones et al., 2003)
- confirm or refute their correlation with predicted trails.
• Study inbound waves in the corona (Verdini et al., 2009; DeForest et al., 2014).
1.2.2.4 Identify and characterize the solar wind reconnection physics in current sheets with
thickness down to the ion scales and smaller. Electron and ion physics near the current sheet.
Magnetic islands formation. Compare microphysics of solar wind reconnection with
magnetospheric reconnection.
1.2.2.5 Magnetic reconnection in the chromosphere, the transition region and the corona, driven by
magnetic field evolutionary processes. Explore reconnection signatures such as brightenings
(Harrison et al., 1997), flows and jets (Innes et al., 1997) and plasma evaporation (Klimchuk, 2006).
Reconnection between small closed loops and open structures supplying energy to the nascent solar
wind in the form of waves and turbulent flows (Tu et al., 2005).
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1.2.2.6 Study fast plasma flows from the edges of solar active regions discovered with Hinode/EIS
(Harra et al., 2008), which are driven by pressure gradient between reconnecting magnetic loops
(Baker et al., 2009; Del Zanna et al., 2011) and produce intermediate-speed solar wind streams (van
Driel-Gesztelyi et al., 2012). Study active region expansion (Uchida et al., 1992).
1.2.2.7 Study the correlation degree between velocity and magnetic field fluctuations in the
interplanetary space: may help in disentangling the source regions of the slow solar wind.
3.2.1.3 1.3 What are the sources of solar wind turbulence and how does it evolve?
Present state of knowledge:
The solar wind is filled with turbulence and instabilities. At large scales, the fast solar wind is
dominated by anti-sunward propagating Alfvén waves thought to be generated by photospheric
motions. At smaller scales, these waves decay and generate an active turbulent cascade, with a
spectrum similar to the Kolmogorov hydrodynamic scaling of f^-5/3. In the slow solar wind,
turbulence does not have a dominant Alfvénic component, and it is fully developed over all
measured scales. There is strong evidence that the cascade to smaller scales is anisotropic, but it is
not known how the anisotropy is generated or driven (Horbury et al. 2008). What do the differences
between the turbulence observed in the fast wind and that observed in the slow wind reveal about
the sources of the turbulence and of the wind itself?
Little is known about what drives the evolution of solar wind turbulence. Slow-fast wind shears,
fine-scale structures, and gradients are all candidate mechanisms (Tu and Marsch 1990; Breech et
al. 2008). To determine how the plasma environment affects the dynamical evolution of solar wind
turbulence it is essential to measure the plasma and magnetic field fluctuations in the solar wind as
close to the Sun as possible, before the effects of mechanisms such as velocity shear become
significant, and then to observe how the turbulence evolves with heliocentric distance.
The dissipation of energy in a turbulent cascade contributes to the heating of the solar wind plasma.
However, while measurements of the properties of solar wind turbulence in near-Earth orbit largely
agree with observed heating rates (Smith et al. 2001; Marino et al. 2008), the details are
controversial and dependent on precise models of turbulent dynamics. In order to establish a full
energy budget for the solar wind, the heating rates as a function of distance and stream properties
must be determined, including turbulence levels before the cascade develops significantly.
The statistical analysis of the fluctuating fields also reveals pervasive fine-scale structure (e.g.,
discontinuities and pressure balanced structures). The origin of these structures is uncertain: are they
the remnant of complex coronal structuring in the form of strands of small-scale flux tubes advected
by the solar wind flow (Borovsky 2008; Bruno et al. 2001), or are they generated locally by
turbulent fluctuations?
At scales around the proton gyroradius and below, turbulent fluctuations interact directly with the
solar wind ions. The precise nature of the turbulent cascade below the proton gyroradius is poorly
understood and might even vary depending on local plasma conditions. Below the electron
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gyroradius, conditions are even less certain and the partitioning of turbulent energy into electron or
ion heating is unknown at this time. In addition, solar wind expansion constantly drives distribution
functions toward kinetic instabilities, where fluctuations with characteristic signatures are generated
(e.g., Marsch 2006). What physical role do kinetic effects play with distance from the Sun? What
role do wave-particle interactions play in accelerating the fast solar wind? What contribution do
minor ions make to the turbulent energy density in near-Sun space?
How Solar Orbiter will address this question:
Solar Orbiter will measure waves and turbulence in the solar corona and solar wind over a wide
range of latitudes and distances, including closer to the Sun than ever before, making it possible to
study turbulence before it is significantly affected by stream-stream interactions. By traveling over a
range of distances, the spacecraft will determine how the turbulence evolves and is driven as it is
carried anti-sunward by the solar wind flow.
Detailed in-situ data will make it possible to distinguish between competing theories of turbulent
dissipation and heating mechanisms in a range of plasma environments and are thus of critical
importance for advancing our understanding of coronal heating and of the role of turbulence in
stellar winds.
By entering near-corotation close to the Sun, Solar Orbiter will be able to distinguish between the
radial, longitudinal, and temporal scales of small-scale structures, determining whether they are the
signatures of embedded flux tubes or are generated by local turbulence.
Solar Orbiter’s magnetic and electric field measurements, combined with measurement of the full
distribution functions of the protons and electrons will fully characterize plasma turbulence over all
physically relevant time scales from very low frequencies to above the electron gyrofrequency.
Because Solar Orbiter is a three-axis stabilized spacecraft, it can continuously view the solar wind
beam with its proton instrument, measuring proton distributions at the gyroperiod and hence making
it possible directly to diagnose wave-particle interactions in ways that are not possible on spinning
spacecraft. By traveling closer to the Sun than ever before, it will measure wave-particle interactions
before the particle distributions have fully thermalized, studying the same processes that occur in
the corona. By measuring how the distributions and waves change with solar distance and between
solar wind streams with different plasma properties, Solar Orbiter will make it possible to determine
the relative effects of instabilities and turbulence in heating the plasma.
The solar wind is the only available plasma laboratory where detailed studies of
magnetohydrodynamic (MHD) turbulence can be carried out free from interference with spatial
boundaries and in the important domain of very large magnetic Reynolds numbers. A detailed
comparison between experimental in-situ data and theoretical concepts will provide a more solid
physical foundation for MHD turbulence theory, which will be of critical importance for
understanding the solar coronal heating mechanism and the role or turbulence in the solar wind.
1.3.1 Solar and local origin of Alfvénic fluctuations
Description of the objective:
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• Outward propagating stochastic fluctuations, Alfvénic modes within the fast streams
(Belcher and Davis, 1971; Tu and Marsch, 1995).
o Excitement by magnetic activity and reconnection in the chromospheric network.
o Large-scale torsional Alfvén waves generated from magnetic restructuring related to
interchange reconnection (Lynch et al., 2014).
o Alfvén waves in the low corona caused by leakage of wave power from
chromospheric oscillations (Tomczyk et al., 2007).
o Presence or not of large-amplitude Alfvénic fluctuations with ⎟B⎜=const (Matteini
et al., 2014) in the inner heliosphere, implication for their stability and their role as a
source for turbulence.
• Identification of the drivers at the "outer scale" (large scale) end of the turbulent cascade.
• Inward propagating Alfvén waves (already active at 0.3 AU as observed by Helios):
o Local production by velocity shears (Roberts et al., 1992).
o Parametric decay mechanisms (Malara et al., 2000).
o Generation of inward component by expansion (relevant for polar wind where shear
effects are negligible) (Velli et al., 1989; Verdini et al., 2009).
o Nonlinear generation by compressive MHD cascade (Marsch & Mangeney, 1987).
o Carefully distinguish real local production of inward modes from compressive events
masking a generation process (Bavassano and Bruno, 1989).
• How does the behavior of MHD waves observed remotely in the core of coronal holes and at
the CH/QS boundaries match what is observed in situ in the solar wind that emanates from
these regions?
Remarks:
• Note that the McIntosh paper (McIntosh et al., Nature, 475, 477) observed waves in
structures seen above the limb (spicules) and in active region loops (on-disc). In coronal
holes seen on the disc, it will be probably easier to observe oscillating structures in Lyman
alpha (HRI) than in the other passbands. In Ly-alpha, there are dark, fibril-like structures
that are seen with a better contrast with respect to the background (optical thickness), while
in 304 or 174, the structures are optically thin and the contrast may be less favorable. But
using HRI and SPICE will probably require off-pointing, so we should implement two
different observational strategies: one with and one without off-pointing.
• For the objective of linking with in-situ observations, the "minimal" option should be
without SPICE or EUI/HRI, and no off-pointing required. EUI/FSI 174 and 304 at 30-s
cadence during 1-2 hours minimum (there is currently no program fitting this cadence with
FSI). Alternating both filters would probably require filter wheel motions incompatible with
EMC requirements. If so, FSI 304 may be preferred. Or 1 hour of FSI 304 followed by 1
hour of FSI 174. To save telemetry, it is possible to store only 1/4 of the FOV as long as it
covers the core and one boundary of the CH. The 30-s cadence is sufficient to observe waves
as in McIntosh et al., Nature, 475, 477. But the spatial resolution of FSI may not be
sufficient to resolve fine structures and their motions, especially for observations on the disc,
which are necessary as the objective is to measure the wind emanating from the RS
observations.
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• "Optimal" option: with SPICE and HRI; off-pointing and mosaic may be necessary to
observe the source region of the wind that will be measured at the spacecraft within a few
days from the RS observations. The observed region must cover both the core of the CH and
at least one of the CH/QS boundaries.
• HRI 174 and HRI Ly-alpha at 30-s cadence during 1-2 hours (EUI Coronal hole mode (C)
can fit). Observations in coronal holes require deep exposures (at least in HRI 174), which
requires minimal jitter.
• FSI 304 at 30-s cadence.
• SPICE rasters for measurement of the line widths (to determine the non-thermal motions,
supposed to correspond to LOS-integrated effect of Alfvén waves), for at least transition
region and coronal lines (optionally: also, lines for abundance measurements, for which the
total intensity may be enough and not the full profile). Cadence is not an issue as we
measure the integrated effect in the line width (for observation of the wave motions in
Doppler shifts, sit-and-stare (i.e. no slit scan) may be preferable).
1.3.2 How is turbulent energy dissipated and how does turbulence evolve within the heliosphere?
Description of the objective:
• Measure the turbulence dissipation range and understand its scaling with heliocentric
distance and plasma properties (Alexandrova et al., 2007; Bruno & Trenchi, 2014; Chen et
al., 2014). Separate the analysis between fast and slow wind and in particular for different
plasma beta.
• Distinguish between various heating and dissipation mechanisms. Various kinetic processes:
Alfvénic or magneto sonic damping, kinetic Alfvén waves (see above), whistler dispersion,
Hall MHD dispersive cascade.
• Study the evolution of the intermittency of the magnetic and plasma quantities (Bruno et al.,
2003; 2014).
• Study the evolution of the effective magnetic Reynolds number (Bruno et al., 2015).
• Study the evolution of the MHD rugged invariants (magnetic helicity, cross-helicity, and
residual energy).
• Relate the localization of the spectral break between the fluid and the kinetic regimes to the
amplitude of the fluctuations at MHD scales: are the scales at which dissipation mechanisms
become important related to the energy contained in the inertial range? (Bruno and Trenchi,
2014).
• Explore the radial evolution of the compressible and incompressible third order moment
scaling within the inertial range, in order to gain information on the status of the turbulent
cascade (Sorriso-Valvo et al., 2007), including: the evaluation of the mean energy transfer
rate, the development of the cascade relative to the balance between the inward/outward
fluctuations, the dependence on local parameters.
• Understand the origin and the radial evolution of the low frequency 1/f spectrum in fast solar
wind and its implications on the characteristics of turbulence, intermittency and dissipation.
• Dissipation at the corona.
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1.3.3 Plasma turbulence variability
Description of the objective:
Quantify the evolution of the turbulent spectrum with distance, heliomagnetic latitude and stream
structure in the inner heliosphere.
• Determine in what conditions the radial evolution is a matter of age of turbulence or in situ
heliospheric sources.
• Try to connect in situ spectrum with source spectrum at the Sun (together with sophisticated
models).
• Get insight into the observed correlation between speed, temperature and fluctuations
amplitude (Grappin et al., 1990) by comparing in situ properties and source region
properties.
1.3.4 Plasma turbulence anisotropy
Description of the objective:
Quantify the anisotropy of the turbulence in both slow and fast streams; determine its effects on
particles at a range of energies and relate it to the solar origin of the fluctuations (Horbury et al.,
2005; Dasso et al., 2005).
• Determine if different levels of anisotropy result from differences in the source region or
from slower/faster evolution in fast/slow streams (age of turbulence).
• Quantify the effect of expansion (radial evolution) on the 3D anisotropy of turbulence as
observed at 1AU (Narita et al., 2010) and predicted by MHD simulations (Dong et al.,
2014).
• Compare turbulence anisotropy to measured diffusion coefficients of energetic particles
during solar events (Horbury and Balogh, 2001; Bieber et al., 1996).
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3.2.2 Objective 2: How do solar transients drive heliospheric variability?
The dynamic Sun exhibits many forms of transient phenomena, such as flares, CMEs, eruptive
prominences, and shock waves. Many directly affect the structure and dynamics of the outflowing
solar wind and thereby also eventually affect Earth’s magnetosphere and upper atmosphere, with
significant consequences for society through hazards to, for example, space-based technology
systems and surface power systems. Understanding these impacts, with the ultimate aim of
predicting them, has received much attention during the past decade and a half under the banner of
‘space weather.’ However, many fundamental questions remain about the physics underpinning
these phenomena and their origins, and these questions must be answered before we can realistically
expect to be able to predict the occurrence of solar transients and their effects on geospace and the
heliosphere. These questions are also pertinent, within the framework of the ‘solar-stellar
connection,’ to our understanding of other stellar systems that exhibit transient behavior such as
flaring (e.g., Getman et al. 2008).
Solar Orbiter will provide a critical step forward in understanding the origin of solar transient
phenomena and their impact on the heliosphere. Located close to the solar sources of transients,
Solar Orbiter will be able both to determine the inputs to the heliosphere and to measure directly the
heliospheric consequences of eruptive events at distances close enough to sample the fields and
plasmas in their pristine state, prior to significant processing during their propagation to 1 AU. Solar
Orbiter will thus be a key augmentation to the chain of solar-terrestrial observatories in Earth orbit
and at the libration points, providing a critical perspective from its orbit close to the Sun and out of
the ecliptic.
In the following sections, we discuss in more detail three interrelated questions which flow down
from this top-level question:
How do CMEs evolve through the corona and inner heliosphere? How do CMEs contribute to solar
magnetic flux and helicity balance? How and where do shocks form in the corona and inner
heliosphere?
3.2.2.1 2.1 How do CMEs evolve through the corona and inner heliosphere?
Present state of knowledge:
Following earlier observations by space-based white-light coronagraphs, considerable progress in
understanding CMEs has been achieved using data from the SOHO mission, which provides
continuous coverage of the Sun and combines coronagraphs with an EUV imager and off-limb
spectrometer. Other spacecraft, such as ACE, WIND, Ulysses and STEREO, which carried
comprehensive in-situ instrumentation, have contributed significantly to our understanding of the
interplanetary manifestation of these events. With a full solar cycle of CME observations, the basic
features of CMEs are now understood. CMEs appear to originate from highly-sheared magnetic
field regions on the Sun known as lament channels, which support colder plasma condensations
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known as prominences. Eruptions develop in the low corona within 10-15 minutes, while the
associated shocks cross the solar disk within 1 hour. CMEs reach speeds of up to 3000 km/s and
carry energies (kinetic, thermal and magnetic) of ~1025 J (= 1032 ergs). They can also accelerate
rapidly during the very early stages of their formation, with the CME velocity being closely tied, in
time, to the associated flare’s soft X-ray light profile (Zhang and Dere 2006). High-resolution
SOHO and STEREO coronagraph images have provided evidence for a magnetic flux rope structure
in some CMEs as well as for post-CME current sheets. Both features are predicted by CME
initiation models (e.g., Lin & Forbes 2000; Lynch et al. 2004).
STEREO observations have made it possible to chart in three dimensions the trajectories of CMEs
in the corona and heliosphere, thereby improving our understanding of CME evolution and
propagation. STEREO data have supported detailed comparison both of in-situ measurements with
remote-sensing observations and of MHD heliospheric simulations with observations. The
combination of high-cadence coronagraphic and EUV imaging simplifies the separation of the CME
proper from its effects in the surrounding corona (Patsourakos and Vourlidas 2009) and allows a
more accurate determination of its dynamics.
Despite the advances in our understanding enabled by SOHO and STEREO, very basic questions
remain unanswered. These concern the source and initiation of eruptions, their early evolution, and
the heliospheric propagation of CMEs. All current CME models predict that the topology of ICMEs
is that of a twisted flux rope as a result of the flare reconnection that occurs behind the ejection.
Observations at 1 AU, however, find that less than half of all ICMEs, even those associated with
strong flares, have a flux rope structure (Richardson and Cane 2004). Many ICMEs at 1 AU appear
to have a complex magnetic structure with no clearly-defined topology. Moreover, for ICMEs that
do contain flux ropes, the orientation is often significantly different from that expected on the basis
of the orientation of the magnetic fields in the prospective source region. CMEs are believed to
originate from prominence eruptions, yet in ICMEs observed at 1 AU prominence plasma is very
rarely detected. These major disconnects between theoretical models (of prominence eruption and
CME propagation) and observations (remote and in situ) need to be resolved if any understanding of
the CME process is to be achieved.
How Solar Orbiter will address this question:
To advance our understanding of the structure of ICMEs and its relation to CMEs at the Sun beyond
what has been achieved with SOHO and STEREO requires a combination of remote-sensing and in-
situ measurements made at close perihelion and in near-corotation with the Sun. Through combined
observations with its magnetograph, imaging spectrograph, and soft X-ray imager, Solar Orbiter
will provide the data required to establish the properties of CMEs at the Sun and to determine how
coronal magnetic energy is released into CME kinetic energy, flare-associated thermal/non-thermal
particle acceleration, and heating. Observations with the imaging spectrograph will be used to
determine the composition of CMEs in the low corona and to establish how they expand and rotate
and will also provide vital clues to the energy partition within a CME once it is released. Solar
Orbiter will make comprehensive in-situ measurements of the fields and plasmas (particularly
composition) of ICMEs following their release and, critically, prior to their processing during
propagation in the heliosphere. These measurements will allow the properties of an ICME to be
related to those of the CME at the Sun and to the conditions in the CME source region as observed
by Solar Orbiter’s remote-sensing instruments and will make it possible to examine the evolution of
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CMEs in the inner heliosphere. Solar Orbiter’s combination of remote-sensing and in-situ
observations will also establish unambiguously the magnetic connectivity of the ICME and reveal
how the magnetic energy within flux ropes is dissipated to heat and accelerate the associated
particles. Solar Orbiter data will also reveal how the structure of the magnetic field at the front of a
CME evolves in the inner heliosphere – a critical link in understanding, and eventually predicting,
how transient events on the Sun may determine the geoeffective potential of the event.
To fully understand the physical system surrounding CME ejection, the temporal evolution of active
regions and CME-related shocks and current sheets must be tracked from their formation in the
corona to their expulsion in the solar wind. During the mission phases when the spacecraft is in
near-corotation with the Sun, Solar Orbiter will continuously observe individual active regions, free
from projection complications, over longer periods than are possible from Earth orbit. Solar Orbiter
will thus be able to monitor the development of sheared magnetic fields and neutral lines and to
trace the flux of magnetic energy into the corona. Observations from the vantage point of near-
corotation will make it possible to follow the evolution of the current sheet behind a CME with
unprecedented detail and to clarify the varying distribution of energy in different forms (heating,
particle acceleration, kinetic).
2.1.1 What are the global structure, initiation, and evolution of CMEs?
2.1.1.1 CME initiation
Description of the objective:
CME initiation has been a core space physics problem for the last three decades. The two current
paradigms are distinguished primarily by the topology of the pre-eruption magnetic field: twisted
flux rope (e.g., Roussev et al., 2004) or sheared arcade (Antiochos et al., 1994). Irrespective of the
pre-eruption topology, all models predict that as a result of the flare reconnection occurring below
the ejection, CMEs in the heliosphere must have a twisted flux rope topology, as commonly
observed (Gosling et al., 1995). If the pre-eruption topology is that of a twisted flux rope, then the
innermost part of its structure should exhibit relatively undisturbed filament plasma parameters.
However, if the twist forms only as a result of flare reconnection, then the whole twisted structure in
the heliosphere should exhibit the properties of flare-reconnection-heated plasma, hot beamed
electrons, high charge states of Fe, as well as compositional anomalies of heavy ions including He.
By measuring the electron and ion properties of a CME along with its magnetic structure, we
determine the pre-eruption topology and the initiation mechanism. Solar Orbiter will provide the
opportunity to perform these measurements near the Sun, minimizing propagation effects such as
internal reconnection, which homogenizes the CME structure (e.g., Lynch et al., 2005).
2.1.1.2 CME structure
Description of the objective:
In order to better understand the structure of CMEs, we have to explore the following:
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• Three-part structure: bright leading shell, dark cavity, bright core (Illing and Hundhausen,
1986; Plunkett et al., 2000).
• Presence of large-scale helical magnetic flux rope in the CME cavity (Vourlidas et al.,
2013).
• Counter-streaming electrons observations as signatures of field lines tied at both ends to the
Sun (Crooker and Owens, 2012).
• Electron heat flux dropouts as signatures of magnetic flux disconnection from the Sun
(Crooker and Owens, 2012).
• Longitudinal distribution of the emitted plasma and detection of geoeffective events.
2.1.1.3 CME evolution
Description of the objective:
In order to better understand the structure of CMEs, we have to explore the following:
• Self-similar expansion (Cremades and Bothmer, 2004; Gibson and Low, 1998).
• Cylindrical geometry with the axis of symmetry corresponding to the long axis of a large-
scale helical magnetic flux rope that originates in the CME source region (Thernisien et al.,
2006).
• Interaction with the background velocity and density structures (Odstrcil et al., 2005; Riley
et al., 2003; Colaninno et al., 2014).
• Aerodynamic drag force and equalization of ICME and solar wind speeds (Gosling and
Riley, 1996; Cargill, 2004; Vrsnak et al., 2012; Subramanian et al., 2012).
• Relationship between the three-part structure of a CME and the ICME counterparts
(magnetic cloud with a flux rope or ‘complex ejecta’ with disordered magnetic fields)
(Vourlidas et al., 2013).
• Bright front evolution to become the sheath of compressed solar wind.
• Dark cavity corresponds to the flux rope.
• Presence of an in situ ‘plug’ of cold, dense plasma trailing the flux rope, interpreted as
remnant material from the erupting filament (bright core).
• What is the fate of the erupting filament as the CME propagates?
o Only a small fraction escapes with the CME?
o Filament material flows back along magnetic field lines or falls back due to
Rayleigh-Taylor instability (Innes et al., 2012; Carlyle et al.,2014; van Driel-
Gesztelyi et al., 2014).
o Filament is present but has lost its expected low-charge state signature because of
heating (Skoug et al., 1999; Rakowski et al., 2007)?
• Relationship between the three-part structure of a CME and shocks in the low corona and in
situ.
• Interaction between CMEs and its effects on SEPs (Gopalswamy et al., 2004; Richardson et
al., 2003; Rodriguez-Pacheco et al., 2003).
• Dynamics in transient coronal holes and recovery phase of the eruption.
• Improve CME arrival time estimates and predictions of geomagnetic activity.
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3.2.2.2 2.2 How do CMEs contribute to solar magnetic flux and helicity balance?
Present state of knowledge:
Magnetic flux is transported into the heliosphere both by the solar wind, in the form of open flux
carried mostly by the fast wind from polar coronal holes, and by coronal mass ejections, which drag
closed flux with them as they propagate into the heliosphere. At some point, the closed flux
introduced by CMEs must be opened to avoid an unsustainable buildup of magnetic flux in the
heliosphere. Measurements of the magnetic flux content of the heliosphere from near the Earth
show that the total amount of magnetic flux in the solar system changes over the solar cycle.
Longer-term variations are also known to occur. Proxies such as geomagnetic activity and cosmic
ray fluxes provide evidence that the average IMF strength has increased substantially in the last 100
years, perhaps by as much as a factor of 2. Surprisingly, however, during the recent solar minimum,
the IMF strength is lower than at any time since the beginning of the space age.
The relative contribution of the solar wind and CMEs to the heliospheric magnetic flux budget is an
unresolved question, as is the process by which the flux added by the CMEs is removed. Models to
explain the solar cycle variation assume a background level of open flux, to which CMEs add extra
flux during solar maximum, increasing the intensity of the IMF. The exceptionally low intensity of
the IMF during the last minimum has been attributed to the low rate of CME occurrence [Owens et
al. 2008]. Alternatively, there may simply be no ‘background’ open flux level.
There is evidence that the flux introduced into the heliosphere by CMEs may be removed by
magnetic reconnection within the trailing edges of CMEs, which disconnects the CME from the Sun
or by interchange reconnection closer to the solar surface [e.g., Owens and Crooker 2006]. Recent
observations show that the reconnection process occurs quite often in the solar wind, even when the
magnetic field is not under compression. However, the rate and/or locations at which reconnection
generally removes open flux are not at present known.
Together with magnetic flux, the solar wind and CMEs carry magnetic helicity away from the sun.
Helicity is a fundamental property of magnetic fields in natural plasmas, where it plays a special
role because it is conserved not only by the ideal dynamics but also during the relaxation which
follows instabilities and dissipation. Helicity is injected into the corona when sunspots and active
regions emerge, via the twisting and braiding of magnetic flux. During the coronal heating process,
the overall helicity is conserved and tends to accumulate at the largest possible scales. It is natural to
assume that critical helicity thresholds may be involved in the triggering of CMEs, but how solar
eruptions depend on the relative amounts of energy and helicity injection during active region
emergence and evolution is unknown. Yet this understanding could be a crucial element in the
prediction of large solar events.
How Solar Orbiter will address this question:
Fundamental to the question of the contribution of CMEs to the heliospheric flux budget is the flux
content of individual events. Encountering CMEs close to the Sun before interplanetary dynamics
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affect their structure, Solar Orbiter will measure their magnetic flux content directly; comparisons
with remote-sensing measurements of their source regions will clarify the relation between CME
flux and the eruption process. As Solar Orbiter moves through the inner heliosphere, it will
encounter CMEs at different solar distances, making it possible to quantify the effect of
interplanetary dynamics on their apparent flux content.
The flux carried outwards by CMEs must eventually disconnect completely from the Sun, or
interchange reconnects with existing open field lines. Solar Orbiter will diagnose the magnetic
connectivity of the solar wind and CME plasma using suprathermal electron and energetic particle
measurements. These particles, which stream rapidly along the magnetic field from the Sun, indicate
whether a magnetic flux tube is connected to the Sun at one end, at both ends, or not at all. These
particles disappear when the field is completely disconnected, or may reverse their flow direction as
a result of interchange reconnection. However, scattering and reflection due to curved, tangled, or
compressed magnetic field lines act to smear out these signatures with increasing solar distance,
leading to ambiguity in connectivity measurements. Solar Orbiter, by traveling close to the Sun
before this scattering is significant, will determine the original level of magnetic connectivity;
covering a wide range of distances in the inner heliosphere, the spacecraft will measure how the
connectivity changes as field lines are carried away from the Sun.
Solar Orbiter will also directly sample reconnection regions in the solar wind as they pass the
spacecraft, determining their occurrence rates in the inner heliosphere as a function of distance and
testing theories of CME disconnection by searching for reconnection signatures in the tails of
CMEs.
The contribution to the heliospheric magnetic flux of small scale plasmoids, ejected from the tops of
streamers following reconnection events, is unclear. Solar Orbiter, slowly moving above the solar
surface during perihelion passes, will determine the magnetic structure, connectivity, and plasma
properties including the composition of these ejecta, using spectroscopic imaging observations to
unambiguously link them to their source regions.
To assess the role of CMEs in maintaining the solar magnetic helicity balance, Solar Orbiter will
compare the helicity content of active regions as determined from remote sensing of the
photospheric magnetic field with that of magnetic clouds measured in situ. Such a comparison
requires both extended remote-sensing observations of the same active region over the region’s
lifetime and in-situ measurements of magnetic clouds from a vantage point as close to the solar
source as possible. Around its perihelia, Solar Orbiter will ‘dwell’ over particular active regions and
observe the emergent flux for a longer interval (more than 22 days) than is possible from 1 AU,
where perspective effects complicate extended observations. The resulting data will be used to
calculate the helicity content of an active region, track its temporal variation, and determine the
change in helicity before and after the launch of any CMEs. Should a magnetic cloud result from an
eruptive event in the active region over which Solar Orbiter is dwelling, the relatively small
heliocentric distances between the solar source and the spacecraft will make it highly probable that
Solar Orbiter will directly encounter the magnetic cloud soon after its release. Determination of the
cloud’s properties and connectivity through Solar Orbiter’s in-situ particle-and-fields measurements
will enable the first-ever comparison of a magnetic cloud in a relatively unevolved state with the
properties of the solar source, an impossibility with measurements made at 1 AU. The comparison
of the helicity change in the source region with the value measured in the magnetic cloud will
provide insight into the role of CMEs in the helicity balance of the Sun.
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2.2.1 How do CMEs contribute to the global evolution of magnetic flux in the heliosphere?
Description of the objective:
• Temporal addition of closed flux by CMEs, which opens via interchange reconnection,
conserves open magnetic flux while accounting for the rise in the observed heliospheric field
strength near solar maximum (Owens and Crooker, 2006).
• CME induced interchange reconnection may also drive the reversal of the heliospheric open
field (Owens et al., 2006).
• For this we need to understand:
• How much axial magnetic flux is carried out by CMEs?
• The opening time of closed CME flux through interchange reconnection.
• The latitudinal separations of CME foot-points.
• How does the in situ ICME magnetic flux compare to photospheric magnetic flux?
2.2.2 What is the role of ICMEs in the Sun’s magnetic cycle?
Description of the objective:
• Measure the helicity of ICMEs leaving the Sun and link them directly to changes in the solar
field thereby quantifying the effect of ICMEs on the solar cycle (Owens et al., 2007).
• Study the evolution of ICMEs in the inner heliosphere to determine the effects of local
dynamics in changing their characteristics and apparent helicity, relating them directly to
observed solar events (Zurbuchen and Richardson, 2006).
3.2.2.3 2.3 How and where do shocks form in the corona?
Present state of knowledge:
The rapid expulsion of material during CMEs can drive shock waves in the corona and heliosphere.
Shocks in the lower corona can also be driven by flares, and in the case of CME/eruptive flare
events, it may be difficult to unambiguously identify the driver (Vršnak and Cliver 2008). CME-
driven shocks are of particular interest because of the central role they play in accelerating coronal
and solar wind particles to very high energies in SEP events.
Shocks form when the speed of the driver is super-Alfvénic. The formation and evolution of shocks
in the corona and the inner heliosphere thus depend (1) on the speed of the driving CME and (2) on
the Alfvén speed of the ambient plasma and its spatial and temporal variations. According to one
model of the radial distribution of the Alfvén speed in the corona near active regions, for example,
shocks can form essentially in two locations, in the middle corona (1.2-3 RSun), where there is an
Alfvén speed minimum, and distances beyond an Alfvén speed maximum at ~4 RSun (Mann et al.
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2003). A recent study of CMEs with and without type II radio bursts (indicative of shock formation)
has shown that some of the fast and wide CMEs observed produced no shock or only a weak shock
because they propagated through tenuous regions in the corona where the Alfvén velocity exceeded
that of the CME (Gopalswamy et al. 2008). CME shock formation/evolution can also be affected by
the interaction between an older, slower-moving CME and a faster CME that overtakes it.
Depending on the Alfvén speed in the former, the interaction may result in the strengthening or
weakening of an existing shock driven by the overtaking CME or, if there is no existing shock, the
formation of one (Gopalswamy 2001; 2002).
Studies of LASCO images obtained during the rising phase of solar cycle 23 have demonstrated the
feasibility of detecting CME-driven shocks from a few to ~20 RSun and of measuring their density
compression ratio and propagation direction (Vourlidas et al. 2003; Ontiveros and Vourlidas 2009).
This development has opened the way for the investigation of shock formation and evolution in the
lower corona and heliosphere through Solar Orbiter’s combination of remote-sensing observations
and in-situ measurements.
How Solar Orbiter will address this question:
Understanding shock generation and evolution in the inner heliosphere requires knowledge of the
spatial distribution and temporal variation of plasma parameters (density, temperature, and magnetic
field) throughout the corona. Solar Orbiter’s remote-sensing measurements – in particular electron
density maps derived from the polarized visible-light images and maps of the density and outflow
velocity of coronal hydrogen and helium – will provide much improved basic plasma models of the
corona, so that the Alfvén speed and magnetic field direction can be reconstructed over the distance
range from the Sun to the spacecraft. Remote sensing will also provide observations of shock
drivers, such as flares (location, intensity, thermal/non-thermal electron populations, time-profiles),
and manifestations of CMEs (waves, dimmings, etc.) in the low corona with a spatial resolution of a
few hundred kilometers and cadence of a few seconds. It will measure the acceleration profile of the
latter and then track the CMEs through the crucial heights for shock formation (2-10 RSun) and
provide speed, acceleration, and shock compression ratio measurements.
Type II bursts, detected by Solar Orbiter, will indicate shock-accelerated electron beams produced
by the passage of a CME and thus provide warning of an approaching shock to the in-situ
instruments. These in-situ plasma and magnetic field measurements will fully characterize the
upstream and downstream plasma and magnetic field properties and quantify their microphysical
properties, such as turbulence levels and transient electric fields (while also directly measuring any
SEPs). Spacecraft potential measurements also allow for rapid determinations of the plasma density,
and of electric and magnetic field fluctuations, on microphysical scales, comparable to the Doppler-
shifted ion scales, which are characteristic of the spatial scales of shocks. The evolution of such
parameters will provide insight into the processes dissipating shock fronts throughout the range of
magnetic/velocity/density and pressure parameter space. Because of Solar Orbiter’s close proximity
to the Sun, the measurements of the solar wind plasma, electric field, and magnetic field will be
unspoiled by the dynamical wind interaction pressure effects due to solar rotation and will provide
the first reliable data on the magnetosonic speed, the spatial variation of the plasma pressure and
magnetic field in the inner heliosphere. MHD modeling studies have shown that interactions among
recurring CMEs and their shocks occur typically in the distance range around 0.2-0.5 AU (Lugaz et
al. 2005). Solar Orbiter will spend significant time in the regions of recurring CME interactions and
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so will be able to investigate the effects of such interactions on the evolution of CME-driven
shocks.
2.3.1 Coronal shocks
Description of the objective:
• Identify coronal shocks and characterize their spatial distribution and outward propagation
velocity.
• Study interaction with coronal plasma.
• Characterize the longitudinal distribution of coronal shocks during high latitude orbits.
2.3.2 What are the properties and distribution of heliospheric shocks?
2.3.2.1 Understand coronal conditions under which the shocks form and determine the
interplanetary conditions where they evolve
2.3.2.2 Identify interplanetary shocks and characterize their spatial and temporal evolution
• Identify interplanetary shocks and characterize their spatial and temporal evolution.
• Characterize shock structure and determine whether they are locally quasi-perpendicular or
quasi-parallel, and study kinetic properties of shock-related waves and turbulence that
control ion scattering mean free paths near the shock.
2.3.2.3 Study heating and dissipation mechanisms at shocks with radial distance
2.3.2.4 Identify mechanisms that heat the thermal solar wind particle populations near shocks and
determine their energy partition
2.3.3 Resolve the interplanetary shock field and plasma structure down to the spatial and temporal
scales comparable and smaller than the typical ion scales.
2.3.4 Shock-surfing acceleration mechanism
Description of the objective:
• Measure the relative strength of the convection vs cross-shock electric field as a function of
heliocentric distance.
• Fine-scale structure of shocks with fast density measurements (using satellite potential
measurements).
2.3.5 Understand the radio emissions from the ICME driven shocks
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Description of the objective:
• Type II and III radio bursts.
• Langmuir waves and electromagnetic mode conversion.
• Characterize the energy balance between electron beams, Langmuir waves and
electromagnetic radio waves at several heliocentric distances.
2.3.6 Identify shock accelerated particles
• Infer the rates of particle acceleration and the injection energies as a function of distance and
along different parts of the shock.
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3.2.3 Objective 3: How do solar eruptions produce energetic particle radiation that fills the
heliosphere?
Astrophysical sites throughout the solar system and galaxy have the universal ability to
accelerate ions and electrons to high speeds, forming energetic particle radiation. Detected remotely
from radio and light emission around supernovae remnants, the Sun, and planets, or directly from
particles that reach our detectors, this radiation arises from the explosive release of stored energy
that can cause magnetic fields to rearrange, or can launch shock waves which accelerate particles by
repeatedly imparting many small boosts to their speed. The nearly universal occurrence of energetic
particle radiation, along with the effects it can have on planetary environments, evolution of life
forms, and space systems has fostered a broad interest in this phenomenon that has long made it a
high priority area of investigation in space science. Since remote sites in the galaxy cannot be
studied directly, solar system sources of energetic particles give the best opportunity for studying all
aspects of this complex problem.
The Sun is the most powerful particle accelerator in the solar system, routinely producing
energetic particle radiation at speeds close to the speed of light, sufficiently energetic to be detected
at ground level on Earth even under the protection of our magnetic field and atmosphere. SEP
events can severely affect space hardware, disrupt radio communications, and cause re-routing of
commercial air traffic away from polar regions. In addition to large events, which occur roughly
monthly during periods of high sunspot count, more numerous, smaller solar events can occur by
thousands each year, providing multiple opportunities to understand the physical processes
involved. In the following sections, we have divided in more detail three interrelated questions that
flow down from this top-level question: How and where are energetic particles accelerated at the
Sun? How are energetic particles released from their sources and distributed in space and time?
What are the seed populations for energetic particles?
3.2.3.1 3.1 How and where are energetic particles accelerated at the Sun?
One of the two major physical mechanisms for energizing particles involves particles
interacting with moving or turbulent magnetic fields, gaining small amounts of energy at each
step and eventually reaching high energies. Called Fermi or stochastic acceleration, this mechanism
is believed to operate in shock waves and in turbulent regions such as those associated with
reconnecting magnetic fields or in heated coronal loops. The second major physical mechanism is a
magnetic field whose strength or configuration changes in time, producing an electric field which
can directly accelerate particles in a single step. At the Sun, such changes occur when large
magnetic loop structures reconnect or are explosively rearranged due to the stress from the motion
of their footpoints on the solar surface (e.g., Aschwanden 2006; Giacalone and Kota 2006).
Multiple processes may take place in SEP events, and while it is not possible to cleanly
separate them, they can be split into two broad classes, the first being events associated with shock
waves. As a CME moves into space, it drives a shock creating turbulence that accelerates SEPs from
a seed population of ions filling the interplanetary medium. Mixed into this may be particles from
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an associated solar flare. CMEs often accelerate particles for hours as they move away from the
Sun, and in some cases, are still accelerating particles when they pass Earth orbit in a day or two.
Since CMEs can be huge, it is easy to see how they can fill a large portion of the heliosphere with
SEPs. However, the correlation of the observed radiation intensities with CME properties is poor,
indicating that additional aspects of the mechanism such as seed populations or shock geometry
must play important roles that are not yet well understood (Gopalswamy 2006; Desai et al. 2006;
Mewaldt 2006).
The second class of events is associated with plasma and magnetic field processes in loops
and active regions that accelerate particles. Reconnecting magnetic loops, and emerging magnetic
flux regions provide sites for stochastic energetic particle acceleration or acceleration by electric
fields. Because these regions are relatively small, the acceleration process is quick: on the order of
seconds or minutes, but the resulting event is small and often difficult to observe. Since the
energized particles are in the relatively high-density regions of the corona, they collide with coronal
plasma, producing ultraviolet (UV) and X-ray signatures that make it possible to locate their
acceleration sites and probe the local plasma density. Most of these particles remain trapped in their
parent loops, traveling down the legs to the solar surface where they lose their energy to the ambient
material, producing X- and gamma-rays. A few escape on magnetic field lines leading to
interplanetary space, traceable by their (‘type III’) radio signatures, electrons, and highly
fractionated ion abundances where the rare 3He can be enhanced by 1000-10,000 times more than in
solar material.
The energetic particles from these events reach our detectors at Earth orbit after spiraling
around the IMF, which is an Archimedes spiral on average. But since the IMF meanders, and has
many kinks, the length of the particle’s path has a good deal of uncertainty, and the particles
themselves scatter and mix, smearing and blurring signatures of the acceleration at the Sun.
Although we can enumerate candidate mechanisms for producing SEPs, a critical question is: what
actually happens in nature? Which processes dominate? How can shocks form fast enough to
accelerate ions and electrons to relativistic energies in a matter of minutes?
3.1.0 Explore in depth the SEP properties
Our current knowledge of the properties of the two different kinds of SEP events, gradual and
impulsive can be summarized as follows:
Gradual SEP events:
▪ electrons up to tens of MeV, ions at GeV
▪ once per month near solar maximum
▪ dominated by protons (small e/p fluence ratios)
▪ variable composition and charge states
▪ extend over > 100° in solar longitude
▪ generally associated with large solar flares & fast CMEs
▪ believed to be due to acceleration by CME-driven shocks and not by flares.
Impulsive SEP events:
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• Low energy electrons: 1-100 keV
• Ions: 0.01-1 MeV/nucleon
• >104 events per year near solar maximum
• large e/p ratios
• enhanced abundances of 3He and of heavy ions
• high charge states
• enhanced alpha/proton ratio
• extend over 30° in solar longitude
• associated with type III bursts
• many clearly come from flares/microflares, but often not even a microflare or soft X-ray
burst is observed.
3.1.1 CME and shock associated SEP sources
Moving from the lower corona to the interplanetary medium, shocks evolve rapidly since the
sound speed drops as plasma density and magnetic field strength decline as ~1/r^2. Solar Orbiter’s
coronagraphs will remotely identify shock front location, speed, and compression ratios through
this critical region within ~10 RSun. Combining this information with local electron densities as
well as coronal ion velocities given by Solar Orbiter radio and light polarization observations will
provide critical constraints on shock evolution models in regions too close to the Sun for direct
sampling.
In the regions explored by Solar Orbiter close to the Sun, the IMF is almost radial with much
less variation (uncertainty) in length than is the case at 1 AU, so the knowledge of the actual path
length improves by a factor of 3-5 as the length shortens. Having observed the CMEs and their radio
signatures in the corona and the X-ray signatures of the energetic particles near the Sun, Solar
Orbiter will then determine the subsequent arrival time of the particles in situ that can be accurately
compared to CME position. As the shock then rolls past the spacecraft, Solar Orbiter will measure
the shock speed and strength as well as the associated plasma turbulence, electric, and magnetic
field fluctuations. This will give a complete description of the acceleration parameters in the inner
heliosphere where much of the particle acceleration takes place. Indirect evidence from 1 AU
indicates that shock acceleration properties depend on the longitude of the shock compared to
the observer; close to the Sun, Solar Orbiter can cleanly test this property since the IMF is nearly
radial, the CME lift-off site is known, and the accelerated particles will have little chance to mix. In
the high-latitude phase of the mission, Solar Orbiter will be able to look down on the longitudinal
extent of CMEs in visible, UV, and hard X-rays, allowing first direct observations of the
longitudinal size of the acceleration region. This will make it possible to test currently
unconstrained acceleration and transport models by using measured CME size, speed, and shape to
specify the accelerating shock.
3.1.1.1 Where and when are shocks most efficient in accelerating particles?
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Type II radio bursts show that CME-driven shocks form in the corona, but we do not know
when they start accelerating ions. We do know from energetic storm particle (ESP)/shock spike
events that shocks accelerate particles at 1 AU, but rarely to >~ a few 10s of MeV. As the shock
propagates and evolves, when is it most potent in accelerating particles? How is the acceleration
related to the changing magnetic geometry? Comparisons of the kinetic energy of fast CMEs
measured by SOHO/LASCO with ACE and GOES measurements of accelerated particle energy
spectra indicate that as much as 10% of the CME kinetic energy often goes into accelerated
particles. That is about the efficiency needed for the acceleration of galactic cosmic rays by
supernovae shocks.
Correlating the SEP/CME efficiency estimates with interplanetary conditions near the Sun
will probe why and how the CME acceleration efficiency is so variable. With the wide range of
other multi-point imaging and in situ data available, it will be possible to significantly improve on
estimates of the CME energy, and of the global energy budget for solar eruptions, including plasma,
magnetic field, energetic particle, and photon contributions. (e.g., Emslie et al. 2004).
There is now a well-developed theory of SEP shock acceleration and the capability to
perform sophisticated simulations of these events. The models, however, depend on assumptions
about conditions in the inner heliosphere, and many predictions of the models cannot be verified
using data from 1 AU. By going close to the Sun, within ~1-2 λ, SolO provides a unique opportunity
to probe shock acceleration of SEPs to high energies with powerful new in situ and imaging
diagnostic measurements. Only by measuring the accelerated particles and the plasma and magnetic
field properties close to the acceleration source can we test and improve current models of the SEP
acceleration and transport.
We, therefore, need to explore the following questions:
• As the shock propagates and evolves, when is it most potent in accelerating particles?
o When/where does the shock start accelerating ions?
• How is the acceleration related to the changing magnetic geometry?
• Improve SEP kinetic energy estimates (variable CME acceleration efficiency) (Mewaldt et
al., 2005)
o As much as 10% of the CME kinetic energy often goes into accelerated particles.
o Show why the efficiency is so variable (by better constraining CME energies?).
• What is the relation between shock acceleration, turbulence properties, and anomalous
diffusion mechanisms? (Perri and Zimbardo, 2012)
3.1.1.2 Why are gradual SEP events so variable?
One of the most intriguing questions about gradual events is their great variability in
intensity: while the CME speed varies by a factor of 4, the SEP intensities of >20 MeV protons vary
by 3 to 4 orders of magnitude for a given CME speed (Kahler and Vourlidas 2005). In addition,
there are wide variations in the composition of gradual SEP events, including Fe/O ratios that vary
by a factor of ~100 (Mewaldt et al. 2006).
Observations from solar cycle 23 suggest that CME-driven shocks accelerate mainly
suprathermal ions rather than the bulk solar wind (Mason et al. 1999). The intensity and
composition of suprathermal ions with 10 to 100 keV/nuc is known to be much more variable than
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the solar wind. One reason for this variability is that suprathermal particles are believed to arise
from a variety of sources, including previous gradual and impulsive SEP events, ubiquitous micro-
and nano-flaring at the Sun, and interplanetary acceleration of solar wind by stochastic processes
(Fisk and Gloeckler 2006). In a steady state solution, particles are assumed to be continually
injected at the shock and a power law spectrum is achieved when there is enough time. However, at
a propagating CME-driven shock, the amount of time available for acceleration is limited. Thus, for
the acceleration to successfully operate, it is crucial that there are enough seed particles.
3.1.1.2.1 Intensity variability
• Study and explain the SEP intensity variability (Kahler and Vourlidas, 2005): while the
CME speed varies by a factor of 4, the SEP intensities of >20 MeV protons vary by 3 to 4
orders of magnitude for a given CME speed.
3.1.1.2.2 Composition variations
The intensity and composition of suprathermal ions with 10 to 100 keV/nuc is known to be
much more variable than the solar wind. One reason for this variability is that suprathermal particles
are believed to arise from a variety of sources, including previous gradual and impulsive SEP
events, ubiquitous micro- and nano-flaring at the Sun, and interplanetary acceleration of solar wind
by stochastic processes (Fisk and Gloeckler 2006).
In a steady state solution, particles are assumed to be continually injected at the shock and a
power law spectrum is achieved when there is enough time. However, at a propagating CME-driven
shock, the amount of time available for acceleration is limited. Thus, for the acceleration to
successfully operate, it is crucial that there are enough seed particles.
We need observations of the suprathermals and SEPs before and during CMEs, to know
what material is actually accelerated. Simultaneous observations of the near Sun seed particles with
SPP would be very helpful.
3.1.1.2.3 Warped shock fronts
3.1.1.2.4 Turbulence and inhomogeneities
3.1.1.3 How are energetic particles accelerated continuously in the corona and solar wind?
Observations of quiet-time suprathermal ions and electrons (the "superhalo") at 1 AU and
beyond find that the suprathermal tails of the solar wind are always present, suggesting that
continuous acceleration is taking place in the solar wind and/or in the corona. Fisk and Gloeckler
(2006) have suggested that the ubiquitous suprathermal ion tails can result from thermal particles
interacting with compressional turbulence embedded in the solar wind. This mechanism predicts a
power-law steady state spectrum with an index of -1.5, consistent with a number of ion observations
at 1 AU and in the outer heliosphere. However, the mechanism of the acceleration takes time, so it
is likely that the suprathermal spectrum will not be fully developed and more variable inside 1 AU.
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3.1.1.4 How can SEPs be accelerated to high energies so rapidly?
The January 20, 2005 event provided a wake-up call to space weather forecasters when the
>100 MeV proton intensities reached their highest levels in more than 15 years. This event was also
the 2nd largest ground-level event of all time, reaching peak intensity only 18 minutes after the
onset of the associated X-ray flare. If this event was accelerated by a CME-driven shock, the shock
must have formed very low in the corona and accelerated particles to GeV energies within ~10
minutes or less (Mewaldt et al. 2005a; Saiz et al. 2005), presenting a significant challenge to shock
acceleration models. On the other hand, Simnett (2006) and others have suggested that the January
20 SEPs were accelerated by the X7 flare.
Shock geometry may be a key parameter for understanding how particles can be rapidly
accelerated to very high energies. According to simulations by Giacalone (2005), particles can be
accelerated much more rapidly at quasi- perpendicular shocks than they can at quasi-parallel shocks.
Although it has also been claimed that quasi-perpendicular shocks have a higher injection threshold
than quasi-parallel shocks and therefore favor the acceleration of pre-existing suprathermal ions
over thermal seed particles (Tylka et al. 2005), this is a subject of controversy (Giacalone 2005).
These controversies are best tested close to the Sun, where particles are accelerated to higher
energies than at 1 AU.
We, therefore, have to test whether:
• there is a correlation between shock geometry and the composition and energy spectra of
particles when the shock is still close to the Sun (Giacalone 2005; Mewaldt et al., 2005; Saiz
et al. 2005),
• anomalous diffusion mechanisms can give shorter acceleration times (Zimbardo and Perri,
2013).
3.1.1.5 Do proton-amplified Alfvén waves play a role in accelerating particles at shocks?
At quasi-parallel shocks, the diffusive shock acceleration mechanism operates when
particles repeatedly cross the shock front and gain energy by scattering off magnetic irregularities
carried by the plasma flow. One of the essential elements of this theory is that protons streaming
upstream from the shock can amplify anti-sunward propagating Alfvén waves in the solar wind. The
amplified waves can then resonantly scatter subsequent ions escaping upstream, trapping them near
the shock and increasing the efficiency of the shock acceleration process (Lee 1983). Although the
existence of these waves is a cornerstone of the theory (also invoked for galactic cosmic ray
acceleration by supernovae shocks), the proton beams and the amplified waves that are required for
diffusive shock acceleration of SEPs to high energies (>~10-100MeV) at a CME shock have never
been observed.
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The waves responsible for accelerating MeV ions at shocks are seldom observed at 1 AU,
where they are difficult to observe against the background turbulence (e.g., Bamert et al. 2004), but
in the inner heliosphere, wave growth is expected to proceed more rapidly, and the wave intensities
are predicted to be significantly greater (Ng, et al, 2003) and should be easily detectable at 0.35 AU.
Solar Orbiter observations of the amplified waves by the magnetometer (MAG) and RPW search
coil (for electrons), in association with the proton (and electron) beams and simultaneous
accelerated particle spectra (by EPD), would be a crucial test of this fundamental theory.
3.1.1.6 What causes SEPs’ spectral breaks?
The energy spectra of most gradual SEP events start out as power laws at low energy with a
gradual spectral steepening in the energy range from ~3 to 30 MeV, indicating that the process has
become less efficient at high energies. The “spectral breaks” of heavier species occur at lower
energy/nucleon than for protons, suggesting a rigidity-dependent process (Tylka et al. 2001, Cohen
et al. 2005). It has been suggested that the breaks are organized by diffusion and that they occur at
the locations where there are sudden decreases in the wave intensity. Thus, the proton-amplified
Alfvén waves may also be responsible for determining the Q/A-dependence of spectral breaks in
large SEP events (Cohen et al. 2005, Mewaldt et al. 2005b, Li et al. 2005). Spectral breaks are also
seen in flare accelerated SEPs (Mason et al. 2002), so we require observations of the source as well.
3.1.1.7 Are there favorable environments for particle acceleration?
Statistical studies suggest that CMEs that erupt soon after a previous CME are more efficient
in accelerating particles than those erupting into a pristine environment (Gopalswamy et al. 2004).
Model calculations (Li and Zank, 2005) suggest that particles can be accelerated to energies ~30
times higher at a second shock following a previous shock. This more efficient acceleration may be
due to a stronger turbulence level and a larger population of the seed particles at the second shock,
but there are also other suggested explanations for the observations (see Gopalswamy et al. 2004).
3.1.2 SEPs associated with flares, coronal loops and reconnection regions
As Solar Orbiter approaches the Sun, the photon and particle signatures from small events
will increase by 1/r^2, making it possible to observe events 15-20 times smaller than ever before, in
effect opening a new window for SEP processes. We may detect for the first time energetic particle
populations from X-ray microflares, a candidate mechanism for coronal heating that cannot be
studied further away from the Sun due to background problems. For the small flares that produce X-
ray, electron, and 3He-enrichments we will observe with great accuracy events that at 1 AU are not
far above the level of detection: the timing of particle and radio signatures, the composition and
spectra, etc., providing strong new constraints on the process operating in these events. Particle
acceleration on coronal loops will have new insights since the 1/r^2 sensitivity advantage and
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viewing geometry will make it possible to view the X-ray emission from the tops of loops in
numerous cases where the much stronger footpoint sources are occulted behind the solar limb.
These studies of faint coronal sources that are only rarely observable from 1 AU will give crucial
information about the location and plasma properties of suspected electron acceleration sites in the
high corona.
3.1.2.0 Impulsive SEP event sources
Small impulsive events dominated by ~1-100 keV (sometimes down to ~0.1 keV) solar
electrons are detected hundreds of times a year at ~1 AU near solar maximum (see Lin 1985). Over
the last solar cycle, the occurrence frequency rate for impulsive electron events observed at ~ 1AU
varies with the electron peak flux as a power law with average exponent -1.4, suggesting that many
more events are likely below present detection thresholds. Associated increases in ~0.01 to 1
MeV/nuc ions are often observed (may always be present, but with fluxes too low to be always
detected), with factors of 10-10000 enrichments of 3He (sometimes 3He/4He >1) along with factors
of ~10 enrichment of heavy ions up to Fe, together with high charge states. Curiously, analysis of
the nuclear lines and continuum for the large γ-ray flares that accelerate ions to >~100 MeV shows
that those flares often also have high e/p ratios, enhanced heavy ions, and enhanced He/p ratios
(Ramaty et al. 1993), even though impulsive SEP events at 1 AU are almost never associated with
such large flares.
3.1.2.1 Understand energy release and particle acceleration process.
Large solar flares rapidly accelerate ions up to many GeV and electrons up to 10s of MeV,
and they could be a significant contributor to the SEPs observed in space in gradual events. Imaging
in hard X-rays generated by accelerated electrons, and in the 2.223 MeV neutron-capture γ-ray line
produced by accelerated >~30 MeV ions, show that these emissions come from the footpoints of
newly formed loops (Hurford et al, 2006). Furthermore, the accelerated >~20 keV electrons and >~
few MeV ions often contain >~10-50% of the energy released in the flare (Lin & Hudson 1976; Lin
et al 2003), suggesting that the particle acceleration is intimately related to the magnetic
reconnection process that appears to release the flare energy. Neither the energy release or particle
acceleration processes are understood, but some accelerated particles could end up on magnetic field
lines connected to the heliosphere, as observed in many impulsive events.
Understand the energy release and particle acceleration process:
• Origin from the footprints of newly formed loops (Hurford et al., 2006).
• Acceleration relating to the magnetic reconnection process.
• Properties of flare energy release (in high spatial and temporal resolution).
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3.1.2.2 Evaluate how significantly large flares contribute directly to gradual SEP events
For the magnetically well-connected 2 November 2003 and 20 January 2005 flares, the
spectra inferred for the energetic protons that produce the γ-ray lines observed by RHESSI are
essentially the same, within measurement uncertainties, as the spectra of SEP protons at 1AU. As
already mentioned, the extremely rapid rise of SEP fluxes at 1 AU after the 20 January 2005 flare
X-ray peak raises serious questions about CME-driven shock acceleration. Measurements of
SEPs close to the Sun will provide the timing and compositional information to evaluate how
significantly do large flare contribute directly to gradual SEP events.
3.1.2.3 Flare seed particles
For SEP events in general, the intensity of the photon emission appears poorly correlated
with the particle intensities observed in interplanetary space, so it is not clear whether many of these
particles escape. Particles accelerated by flares of all sizes down to microflares, however, may
provide the seed particles that are preferentially accelerated by interplanetary processes, thus
inseparably linking the two acceleration sites and processes. It has been proposed that SEP events
with large initial Fe/O ratios are a signature of “flare particles” followed by more typical Fe/O ratios
from subsequent interplanetary acceleration. Solar Orbiter observations can test these differing
models since its proximity to the flare site removes most of the uncertainty in magnetic connection,
and the timing differences between the flare acceleration and the interplanetary acceleration phase
would be much clearer than at 1 AU.
3.1.2.4 Explore the fact that only some of the hard X-ray peaks are related to escaping
electrons, while others are not (Benz et al., 2005).
3.1.2.5 X-ray prompt events
In X-ray literature "prompt event" is defined as an event where the estimated release time of
the electrons that are later seen at 1 AU agrees with the hard X-ray flare peak time within the
uncertainty of a few minutes. (From the SEP point of view this is called "prompt impulsive". For
delayed impulsive see below and next section.)
Also, the spectral shapes of the in situ observed electron spectrum and hard X-ray photo
spectrum are statistically correlated. This suggests that a common accelerator produces both the
hard X-ray emitting electrons and the escaping electrons (Droege 1996, Krucker et al. 2007).
However, the observed correlation does not agree with the prediction of the thick-target model
(Brown 1971) that would be expected for footpoint emission. This is currently not understood.
Hard X-ray imaging of the solar source region of prompt events often reveals three sources.
Frequently jets are observed in EUV and soft X-rays that appear to emanate from the hard X-ray
sources. This geometry is similar to that of interchange reconnection models, where closed field
lines reconnect with open fields. The jet indicates the direction of open field lines along which
electrons escape. Prompt events are well correlated with the occurrence of 3He rich SEP events.
However, arrival time studies of ions seen at 1 AU suggest a delayed release of the ions
relative to electrons, at least when assuming scatter-free transport. This is rather puzzling as it
suggests a different accelerator for electrons and ions despite the closely connected occurrence.
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However, electrons and ions could still be released simultaneously with propagation effects
explaining the observed delays.
This objective is about exploring if the comparative delay is real or due to propagation
effects. It should be studied in different heliocentric distances and hopefully near the perihelion in
order to better establish the nature of the flare region (with STIX & EUI).
Complete set of observations is required: STIX, EUI, RPW + ground-based radio
observations (good to have). That will allow studying the magnetic structures along which flare-
accelerated particles escape into interplanetary space.
3.1.2.6 Delayed events (between X-ray peak and electron release time)
These are events that show sufficiently long delay times (>10 minutes) between the peak
hard X-ray emission and the calculated electron release times.
• Is there indeed a delay between the peak hard X-ray emission and the calculated electron
release times or propagation effects are producing the observed delayed onsets?
• The origin of the delayed release times is currently not understood, but two main ideas are
discussed:
• The delay could be due to time-extended electron acceleration and/or electron
storage at high coronal altitudes during solar flares in combination with a delayed
access of these electrons to magnetic field lines open to interplanetary space
(Laitinen et al., 2000; Classen et al., 2003; Klein et al., 2005; Aurass et al., 2006).
• The delay could be the result of electron acceleration by coronal shocks that move
away from the flare site (see also next section 3.2) (Krucker et al., 1999; Haggerty
and Roelof, 2002; Simnett et al., 2002). The delayed release would be produced by
the time that the shock takes to form and to efficiently accelerate electrons (Mann et
al. 1995; Warmuth & Mann 2005), and/or by the time it takes for the shock to reach
magnetic field lines that are connected to the spacecraft. Problem: shocks not very
efficient at accelerating electrons at higher (MeV) energies (but it might work for
shocks that further accelerate a previously produced population of energetic
electrons, e.g. Selkowitz & Blackman 2007).
• In both cases, the related hard X-ray emissions are expected to be seen after the main
hard X-ray bursts. They are likely to be much fainter and to originate from a different
location. However, detection of these faint emissions will be difficult in the presence
of decaying thermal X-ray emissions at these later times. The best chance of
detection will be for occulted flares where the main flare emission is not yet visible.
The faint emission related to the delayed release may be seen on the disk or at high
altitude above the disk.
3.1.2.7 How are so many electrons accelerated on such short time scales to explain the
observed hard X-ray fluxes?
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Hard X-ray imaging observations show most prominent emissions from footpoints of flare
loops in the chromosphere where the ambient density is high enough to stop flare-accelerated
electrons by collisions (e.g. Hoyng et al. 1981). However, fainter, co-temporal hard X-ray sources
are also seen in the corona (e.g. Frost & Dennis 1971, Masuda et al. 1994, Veronig & Brown 2004,
Battaglia & Benz 2006, Krucker et al. 2007) consistent with electron acceleration in the corona. In
particular, RHESSI observations of partially-occulted flares show that at least 90% of all flares have
coronal hard X-ray sources (Krucker & Lin 2008). Further evidence for a coronal acceleration
region comes from radio observations (e.g. Benz 1985, Aurass et al. 2004, Mann et al. 2006). The
details of the transport of electrons from the coronal acceleration site down to the hard X-ray
footpoints are still unclear (e.g. Miller et al. 1997, Önel et al. 2007, Battaglia & Benz 2007).
It is becoming increasingly clear that propagation of energetic electrons does not follow a
simple collisional thick-target scenario, so more sophisticated models of electron transport are
required including effects of non-uniform plasma ionization (e.g. Kontar et al. 2003), return current
(e.g. Zharkova & Gordovskyy 2006), and beam-plasma interaction via various plasma waves (e.g.
Kontar 2001). The deposited energy of the non-thermal electrons heats the chromospheric plasma
and the resulting overpressure drives the hot plasma up the legs of the magnetic loops (e.g. Brown
1973) in the process termed chromospheric evaporation. Hard X-ray observations provide thermal
diagnostics of the heated flare loops.
Observing plans for this objective should include in particular partially limb-occulted flare
observations and observations of hard X-ray emissions associated with CMEs:
Partially limb-occulted flare observations
Partial limb-occultation will frequently provide view angles from which purely coronal
emission (e.g. Hudson 1978) can be readily seen because the footpoint sources are occulted. This
will allow us to study faint coronal sources that otherwise would have been lost in the limited
dynamic range (≤20) of indirect imaging instruments. Coronal hard X-ray emission in the absence
of hard X-ray footpoint emissions, but with the thermal flare loop still partially visible, is best
studied for flares occurring up to ~20° behind the limb. This suggests that for a single spacecraft
about 20% of the observed flares have occulted footpoints with the main flare loop still partially
visible above the limb. In at least 90% of all such events, non-thermal coronal emission is observed,
most prominently during the rise of the thermal emission (Krucker & Lin 2008). Most often the
non-thermal emission is seen close to the thermal loop, but occasionally from above the thermal
flare loops similar to what was reported for the Masuda flare (Masuda et al. 1994). As electron
acceleration is thought to occur in the corona, hard X-ray imaging spectroscopy provides crucial
information on the location and spectrum of energetic electrons before they precipitate to the
chromosphere. STIX will provide frequent partially limb-occulted observations at very high
sensitivity and will be able to see up to 15 weaker coronal emissions than RHESSI. This will allow
us to image hard X-ray emission from the corona at the highest ever sensitivity, free from the
intense footpoint sources, thus providing unique information about the suprathermal electrons
closest to the site in the corona where their acceleration is believed to occur.
Coronal phenomena in hard X-rays associated with CMEs
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The partial occultation of a solar flare by the solar limb is an excellent opportunity for
studying faint coronal HXR emissions without competition from the very bright emission of the
footpoint sources. For flares occurring >25° behind the solar limb (i.e. occultation height larger than
0.1 of a solar radius), not are only the hard X-ray footpoints occulted, but also the main thermal and
non-thermal emissions from the corona. Nevertheless, these highly occulted events also show hard
X-ray emission (e.g. Kane et al. 1992) that is associated with fast backside CMEs (Krucker et al.
2007). The emission is faint but has a rather flat (hard) spectrum indicating that the emission is
produced by non-thermal bremsstrahlung of energetic electrons. Multi-spacecraft observations
reveal that the onset of the emission is simultaneous with the onset of the main hard X-ray emission
seen in footpoints (e.g. Kane et al. 1992), but has a much larger source size that expands with time.
These sources move rapidly (~1,000 km s-1) upwards (Hudson et al. 2001, Krucker et al. 2007) in
the same direction as the associated CME. The high coronal emissions may be produced by flare-
accelerated energetic (>10 keV) electrons trapped in magnetic structures related to the CME, or they
may be accelerated in CME current sheets or other coronal magnetic restructuring related to the
CME. However, the details are not understood. The relative number of non-thermal electrons is
observed to be about 10% of the number of thermal electrons in the high coronal source and the
pressure exerted by the non-thermal electrons may, therefore, be comparable to that of the thermal
plasma itself. High-altitude coronal hard X-ray sources are believed to be a common phenomenon
since RHESSI detected them in association with all fast (>1500 km/s) backside CMEs with flare
locations between ~25° and ~50° behind the limb (Krucker et al. 2007). However, present-day
observations are only sensitive to high coronal emissions related to large X-class flares. With
the 15-times enhanced sensitivity of STIX when near the Sun, high coronal emission from backside
CMEs related to M-class flares will regularly be detected.
3.1.2.8 Explore the type III radio bursts delays
Impulsive SEP events are generally accompanied by solar type III radio bursts, indicative of
electrons escaping from the Sun, that drift down to near the local plasma frequency (tens of kHz).
At 1 AU, for most impulsive events the injection of the >25 keV electrons at the Sun, inferred from
the observed velocity dispersion, is delayed by ~10-30 min after the type III radio bursts (Krucker et
al., 1999; Haggerty & Roelof, 2002). These delays have been suggested to be due to propagation
effects in the interplanetary medium (Cane & Erickson 2003; Cane 2003), or to delayed acceleration
by large-scale coronal transient (EIT or Moreton) waves or by shock waves associated with narrow
CMEs (Simnett et al. 2002; Rouillard et al., 2012), or by coronal magnetic restructuring in the
aftermath of CMEs (Maia & Pick 2004; Klein et al. 2005; van Driel-Gesztelyi et al., 2014). Recent
analyses of highly scatter-free impulsive electron events show evidence for two injections (Wang et
al. 2006): a low-energy (~0.4 to ~12 keV) injection that begins ~10 min earlier than the type III
radio burst, and a high-energy (>~13 to 300 keV) injection that starts ~10 min after. This confirms
that type III bursts are produced by ~1-12 keV electrons, consistent with the type III-producing
Langmuir waves being detected in situ at the time of their arrival at 1 AU. Recent theoretical studies
suggest that some of the delays for higher energy electrons may be due to propagation effects.
In a few events, the injection of energetic ions can be inferred, and they lag behind the
electrons by ~30-60 minutes more. SOHO EIT and LASCO images suggest an association of
impulsive events with jets or narrow fast CMEs (Wang et al 2006); perhaps they
accelerate electrons lower down and ions higher up in the corona. However, the pattern of high
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charge states observed in impulsive events suggests the ions are accelerated in relatively high-
density regions.
3.1.3 Relativistic electron acceleration (including sub-objectives 3.1.3.1 & 3.1.3.2)
Relativistic electrons were observed up to tens of MeV energies in the 1970s and 80s
(Datlowe 1971, Moses et al. 1989), but in the past 30 years, there have been no measurements at
those high energies to compare with the much more sophisticated observations of ions and
electromagnetic emission. In contrast to keV - tens of keV electrons, relativistic electrons generally
exhibit a diffusive flux-time profile, indicating that substantial scattering has occurred in their
propagation to ~1 AU. Electrons are as fundamental as ions to the understanding of the energy
release process. Their spectral shapes provide a nearly perfect diagnostic of the SEP event type
(Moses et al. 1989) - gradual events show single power laws in momentum, while impulsive events
all show a spectral hardening starting at around 5 MeV/c. Although the hardening of the spectrum is
universally observed, the spectral indices vary from event to event. Another surprising result (from
combined Helios/ISEE-3 electron spectra) is that the spectral shapes are apparently invariant as a
function of the azimuthal distance to the flare if the fluxes are adjusted for different radial distances.
Little is known as to the origin of these electrons or the reason for the hardening. The double power
law spectra have been discussed in terms of a superposition of two electron populations, one
accelerated in flaring loops by a stochastic mechanism, and the other by a shock in the high corona
(Dröge 1996a). How shocks can accelerate electrons to relativistic energies (never observed for
shocks near 1 AU) near the Sun is still an unsolved mystery: can coronal shocks accelerate electrons
to relativistic energies starting from a quasi-thermal population, or is a more energetic seed
population necessary?
3.1.4 Other high-sensitivity X-ray studies
3.1.4.1 Hard X-ray emission of escaping electron beams (thin-target emission)
Electron beams escaping into interplanetary space produce faint hard X-ray emission along
their coronal path in the so-called thin-target approximation at a level several orders of magnitude
smaller than the main flare emissions. The only chance of observing this emission is in events that
are well over the limb such that both the hard X-ray footpoints and the thermal X-ray sources are
occulted. Theoretical calculations show that, under extremely favorable conditions, RHESSI
observations could have enough sensitivity to detect thin-target emission from escaping electrons
(Saint-Hilaire et al. 2008). Hard X-ray emission temporally correlated with radio type III bursts was
observed by RHESSI from an elongated hard X-ray source in the corona, possibly outlining the
electron escape path. However, the emission is about an order of magnitude too bright for purely
thin-target emission from the number of escaping electrons seen near 1 AU by WIND/3DP. STIX
will also not be able to regularly observe thin-target emission from escaping electrons, but, under
favorable conditions such as a hard electron spectrum and high ambient plasma density, STIX will
provide the first clean detections of purely thin-target emission from escaping electron beams.
3.1.4.2 X-ray emission from electrons accelerated at CME shocks
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Solar Orbiter will also make it possible to search for hard X-ray emission from electron
beams accelerated at the CME shocks that produce type II radio bursts. As the density of the
emission region of a type II burst is low (<10^9 cm-3), the related hard X-ray emission is expected
to be very faint and only upper limits have been derived so far (Klein et al. 2003). In any case, the
detection will only be possible in highly occulted flares. The high sensitivity of STIX down to low
energies (~4 keV) together with the large number of partially occulted flares with type II bursts will
greatly help to find an event or at least provide more stringent upper limits. Such observations
would identify where along the shock electrons are accelerated, and would provide the first
quantitative measurements of current unknowns, such as the spectrum and energy content of the
electrons associated with radio type II bursts (e.g. Mann & Klassen 2005). Because there have never
been hard X-ray observations of these shock fronts, these observations could provide decisive tests
for theoretical models (e.g. Cairns et al. 2003).
3.2.3.2 3.2 How are energetic particles released from their sources and distributed in space
and time?
Present state of knowledge:
SEPs associated with CME-driven shocks have been long known to often arrive at Earth
orbit hours later than would be expected based on their velocities. There are two alternate processes
that might cause this. (1) The acceleration may require significant time to energize the particles
since they must repeatedly collide with the shock to gain energy in many small steps, so the process
may continue for many hours as the shock moves well into the inner solar system. Or (2) the particle
intensities near the shock may create strong turbulence that traps the particles in the vicinity of the
shock, and their intensity observed at Earth orbit depends on the physics of the particles escaping
from the trapping region. Once free of the vicinity of the shock, SEPs may spiral relatively freely on
their way to earth orbit, or more usually will be scattered repeatedly from kinks in the IMF, delaying
their arrival further. The amount of scattering in the interplanetary space varies depending on other
activity such as recent passage of other shocks or solar wind stream interactions. By the time the
particles reach Earth orbit, they are so thoroughly mixed that these effects cannot be
untangled (Gopalswamy et al. 2006; Cohen et al. 2007).
Particles accelerated on magnetic loops can reach very high energies in seconds after the
onset of flaring activity, and then collide with the solar surface where they emit gamma radiation.
There is a poor correlation between the intensity of the gamma radiation and the SEP intensities
observed at Earth orbit, so most particles from this powerful acceleration process do not escape.
Much more common are flare events observed in UV and X-rays that produce a sudden acceleration
of electrons. The electrons can escape from the corona, producing nonthermal radio emission as
they interact with the local plasma. Moving from higher to lower frequencies as the local plasma
density decreases with altitude, the (type III) radio emission makes it possible to track the energetic
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electron burst into interplanetary space where it may pass by the observer. Energetic ions, greatly
enriched in 3He and heavy nuclei, accompany these electron bursts (Lin 2006; Mason 2007).
Key open questions in shock associated events are whether particle arrival delays at 1 AU
are due to the length of time needed to accelerate the particles, or due to trapping in the turbulence
near an accelerating shock or a combination of both? For particles accelerated on loops, are the
electrons and ions accelerated from sites low in the corona or at higher altitudes, and how are they
related to the X- and gamma-ray signatures?
How Solar Orbiter will address this question:
Solar Orbiter will revolutionize our understanding of SEP acceleration associated with
CME-driven shocks by probing the inner heliospheric sites where particle acceleration and release
take place. Solar Orbiter will observe how shocks evolve, and whether they are still accelerating
particles as they pass by the spacecraft. If particle arrivals are controlled by the time it takes the
shock to accelerate them, then the highest energy particles will be delayed since they require many
more interactions with the shock. If trapping and release control the timing, then as the shock moves
by the faster and slower particles will have similar intensity changes. Since Solar Orbiter will
simultaneously measure the turbulence properties in the shock acceleration region, it will be
possible to construct a complete theory and models of the acceleration process, and its radial
dependence in the inner heliosphere.
For SEPs accelerated on loops or in reconnection regions, Solar Orbiter will see the coronal
location from UV and X-rays, and then trace the progress of released electrons by radio emission
that will drift to the plasma frequency at the spacecraft for those bursts that pass by. This
unambiguously establishes that the magnetic field line at Solar Orbiter connects to the coronal UV
and X-ray emission site. Since Solar Orbiter can be connected to active regions for periods of days,
this will provide multiple tracings between the heliospheric magnetic field and its origin in the
corona. The corotation phase of Solar Orbiter will considerably lengthen the periods of connection
to active regions, greatly increasing the number of field line origin sites that can be determined from
a single active region. X-ray emission from the flaring sites can be used to derive the energetic
electron spectrum at the flare site, which in turn can be compared with the escaping population to
see if most of the accelerated electrons are released (usually most do not escape). Thanks to the
1/r^2 intensity advantage, Solar Orbiter will observe thousands of these cases and thereby permit
detailed mapping of coronal sources and the trapping properties of the acceleration sites.
3.2.0 What controls the escape of the particles to the heliosphere?
For gradual SEP events, the self-generated turbulence around CME-driven shocks may be
important. Theoretical modeling has pointed out a few possible mechanisms: the particles escape
before large wave growth occurs, adiabatic focusing in the diverging magnetic field, and/or the
shock arriving at a region less favorable to wave growth. On the other hand, if strong turbulence
exists in the corona (i.e., not self-generated), particle acceleration and escape are likely to be
consecutive processes. For impulsive SEP events, self-generated turbulence is likely to be weak, and
escape will depend on the SEPs reaching magnetic field lines connected to the heliosphere. Timing,
compositional, spectral and anisotropy observations for SEPs, together with in situ wave and
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magnetic field measurements close to the Sun and field line tracking are key to discriminating
between escape mechanisms.
3.2.1 How do energetic particles scatter and move along the interplanetary magnetic field?
(including sub-objectives 3.2.1.1, 3.2.1.2 & 3.2.1.3)
After SEPs leave their acceleration sites - flares or CME-driven shocks or high coronal
structures - they propagate along the open coronal magnetic fields that expand rapidly close to the
Sun, and then along the large spatial scale interplanetary magnetic field (IMF) - Parker's spiral for
steady flows but highly complex in transients such as ICMEs. The IMF appears to be random and
turbulent at smaller spatial scale. Understanding how energetic particles propagate from their
sources to the point of observation is not only essential to better understand the acceleration
processes on or near the sun, but an important problem in its own right. Propagation of charged
particles parallel to the IMF is affected by two competing processes: the adiabatic motion due to the
divergence of the large-scale IMF results in “focusing” wherein particle pitch angles decrease on
average as they move out; and scattering in pitch angle by small-scale magnetic irregularities at a
rate depending on the strength of the turbulent magnetic field.
A crucial parameter to solve the transport equation is the Fokker-Planck coefficient, Duu,
which can be related to the particle's mean free path if the power spectrum of the turbulence is
known. To understand SEP observations at 1 AU, model calculations of the transport equations are
fit to observation, with assumptions on the r-dependence of the turbulent magnetic field along the
Parker spiral from the acceleration site to 1 AU, since we can only obtain wave spectra at 1
AU. Solar Orbiter’s in situ magnetic field and plasma measurements will map the power spectrum
of the turbulent magnetic field as a function of heliocentric distance in to 0.28 AU.
Transport of ions and electrons may be quite different because electrons and ions resonate in
different regions of the turbulence power spectrum, leading to a rigidity dependence of the particle's
mean free path. Whereas high-rigidity particles are suited to probe the geometry of the fluctuations,
low-rigidity particles are sensitive to the dissipation range and dynamical and thermal effects. Some
models (Dröge, 2003), using power spectra at 1 AU, are able to correctly account for the
dependence of the scattering mean free path on the particle's rigidity and have significantly
improved the agreement between scattering theory and observations. It is important to examine if a
similar rigidity dependence also holds at 0.28 AU.
3.2.2 Latitudinal and longitudinal transport of SEPs
The unique perspective of Solar Orbiter from its inclined orbit in the high latitude phase
should aid in the study of the latitude distribution of gradual and impulsive SEP events close to the
Sun. Ulysses measurements in the outer heliosphere (>1 AU) show that the difference in latitude
between Ulysses and the associated flare orders both onset times and times to maximum well (Dalla
et al. 2003). Sanderson et al. (2003) analyzed the same set of events and found that SEPs either
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propagate at the highest heliographic latitudes along magnetic field lines and not across them at
Ulysses’ position. Since onset delay is not correlated with CME parameters, these authors and
Struminsky et al. (2006), conclude that the SEPs either propagate along the distorted magnetic field
with large latitudinal excursions (e.g. Fisk 1996) or diffuse across magnetic field lines close to the
Sun to reach high latitudes. Wibberenz and Cane (2006) suggested lateral coronal transport via low
coronal magnetic loops to explain multi-spacecraft observations including Helios. Thus, high
latitude measurements of SEPs, plus field line tracking, will be crucial to understanding lateral
transport.
We, therefore, need to study the following:
• Propagation at high heliographic latitudes along magnetic field lines vs across them
(Sanderson et al., 2003).
• Propagation along the distorted magnetic field with large latitudinal excursions
(Struminsky et al., 2006).
• Diffusion across magnetic field lines close to the Sun to reach high latitudes
(Dressing et al., 2014).
• Lateral coronal transport via low coronal magnetic loops (Wibberenz and Cane,
2006; Klein et al., 2008).
• Wide angular spread of SEP events. Relative role of cross-field diffusion and shocks
(Dresing et al., 2014; Lario et al., 2014; Gomez-Herrero et al., 2015).
3.2.3 Properties and distribution of near-Sun shocks, their fluctuations and particle
acceleration
The answers to fundamental questions about the sources and acceleration mechanisms of the so-
called large gradual SEPs, whether flares or shocks, depend on many factors including magnetic
connection to active regions that produce flares (e.g., Cane et al., 2006) and on properties of the
CME-driven shocks in the solar corona and the inner heliosphere (Tylka et al. 2005; Desai et al.,
2006a). We need to understand the coronal conditions under which the shocks form and determine
the interplanetary conditions where they evolve. We must then use this information to characterize
shock structure, determine whether they are locally quasi-perpendicular or quasi-parallel, and then
study kinetic properties of shock-related waves and turbulence that control ion scattering mean free
paths near the shock. With geometric and microphysical shock properties determined, we can infer
the rates of particle acceleration and the injection energies as a function of distance and along
different parts of the shock. We need to derive particle acceleration rates near shocks, and by
measuring shocks at many different radii, we will determine the acceleration rates as a function of
distance from the Sun.
• Detect shocks and determine their fine-scale structure (in situ trigger burst mode) at a range
of Mach numbers (Cane and Lario, 2006).
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• Measure the waves and turbulence around the shocks and compare these with particle
observations and models in order to determine the acceleration processes of the particles
(Sandroos and Vainio, 2006).
3.2.4 How do large and small-scale structures modulate particle fluxes?
The sharp decreases of high energy cosmic ray fluxes at the Earth, known as Forbush
decreases (e.g. Ifedili, 2004), are caused the passage of “magnetic barriers” in the solar wind. As
well as magnetic clouds, these barriers are often the sheaths of compressed solar wind ahead of
ICMEs, or compressed CIRs (Clack et al., 2000). Such compressions lead to “planar magnetic
structures” (e.g. Jones et al., 2002: see Figure below) and it appears to be these sheets of magnetic
field that efficiently block the particles (Intriligator et al., 2001). However, the efficiency of these
barriers, as they develop close to the Sun, is not known.
• Study of the planarity and large-scale structures of the magnetic field within and around
CIRs and ICMEs and investigate their effect on the fluxes of solar energetic particles and
cosmic rays as these structures evolve with heliocentric distance.
At a much smaller scale, “particle channels” of dramatically enhanced or reduced particle flux
during solar particle events, lasting only around an hour, as well as comparable duration burst of
Jovian electrons observed several AU from the planet, demonstrate the existence of very small scale
magnetic connections and disconnections from particle sources. Solar Orbiter will travel much
closer to the solar sources of particles and, by nearly co-rotating with the Sun, distinguish temporal
variability from spatial structures.
• Determine the small-scale diffusion of particles around and within particle channels, as well
as the spatial and temporal scales of these structures and their connectivity to the Sun.
3.2.5 Shock-surfing acceleration mechanism
See Objective 2.3.4.
3.2.6 Effects of energetic particles propagating downward in the chromosphere
• Generation of heating, shocks, sunquakes in the underlying photosphere (Martinez-Oliveros
et al., 2007).
This objective does not appear in any SOOP. Should we include it
in R_SMALL_HRES_HCAD_WaveStereoscopy
and/or R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure?
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3.2.3.3 3.3 What are the seed populations for energetic particles?
Present state of knowledge:
The low-energy particles accelerated by CME-driven shocks to SEP energies are called the
seed population. The observed ionization states of SEP ions show temperatures typical of the
corona, ruling out hot material on flare loops as the seeds. But SEPs also show significant
abundances of ions such as 3He and singly ionized He, which are virtually absent from the solar
wind. The observed energetic particle abundances indicate that the suprathermal ion pool, composed
of ions from a few to a few 10 s the speed of the solar wind, is the likely source. At 1 AU, the
suprathermal ion pool is ~100 times more variable in intensity than the solar wind and varies in
composition depending on solar and interplanetary activity. The suprathermal ions are continuously
present at 1 AU, but it is not known if there is a continuous solar source, or if these ions are from
other activities such as acceleration in association with fast and slow solar wind streams. Inside 1
AU, the suprathermal ion pool is expected to show significant radial dependence due to the different
processes that contribute to the mixture, but it is unexplored (Desai et al. 2006; Mewaldt et al. 2007;
Lee 2007; Fisk and Gloeckler 2007).
For SEPs accelerated on loops or in reconnection regions that give rise to electron and type-
III radio bursts, ionization states are coronal-like at lower energies and change over to much hotter
flare-like at high energies. This may be evidence for a complex source, or, more likely, of energetic
particle stripping as the ions escape from a low coronal source. For SEPs accelerated at
reconnection sites behind CMEs abundances and ionization states would be coronal (Klecker et al.
2006).
Critical questions in this area are: what is the suprathermal ion pool in the inner heliosphere,
including its composition and temporal and spatial variations? What turbulence or stochastic
mechanisms in the inner heliosphere accelerate particles to suprathermal energies? Are the source
locations and arrival times of electrons from SEPs on loops or reconnection regions consistent with
a low or high coronal source?
How Solar Orbiter will address this question:
By systematically mapping the suprathermal ion pool in the inner heliosphere with
spectroscopic and in-situ data, Solar Orbiter will provide the missing seed particle data for models
of SEP acceleration associated with shocks. Together with the shock and turbulence parameters also
measured on Solar Orbiter, there will be the first well-constrained models. Since the suprathermal
ion pool composition varies, different shock events will be expected to produce correspondingly
different energetic particle populations that can be examined on a case-by-case basis. The high-
latitude phase of the mission will add an important third dimension to the suprathermal pool
mapping, since it will be more heavily influenced by, e.g., mid-latitude streamer belts, making it
possible to probe the solar and interplanetary origins of the seed particle populations. Taken
together, these observations will make it possible to construct the first complete physics-based
theory and models of particle acceleration close to the Sun.
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For SEPs accelerated on loops or in reconnection regions the 1/r^2 advantage of Solar
Orbiter will again provide a decisive advantage since particle properties will be accurately measured
and compared with much more precise information on the coronal location. This will permit
distinguishing between low coronal sources that result in stripping of escaping particles vs. higher
sources which could mimic stripping properties. SEPs accelerated from reconnection regions in the
back of CME lift-offs will be identified by comparing energetic particle timing with the location of
the CME, and energetic particle composition with that determined spectroscopically for the remote
coronal source.
3.3.1 What are the properties and distribution of suprathermal seed populations?
The suprathermal population contains a mix of solar wind ions, flare populations, gradual
SEP events, pickup ions, and other particle populations (see figure). Since the populations in this
mix have distinct histories, suprathermal elemental and charge-state composition are likely to vary
with energy (Fisk and Gloeckler 2007), time (Desai et al., 2006a, b) and location. The high initial
speeds of suprathermal ions predispose them to efficient injection into acceleration near shocks,
making them the ideal seed population for SEPs. Where shocks are quasi-perpendicular, we expect
injection energies to be high, and the SEP population should resemble the composition of the higher
energy suprathermal ions. Where shocks are quasi-parallel, we expect the SEP composition to
reflect the lower energy portions of the suprathermal population.
3.3.1.1 Characterization of the suprathermal population
▪ Solar wind ions
▪ Flare populations
▪ Gradual SEP events
▪ Pickup ions
▪ Other particle populations
3.3.1.2 Characterize the suprathermal elemental and charge-state composition as a function of
energy, time, and location.
Characterize the suprathermal elemental and charge-state composition as a function of energy (Fisk
and Gloeckler, 2007), time (Desai et al., 2006a, b), and location.
3.3.1.3 Role of shocks in generating SEPs
• Role of shocks in generating SEPs: comparison between the composition of suprathermal
ions near the shocks and that of energetic particles.
• Distinction for quasi-perpendicular and quasi-parallel shocks.
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• Spectroscopic observations of CME shocks’ line profiles.
• Visible light and UV observations of shocks in the inner corona.
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3.2.4 Objective 4: How does the solar dynamo work and drive connections between the
Sun and the heliosphere?
The Sun’s magnetic field dominates the solar atmosphere. It structures the coronal plasma,
drives much of the coronal dynamics, and produces all the observed energetic phenomena. One of
the most striking features of solar magnetism is its ~11-year activity cycle, which is manifest in all
the associated solar and heliospheric phenomena. Similar activity cycles are also observed in a
broad range of stars in the right half of the Hertzsprung-Russell diagram, and the Sun is an
important test case for dynamo models of stellar activity.
The Sun’s global magnetic field is generated by a dynamo generally believed to be seated in
the tachocline, the shear layer at the base of the convection zone. According to flux-transport
dynamo models (e.g., Dikpati and Gilman 2008), meridional circulation, and other near-surface
flows transport magnetic flux from decaying active regions to the poles. There, subduction carries it
to the tachocline to be reprocessed for the next cycle. This ‘conveyor belt’ scenario provides a
natural explanation for the sunspot cycle and characterizing the flows that drive it will provide a
crucial test of our models and may also allow us to predict the length and amplitude of future cycles.
However, current models fail miserably at predicting actual global solar behavior. For example, the
current sunspot minimum has been far deeper and longer than predicted by any solar modeling
group, indicating that crucial elements are missing from current understanding.
A major weakness of current global dynamo models is the poor constraint of the meridional
circulation at high latitudes. The exact profile and nature of the turnover from poleward flow to
subduction strongly affect behavior of the resulting global dynamo (e.g., Dikpati and Charbonneau
1999), but detecting and characterizing the solar flow is essentially impossible at shallow viewing
angles in the ecliptic plane.
In addition to the global dynamo, turbulent convection may drive a local dynamo that could
be responsible for generating the observed weak, small-scale internetwork field, which is ubiquitous
across the surface and appears to dominate the emergent unsigned flux there.
A key objective of the Solar Orbiter mission is to measure and characterize the flows that
transport the solar magnetic fields: complex near-surface flows, the meridional flow, and the
differential rotation at all latitudes and radii. Of particular and perhaps paramount importance for
advancing our understanding of the solar dynamo and the polarity reversal of the global magnetic
field is a detailed knowledge of magnetic flux transport near the poles. Hinode, peering over the
Sun’s limb from a heliographic latitude of 7°, has provided a tantalizing glimpse of the Sun’s high-
latitude region above 70°; however, observations from near the ecliptic lack the detail, coverage,
and unambiguous interpretation needed to understand the properties and dynamics of the polar
region. Thus, Solar Orbiter’s imaging of the properties and dynamics of the polar region during the
out-of-the- ecliptic phase of the mission (reaching heliographic latitudes of 25° during the nominal
mission and as high as 34° during the extended mission) will provide urgently needed constraints on
our models of the solar dynamo.
Most of the open magnetic flux that extends into the heliosphere originates from the Sun’s
polar regions, from polar coronal holes. The current solar minimum activity period, which is deeper
and more extended than previously measured minima, demonstrates the importance of this polar
field to the solar wind and heliosphere.
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There is evidence that the solar wind dynamic pressure, composition and turbulence levels,
as well as the strength of the heliospheric magnetic field, have all changed in the last few years in
ways that are unprecedented in the space age. None of these changes were predicted, and current
solar conditions present a challenge to our understanding of the solar dynamo and its effects on the
solar system at large and the Earth in particular.
In the following sections, we discuss in more detail three interrelated questions that flow
down from this top-level question: How is magnetic flux transported to and reprocessed at high
solar latitudes? What are the properties of the magnetic field at high solar latitudes? Are there
separate dynamo processes in the Sun?
3.2.4.1 4.0 Overall remarks and feasibility concerning Objective 4 observations with Solar
Orbiter
The most important assets of Solar Orbiter/PHI for helioseismology are:
• Combine Earth-based, front-side observations with PHI observations from back-side (and if
possible higher latitudes) to
o be able to observe front and back side sun simultaneously for few days, which would
strongly improve local helioseismology
o improve and calibrate the far side modeling based on helioseismology
o Use PHI’s out-of-ecliptic observations to measure meridional flow at high latitudes
For local helioseismology, PHI would need to observe at 1 min cadence and good resolution. Only
Dopplergrams are needed, i.e. 1 data product out of the 5 in PHI’s standard dataset.
Helioseismology observing strategies & remarks
• The requirements for local helioseismology and feature tracking are a moderate spatial
resolution (f-mode wavelength; a few Mm) with a temporal cadence of 60s.
• For granulation feature tracking, one needs to resolve granules. The cadence should
be 60s between a pair of observations, then a waiting time of 30 min can be introduced
before recording the second pair of two 60s separated images. Onboard processing is
possible to obtain the flow map.
• For supergranulation tracking, images can be taken separated by 1h.
Again, onboard processing is possible. Granulation needs to be suppressed, which requires
averaging on board. Furthermore, the solar rotation needs to be taken into account. High-
resolution images are best taken close to the disk center to minimize limb effects.
• Noise reduction requires long time series. Given the noise level that increases with latitude
and the smaller number of features at the poles, the error on the horizontal flow speed goes
up to >6m/s at the poles with a time series of 30 days length, i.e. the meridional flow might
not be measurable at the poles.
• Compression can help reduce the data rates with little effect on time-distance
helioseismology (JPEG) and local correlation tracking (quantization).
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• Studying the center-to-limb effect in local helioseismology could help understanding the
physics of solar oscillation modes.
• Different observing strategies are suggested in order to address the following science goals:
o Near-surface differential rotation: for measuring differential rotation at high latitudes
one set of observations would be helpful. This requires 30 days of data. Resolving
the torsional oscillations requires repetition of measurements. These 30 days of data
do not need to be recorded at once. Averaging of single power spectra obtained from
observing several sequences of 10 days is possible. Both HRT and FDT could be
used here: HRT is only necessary for granulation tracking but the two other
techniques could be based on FDT data.
o Near-surface meridional flow at high latitudes requires 120 days of data due to the
low amplitude of the flow. Again, averaging of data is possible, but this smears out
cycle effects. Observations for this sub-objective make sense in the later phases of
the mission as a minimum latitude of 15-20 degrees is required to start analyzing the
poles (depending on B angle).
o For studying convection at high latitudes, 7 days of data are needed for obtaining
flow maps and useful power spectra. Weeks to months of data would be needed for
statistical analyses.
o Large-scale convective flows
• A possible orbit for all these measurements is MTP11 - 2024/01/01 - 2024/07/01, with Solar
Orbiter at 45 degrees from Earth.
• Furthermore, observations need to be done with FDT or HRT with 60 days of observation
time: FDT at high latitude far from the Sun (0.7 AU) and HRT at high latitudes close to the
Sun (0.4 AU).
• In summary, observing times are the biggest constraint. As a high cadence of 60s is required
for many science goals, the feasibility of the observing programs highly depend on the
compression that can be applied. For each of the science objectives above, the best methods
need to be defined as well as the most suitable compression.
• Some of the helioseismology goals can be addressed with synoptic observations. PHI
synoptic observations outside the RS windows could be helpful but this is not included in the
current baseline of the mission operations.
Measurements and topology of the polar magnetic field
• PHI is optimally used for measuring the solar polar magnetic field at maximum solar latitude
and at minimum distance.
• Co-observations should be done from Earth, i.e. at a large B0 angle in March and September
when Solar Orbiter is observing the pole visible to Earth. September is preferable as
observations are then possible from the solar telescopes on the Canary Islands. In the current
planning (October 2018 Option E), a suitable orbit for this is MTP14 - 2025/07/01 -
2026/01/01.
• These observations should be combined with other observables, especially EUI observing
polar jets and SPICE.
• Polar magnetic field useful publications include Tsuneta et al., 2008, ApJ 688, 1374 and
Shiota et al., 2012, ApJ, 753, 157.
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Magnetoconvection
• Magnetoconvection is present on the Quiet Sun and in active regions.
• The challenges are the various spatial scales that range from 100 km to 50 Mm, as well as
the various temporal scales ranging from 1 min (small magnetic elements) to 1 week (active
regions).
• Observations from the ground are affected by solar rotation, i.e. the maximum observing
time is 14 days. In addition, projection effects influence the data.
• Zeeman measurements exhibit 180° ambiguity in the azimuthal magnetic field component,
which can be resolved by stereoscopic observations.
• For this objective, it is important to use the good spatial resolution (pixel size of 110 km)
and the large field of view 1000’’x1000’’ of PHI/HRT for combined observations between
Earth-bases telescopes and Solar Orbiter near the almost co-rotating phases to allow long
periods of AR tracking. This should enable studying the long-term behavior of active
regions. The advantage of the almost co-rotation will allow observing many flares and
follow the decay of active regions.
• Solar Orbiter's most important assets for magnetoconvection are:
o longer observations of the same target due to near co-rotation phase close to the Sun,
o Solar Orbiter is an observatory combining many remote sensing instruments,
o combined with ground-based support or SDO, it will provide a new vantage point
and allow stereoscopy.
• It is suggested to use an observing program of 15 days during maximum and declining solar
cycle 25 to follow up AR dynamics.
• An optimal orbit is MTP-07 (0.3 AU, 8-11.5 degree/day, latitude change of 15 degrees).
MTP-10 would be fine too.
• Another observing sequence should include the full remote-sensing package (PHI, EUI,
SPICE) coordinated with DKIST and GREGOR. Observing targets should be quiet sun
regions at the disc center, the limbs and polar regions. Possible orbits are the first orbit and
later orbits of the mission with higher latitudes. The length of the time series should be 3h
per pointing (telemetry intensive observations).
Other related goals
• Measure solar oblateness and luminosity of the Sun. For this, rolls of the spacecraft are
required, with a minimum of 8 positions, preferably close to the equatorial plane. This could
be done during communication or instrument calibration spacecraft rolls.
• Concerning the question whether a small-scale dynamo is acting on the Sun, the difficulty
for observations is that the very convolved field is hard to resolve spatially. Solar Orbiter
could look at the distribution of the magnetic field and the emergence rate per feature over
latitudes. If emergence rate is independent of latitude, probably the feature is produced by
local dynamo rather than local one.
• Regarding the solar irradiance, Solar Orbiter has no measurements of the Sun's luminosity.
However, it can help obtaining synoptic charts with extended viewing windows. Those could
be used for modes of the total solar irradiance.
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• Furthermore, Solar Orbiter can help answer the question whether the observations of the
Sun's chromospheric and photospheric activity is different from that observed from other
stars. As one sees the Sun only from the ecliptic, comparisons with other stars might be
biased, since they are observed from all viewing angles. By leaving the ecliptic, we might be
able to test this hypothesis.
Related SOOPs
• The following SOOPs are defined for helioseismology observations:
o R_SMALL_HRES_HCAD_ModePhysics: for most helioseismology objectives.
o R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure: for addressing flux
emergence in the quiet sun, as well as polar features.
o R_SMALL_MRES_MCAD_AR_LongTerm: for the decay of Active Regions
o R_SMALL_HRES_HCAD_PDF_Mosaic: for the measurement of the probability
distribution function.
3.2.4.2 4.1 How is magnetic flux transported to and re-processed at high solar latitudes?
Present state of knowledge:
In the last decades, the mapping of surface and subsurface flow fields at low and middle
latitudes has seen major advances, largely due to the availability of high-quality data from the
SOHO's Michelson Doppler Imager (MDI) instrument. These data have provided accurate
knowledge of differential rotation, the low latitude, near-surface part of the meridional flows, and
the near-surface torsional oscillations, which are rhythmic changes in the rotation speed that travel
from mid-latitudes both equatorward and poleward (Howe et al. 2006). Local helioseismic
techniques have also reached a level of maturity that allows the three-dimensional structure of the
shallow velocity field beneath the solar surface to be determined.
Despite these advances, progress in understanding the operation of the solar dynamo
depends on how well we understand differential rotation and the meridional flows near the poles of
the Sun. However, because of the lack of out-of-the-ecliptic observations, the near-polar flow fields
remain poorly mapped, as does the differential rotation at high latitudes (see Beck 2000; Thompson
et al. 2003). The meridional flow, in particular, the very foundation of the flux transport dynamo, is
not well characterized above ~50° latitude; it is not even certain that it consists of only one cell in
each hemisphere. The return flow, believed to occur at the base of the convection zone, is entirely
undetermined save for the requirement of mass conservation. All these flows must be better
constrained observationally in order to help solve the puzzle of the solar cycle and to advance our
understanding of the operation of the solar dynamo (and, more broadly, of stellar dynamos
generally).
How Solar Orbiter will address this question:
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Solar Orbiter will measure or infer local and convective flows, rotation, and meridional
circulation in the photosphere and in the subsurface convection zone at all heliographic latitudes
including, during the later stages of the nominal mission, at the critical near-polar latitudes. Solar
Orbiter will reveal the patterns of differential rotation, the geometry of the meridional flow, the
structure of subduction areas around the poles where the solar plasma dives back into the Sun, and
the properties of convection cells below the solar surface. This will be achieved through correlation
tracking of small features, direct imaging of Doppler shifts, and helioseismic observations
(including the first from a high-latitude vantage point). By monitoring the temporal variations over
the course of the mission, it will be possible to deduce solar cycle variations in the flows.
Solar Orbiter will resolve small-scale magnetic features near the poles, even within the
nominal mission phase, and right up to the poles during the extended mission. It will determine the
detailed surface flow field through tracking algorithms. Such algorithms provide only inconclusive
results when applied to polar data obtained from near-Earth orbit due to the foreshortening. Doppler
maps of the line-of-sight velocity component will complement the correlation tracking
measurements and will also reveal convection, rotation, and meridional circulation flows.
Time series of Doppler and intensity maps will be used to probe the three-dimensional mass
flows in the upper layers of the convection zone, at high heliographic latitudes. The flows will be
inferred using the methods of local helioseismology (e.g., Gizon and Birch 2005): time-distance
helioseismology, ring diagram analysis, helioseismic holography, and direct modeling. Using
SOHO/MDI Dopplergrams, it was demonstrated that even complex velocity fields can be derived
with a single day of data (e.g., Jackiewicz et al. 2008).
The deeper layers of the convection zone will be studied using both local and the global
methods of helioseismology. Moreover, Solar Orbiter will provide the first opportunity to
implement the novel technique of stereoscopic helioseismology to probe flows and structural
heterogeneities deep in the convection zone, even reaching down to the tachocline. Combining Solar
Orbiter observations with ground- or space-based helioseismic observations from 1 AU (e.g.,
GONG or SDO) will open new windows into the Sun. Looking at the Sun from two distinct viewing
angles will increase the observed fraction of the Sun’s surface and will benefit global
helioseismology because the modes of oscillation will be easier to disentangle (reduction of spatial
leaks). With stereoscopic helioseismology, new acoustic ray paths can be taken into account to
probe deeper layers in the interior, including the bottom of the convection zone.
4.1.1 Study the detailed solar surface flow patterns in the polar regions, including coronal hole
boundaries.
Solar Orbiter will provide the first detailed view of the polar subsurface layers by carrying out
seismic measurements both from a high-latitude vantage point and "stereoscopically" by combining
PHI data and seismic data from the ground or from NEO (e.g. SDO/HMI). These observations will
reveal the patterns of differential rotation, the geometry of the meridional flow, and the properties of
convection cell below the solar surface. By monitoring the temporal variations over the course of
the mission, it will be possible to deduce solar-cycle variations in the flows.
Thanks to global helioseismology (e.g. Christensen- Dalsgaard 2002), the solar differential rotation
has been mapped as a function of latitude and radial distance throughout most of the convection
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zone. For heliographic latitudes above 70°, however, the global oscillation mode inversions are
uncertain, leaving an incomplete picture of the solar interior. Local helioseismology (e.g. Gizon &
Birch 2005) aims to measure the 3D velocity vectors of the material flows in the solar interior,
allowing studies of convective, rotational and meridional flows, as well as sunspots and active
regions. Local helioseismology together with PHI observations will enable the study of high
heliographic latitudes. Another important goal of Solar Orbiter is to implement stereoscopic
helioseismology by combining PHI data with Doppler measurements from Earth or NEO
instruments. Local helioseismic inversions from techniques such as time-distance helioseismology
or helioseismic holography will be able to probe deeper into the Sun using observations from widely
separated vantage points because skip distances of more than half a circumference will, at last,
become accessible. This will be important for probing the tachocline at the base of the convection
zone, where the dynamo is surmised to be situated.
• Track granules and magnetic features to follow their motions and mutual interaction over
time (Abramenko et al., 2011; Giannattasio et al., 2013; 2014; Gosic et al., 2014; Requerey
et al., 2014).
• Track mesogranules and supergranules and determine lifetime, sizes, horizontal velocities,
and other properties such as helicity of the flows (e.g. Hathaway 2000, Palacios 2012).
4.1.2 Study the subtle cancellation effects that lead to the reversal of the dominant polarity at
the poles
Study the subtle cancellation effects that lead to the reversal of the dominant polarity at the poles
(Wang et al., 1989; Sheeley, 1991; Makarov et al., 2003).
4.1.2.1 Studies of convective, rotational and meridional flows.
4.1.2.2 Local helioseismology of polar regions (Gizon and Birch, 2005).
4.1.2.3 MDI-like medium-l program (see Löptien et al., 2014).
4.1.2.4 By which physical processes is magnetic flux at the poles removed?
4.1.2.5 How does the cancellation process at the poles differ from cancellation at lower latitudes?
Are there different rates?
4.1.3 Explore the transport processes of magnetic flux from the activity belts towards the poles
and the interaction of this flux with the already present polar magnetic field
4.1.3.1 Follow the evolution of the magnetic flux at different latitudes over the solar cycle (full
disk).
4.1.3.2 Follow individual magnetic features flux from lower to high latitudes (high spatial
resolution).
4.1.3.3 Determine flow rates (differential rotation, meridional flow) obtained by tracking magnetic
features and compare with those determined from tracking granules and from Doppler shifts.
4.1.3.4 How do supergranular flows facilitate or impede the latitudinal transport of small-scale
magnetic features?
4.1.3.5 Compare measurements of the magnetic flux from high and low heliographic latitudes (full
disk and high resolution). How strongly do (near) Earth-based magnetic field measurements need to
be corrected at higher latitudes?
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4.1.4 Study the influence of cancellations at all heights in the atmosphere
4.1.4.1 Understand the origin of polar jets.
3.2.4.3 4.2 What are the properties of the magnetic field at high solar latitudes?
Present state of knowledge:
Meridional circulation transports the surface magnetic flux toward the poles, where a
concentration of magnetic flux is expected to occur. However, because of the directional
sensitivity of the Zeeman effect and magnetic polarity cancellation resulting from geometric
foreshortening, present-day observations from the ecliptic at 1 AU can provide only a poor
representation of the polar magnetic field. The high resolution of Hinode’s Solar Optical Telescope
(SOT) can partly overcome the second disadvantage (Tsuneta et al. 2008), but not the first.
Consequently, an accurate quantitative estimate of the polar magnetic field remains a major and as
yet unattained goal.
The polar field is directly related to the dynamo process, presumably as a source of poloidal
field that is wound up by the differential rotation in the shear layer at the base of the convection
zone. The distribution of the magnetic field at the poles drives the formation and evolution of polar
coronal holes, polar plumes, X-ray jets, and other events and structures that characterize the polar
corona. Polar coronal holes have been intensively studied from the non-ideal vantage point offered
by the ecliptic, but never imaged from outside the ecliptic. Consequently, the distribution of the
polar field and the origin of polar structures are only poorly determined. The fast solar wind is
associated with open field lines inside coronal holes, whereas at least parts of the slow solar wind
are thought to emanate from the coronal hole boundaries. Understanding the interaction of open and
closed field lines across these boundaries is of paramount importance for elucidating the connection
between the solar magnetic field and the heliosphere.
The magnetic flux in the heliosphere varies with the solar cycle (Owens et al. 2008). There is
evidence that the heliospheric magnetic flux has increased substantially in the last hundred years,
perhaps by as much as a factor of two (Lockwood et al. 1999; Rouillard et al. 2007), possibly due to
a long-term change in the Sun’s dynamo action. As already noted, however, the interplanetary
magnetic field was dramatically lower than expected during the last solar minimum. Models based
on the injection of flux into the heliosphere by coronal mass ejections cannot explain this reduction,
and it is becoming clear that the processes by which flux is added to and removed from the
heliosphere are more complex than previously thought.
How Solar Orbiter will address this question:
Solar Orbiter’s comprehensive imaging instruments will characterize the properties and
dynamics of the polar regions for the first time, including magnetic fields, plasma flows, and
temperatures. Solar Orbiter will make the first reliable measurements of the amount of polar
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magnetic flux, its spatial distribution and its evolution (by comparing results from different orbits),
providing an independent constraint on the strength and direction of the meridional flow near the
pole. The evolution of Solar Orbiter’s orbit to higher heliographic latitudes will make it possible to
study the transport of magnetic flux from the activity belts toward the poles, which drives the
polarity reversal of the global magnetic field (see Wang et al. 1989; Sheeley 1991; Makarov et al.
2003). From its viewpoint outside the ecliptic, Solar Orbiter will probe the cancellation processes
that take place when flux elements of opposite polarity meet as part of the polarity reversal process.
Joint observations from Solar Orbiter and spacecraft in the ecliptic will determine, with high
accuracy, the transversal magnetic field, which is notoriously difficult to measure, along with
derived quantities such as the electric current density.
Solar Orbiter will measure the photospheric magnetic field at the poles, while
simultaneously imaging the coronal and heliospheric structure at visible and EUV wavelengths. In
addition, as the spacecraft passes through the mid-latitude slow/fast wind boundary at around 0.5
AU, the field and plasma properties of the solar wind will be measured. With the help of magnetic
field extrapolation methods these observations will, for the first time, allow the photospheric and
coronal magnetic field in polar coronal holes to be studied simultaneously and the evolution of polar
coronal hole boundaries and other coronal structures to be investigated.
Solar Orbiter’s observations from progressively higher heliographic latitudes (25° by the end
of the nominal mission) will enable the first coordinated investigation (jointly with spacecraft in the
ecliptic) of the three-dimensional structure of the inner heliosphere. These observations will reveal
the links between the Sun’s polar regions and the properties of the solar wind and interplanetary
magnetic field, in particular the heliospheric current sheet, which is used as a proxy for the tilt of the
solar magnetic dipole. In addition, Solar Orbiter will pass both north and south of the solar
equatorial plane in each orbit, with repeated transits through the equatorial streamer belt and
through the slow/fast wind boundary at mid-latitudes into the polar wind, making it possible to
follow the evolution of the solar wind and interplanetary magnetic field as well as of the sources in
the polar coronal holes. Ulysses has shown that poleward of the edge of coronal holes the properties
of the solar wind are relatively uniform, so that Solar Orbiter only needs to reach heliographic
latitudes just above the coronal hole edge to enter the high-speed solar wind. The orbital inclination
of 25° reached during the nominal mission is sufficiently high to satisfy this constraint.
4.2.1 Probability density function (PDF) of solar high-latitude magnetic field structures
4.2.1.1 Study the shape and morphology of solar high-latitude magnetic field structures as well as
their relation to the flow fields.
4.2.1.2 Study the distribution of field strengths (and magnetic vectors) of solar high-latitude
magnetic field structures.
4.2.1.3 Compare the results from the above objective (4.2.1.2) with the field strength distribution of
sunspots and active regions.
4.2.2 Basic properties of solar high-latitude magnetic field structures
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4.2.2.1 What is the intensity-field strength relation of polar magnetic field structures and how does it
change with the activity cycle?
4.2.2.2 What is the emergence rate and lifetime of polar magnetic field structures and how do they
change with the activity cycle?
4.2.2.3 Are there differences between the network at polar and at low latitudes?
4.2.3 Probe the structure in deep layers of the Sun
4.2.3.1 Probe the tachocline at the base of the convection zone, where the dynamo is surmised to be
situated.
4.2.3.2 Studies of convective, rotational and meridional flows.
4.2.3.3 MDI-like medium-l program
4.2.3.4 Stereoscopic helioseismology
3.2.4.4 4.3 What is the nature of magnetoconvection?
4.3.1 What are the velocity and magnetic vectors in the solar photosphere?
• Requires simultaneous observations from PHI and an instrument on the ground or another
satellite.
• Simultaneous observations at different viewing angles.
4.3.2 What is the 3D geometry of the solar surface in convective and magnetic features?
• Study the undulating 3D structure of the visible solar surface (Lites et al., 2004).
• Wilson depression appearing in magnetic structures such as faculae, pores and sunspots
(Solanki, 2003).
• Exact estimate of the height difference of the surface layer of different structures.
• Local helioseismology of active regions (Gizon and Birch, 2005).
4.3.3 How does the brightness of magnetic features change over the solar disk?
• Variation across the disk of the contrast of magnetic features relative to the quiet Sun (Topka
et al., 1997; Ortiz et al., 2002; Hirzberger and Wiehr, 2005).
4.3.4 What is the long-term behavior of active regions?
• Study sunspot evolution during co-rotation phase.
• Study the large-scale flows (e.g. sunspot moat flows) in and around active regions.
• Magnetic oscillations in sunspots and active regions.
4.3.5 How do magnetic fields emerge on the solar surface (coalescence, magnetic loops, convective
collapse?).
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• See e.g. Lamb et al., 2008; Centeno et al., 2007; Martínez González & Bellot Rubio, 2009;
Nagata et al., 2008.
3.2.4.5 4.4 Are there separate dynamo processes acting in the Sun?
Present state of knowledge:
MHD simulations indicate that a local turbulent dynamo should be acting in the Sun’s turbulent
convection zone (Brun et al. 2004) and even in the near-surface layers (Vögler and Schüssler 2007).
Hinode/SOT has detected ubiquitous horizontal magnetic fields in quiet regions of the Sun (Lites et
al. 2007), which are possibly generated by local dynamo action (Pietarila Graham et al. 2009).
These small, weak features (inter-network fields; Zirin 1987) bring 100 times more magnetic flux to
the solar surface than the stronger features that are known to be the product of the global dynamo,
and have themselves shown to be in cross-scale turbulent equilibrium (Schrijver et al. 1997). Even
the smallest observable features have been shown to be formed primarily by aggregation of yet
smaller, yet more prevalent features too small to resolve with current instrumentation (Lamb et al.
2008, 2009). It is, however, still uncertain whether a separate local, turbulent dynamo really is
acting on the Sun and how strongly it contributes to the Sun’s magnetic flux (and magnetic energy).
In particular, all solar magnetic features, from the smallest observable intergranular flux
concentrations to the largest active regions, have been shown (Parnell et al. 2009) to have a power
law (scale free) probability distribution function, suggesting that a single turbulent mechanism may
be responsible for all observable scales of magnetic activity.
How Solar Orbiter will address this question:
One way to distinguish between the products of a global and a local dynamo is to study the
distribution of small elements of freshly emerging magnetic flux over heliographic latitude. The
global dynamo, presumably owing to the structure of the differential rotation and the meridional
flow near the base of the convection zone, leads to the emergence of large bipolar magnetic regions
(active regions) at the solar surface at latitudes between 5° and 30° and of smaller ephemeral active
regions over a larger range of latitudes, but concentrated also at low latitudes. In contrast, a local
turbulent dynamo is expected to enhance field more uniformly across the surface.
Observations carried out from the ecliptic cannot quantitatively determine the latitudinal distribution
of magnetic flux and in particular the emergence of small-scale magnetic features (inter-network
fields) due to foreshortening and the different sensitivity of the Zeeman effect to longitudinal and
transversal fields. Solar Orbiter, by flying to latitudes of 25° and higher above the ecliptic, will be
able to measure weak magnetic features equally well at low and high latitudes (Martínez Pillet
2006). If the number and size (i.e., magnetic flux) distributions of such features are significantly
different at high latitudes, then even the weak features are probably due to the global dynamo. If,
however, they are evenly distributed, then the evidence for a significant role of a local dynamo will
be greatly strengthened. Current work is confounded by viewing angle restrictions near the poles, by
the ubiquitous seething horizontal field (e.g., Harvey et al. 2007), and by small deflections in near-
vertical fields, which dominate observed feature distributions near the limb of the Sun.
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Detailed sub-objectives:
4.4.1 Compare the distribution of small-scale fields at low and high latitudes (Tsuneta et al.,
2008; Ito et al., 2010).
4.4.1.1 Is the magnetic network equally strong between low and high latitudes? How does this
distribution change between activity maximum and minimum?
4.4.1.2 What is the latitude distribution of the emergence of ephemeral regions? Is this distribution
dependent on the cycle phase?
The emergence, diffusion and decay of ephemeral regions near the poles and below high-latitude
coronal holes should be studied for the aspect of how they feed the magnetic network (see e.g.
Simon et al. 2001, ApJ 561, 49 427; Gosic et al. 2014, ApJ 797). In particular, the latitudinal
dependence of this decay process would be interesting to study.
4.4.1.3 What is the latitude distribution of internetwork magnetic fields and their emergence rates?
Do they depend on the phase of the solar cycle?
4.4.1.4 What is the latitudinal distribution of the linear-polarization features in the quiet Sun? Is
there a solar cycle dependence?
4.4.2 Joy’s law at high latitudes
4.4.2.1 Extend Joy’s law at low latitudes from large to small bipolar features, i.e. from active
regions to ephemeral regions.
4.4.2.2 Compare the tilt angle distributions of ephemeral regions at low and high latitudes. How do
these distributions change over the cycle?
4.4.3 Understand differences in size and internal structuring of magnetic concentrations in
high and low latitudes and relate them to different origins or dynamic environments (e.g. Lagg
et al., 2010; Martínez González, 2012; Requerey et al., 2014; 2015).
3.2.4.6 4.5 How are coronal and heliospheric phenomena related to the solar dynamo?
4.5.1 Explore the quasi-biennial modulation of galactic cosmic rays (Laurenza et al., 2012) and of
flare-CME onset (Telloni et al., 2015).
4.5.2 Explore the effect of the solar polarity reversal (Gnevyshev gap) on the heliosphere during
solar cycle 25 (Storini et al., 2003) and compare with that of previous cycles.
4.5.3 Explore the possible effect of the Centennial Gleissberg Cycle on the heliosphere during solar
cycles 24 and 25 (Feynman & Ruzmaikin, 2014).
4.5.4 Determine the solar wind, magnetic field, energetic particles and radio emission properties
during solar cycle 25 and compare to those found by Helios.
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3.2.5 5. Additional science objectives
This section will be significantly modified: many of the sub-objectives are not additional but instead
part of Objectives 1-4, they have to be reintegrated to the relevant sections.
5.1 Additional Science Objectives of EUI
5.1.1 Study the corona and its phenomena in a high spatial and temporal scale (active regions
during flares or in quiescent conditions, coronal holes, quiet Sun)
Whatever the active region target, the high spatial resolution of HRI is used in a mode close
to A mode with a time resolution of the order of 1 sec. For the polar coronal hole mode, the high
latitude is mandatory and the compression in Lalpha possibly lower than 15 (mode
C). Complementary observations of SPICE in Lbeta would be very useful.
5.1.2 What is the Ly-α emission and absorption in the cool atmosphere (especially in polar
coronal holes)? To what extent is it a function of latitude?
The observing mode is C but the observations must be repeated during the various orbits
(always at perihelion) in order to cover a large range of latitudes. Complementary observations of
SPICE in Lbeta would be very useful.
5.1.3 Study the coronal He abundance
The S mode should be OK since there is overlap between METIS FOV and FSI (304). One
could modify this mode in three ways: not cropping to 4Rs by 4Rs (in order to increase the FOVs
overlap)), decreasing the cadence (down to about an hour) and decreasing the compression by a
factor 5.
5.1.4 Reconstruct the solar EUV irradiance (Haberreiter et al., 2014) for all hemispheric
directions, in particular for higher latitudes
Synoptic F174 (and additionally F304) at a cadence of 1/day; high latitude observations are
of particular interest; coordination with other imagers in space, if possible.
5.1.5 Targets of opportunity (planetary quadratures, comets, …).
For Jupiter and Saturn, the most obvious targets are the aurorae (see observations with the
Hubble Telescope) to be observed in Lalpha. The visibility will depend on the planet-SO distance.
The signal is time variable (from a few kR to MR, Nichols et al. 2007) and detectable (at the typical
distance of 4 AU for Jupiter). The disk would cover about 30*30 pixels of HR Lalpha. This assumes
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an out-of-limb offset pointing … To perform further S/N computations, it is necessary to know the
pointing stability.
If one refers to Bennett comet observations (Bertaux, Blamont, Festou 1973) performed at
distances smaller than 1 AU, the Lalpha signal should be about 10^^-6 the solar signal with a very
large field. This requires out-of-limb offset pointing of HRI Lalpha with perhaps a “reconnaissance”
imaging in FSI 30.4 nm. To perform further S/N computations, it is necessary to know the pointing
stability. If SO is offset, METIS could observe the extended bubble of the comet in Lalpha.
5.2 Additional Science Objectives of EPD
5.2.1 Energetic particles in Venus and Earth environment
5.2.2 Galactic and anomalous cosmic rays. Long-term and short-term modulation and spatial
variation.
5.2.3 Energetic neutral atoms (Wang et al., 2014). (How does EPD team mean to address this?)
5.2.4 Jovian electrons
5.3 Additional Science Objectives of MAG
N/A
5.4 Additional Science Objectives of METIS
5.4.1 Hydrogen Ly-α emission by the atmosphere of planets (e.g. Venus, Jupiter) (Chaufray et
al., 2012; Colwell, 1998; Jaffel et al., 2007; Menager et al., 2010).
5.4.2 Study of Sungrazing comets
5.4.2.1 Understand cometary properties and evolution by mapping the hydrogen Ly-α emission,
proportional to the outgassing rate, along its trajectory close to the Sun (Raymond et al., 1998; Uzzo
et al., 2001; Bemporad et al., 2005; Ciaravella et al., 2010).
5.4.2.2 Investigate the fragmentation of the cometary nucleus from the variation with the
heliocentric distance of the outgassing rate and from the splitting of the cometary tail.
5.5 Additional Science Objectives of PHI
(All of the following objectives require simultaneous measurements from another satellite or from
the ground.)
5.5.1 How strongly does the solar luminosity vary and what is the source of these variations?
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5.5.1.1 Provide magnetic field and continuum brightness data in order to determine Sun’s irradiance
at different heliospheric latitudes and phases of the solar cycle (Krivova et al., 2003).
5.5.1.2 Comparison with Sun-like stars’ larger irradiance variations (Schatten et al., 1993; Knaack et
al., 2001):
• Is it a line of sight effect relative to their rotation axes or
• Will the Sun also display larger variations in the future?
5.5.2 Other Science Objectives
5.5.2.1 Effect of granulation and oscillations, i.e. interaction of modes and convection (stereoscopic
helioseismology).
5.5.2.2 Two components of velocity. What is the relationship between the components of the
velocities in granulation? Supergranulation? Various modes in quiet Sun? (stereoscopic
helioseismology)
5.5.2.3 Shape of the Sun. Extend oblateness study of (Emilio et al., 2007; Kuhn et al., 2012).
5.6 Additional Science Objectives of RPW
5.6.1 Interplanetary dust: spatial distribution, mass and dynamics.
• Required observations: Electric signatures of dust impact on the spacecraft body (RPW).
5.7 Additional Science Objectives of SoloHI
5.7.1 Interplanetary Dust
• What are the sources and properties of dust in the inner heliosphere?
• Do Sun-Grazing comets contribute to the dust and what is the time dependence?
• Does the scattering function change with heliocentric distance?
• Can the evaporation of dust be detected?
5.7.2 Streamers
• What is the 3-D structure and extent of streamers?
• Can a filamentary nature be detected?
• What is their radial distribution?
• What is the difference between pseudo- and normal streamers?
5.7.3 CMEs
• What is the 3-D structure and extent of CMEs?
• Are flux ropes imbedded in the plasma sheet?
• What are the radial and longitudinal distributions of "blobs", magnetic islands/plasmoids?
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5.8 Additional Science Objectives of SPICE
5.8.1 Obtain spectral atlas of representative features in the solar atmosphere.
5.9 Additional Science Objectives of STIX
TBW
5.10 Additional Science Objectives of SWA
5.10.1 What is the temporal variability of 3D thermal particle distributions?
• Burst mode resolution of EAS and PAS will provide fastest thermal distributions recorded at
different radial distances in the solar wind. At high cadence, test:
• how anisotropic and irregular are proton and electron distributions? This can test whether
resonant heating or stochastic heating is the dominant heating mechanism.
• how does the anisotropy change in time / space?
• how are distributions related to local turbulence and global expansion?
5.10.2 How are proton and helium temperatures related to their relative drift, is there any evidence
for resonant heating (Kasper et al., 2013)?
5.10.3 How common are proton beams, where do they come from and what impact do they have on
ambient conditions (wave generation, heating)?
5.10.4 Identify and characterize the various forms of free energy in the particle distribution
functions. Analyze their formation and relaxation processes at sub-second resolution. Characterize
and measure the degree of reversibility and irreversibility in these processes and generalize to the
understanding of the dynamic behavior of weakly collisional media.
5.10.5 Fully characterize the radial and latitudinal variability of 3D thermal particle distributions
under different solar wind conditions (fast, slow wind, in and around CMEs, CIRs, etc.)
5.10.6 Confirm the origin of the electron halo population – is it derived from the scattering of the
strahl (c.f. Maksimovic et al., 2005) and if so what are the physical mechanisms responsible for this
and their radial and latitudinal dependencies?
5.10.7 How do the combined electron strahl and composition signatures vary across reconnecting
solar wind current sheets? Are these consistent with expectations of interchange reconnection
models? What are the combined signatures in other solar wind structures (e.g. shocks, CIRs,
CMEs), and how definitive are these in determining solar origins and/or defining processes
occurring in the solar wind?
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4 INSTRUMENT DESCRIPTIONS AND THEIR OPERATIONAL CONSTRAINTS
4.1 EPD
The EPD (Energetic Particle Detector) investigation on Solar Orbiter provides the key in situ
measurements of energetic particles to address the major science objective 3 (How do solar
eruptions produce energetic particle radiation that fills the heliosphere?). EPD also addresses the
other major science objectives by tracking magnetic field lines from the spacecraft to their solar
source, by determining magnetic field line lengths and measuring in situ the interplanetary effects of
evolving Interplanetary Coronal Mass Ejections (ICMEs), and by measuring the latitude and
longitudinal distribution of Solar Energetic Particles (SEPs).
The Energetic Particle Detector suite consists of four sensors measuring electrons, protons, and
heavy ions from helium to iron, and operating at partly overlapping energy ranges from 2 keV up to
450 MeV/n.
The EPD sensors are:
a) Suprathermal Electrons and Protons (STEP)
b) Suprathermal Ion Spectrograph (SIS)
c) Electron Proton Telescope (EPT)
d) High Energy Telescope (HET)
EPT and HET are combined in one unit, and there are two identical EPT-HET units on the
spacecraft.
STEP relies on a magnetic deflection system in one of the two detector units to suppress electrons in
that particular unit. The particles are measured by the use of solid state detectors with ultra-thin
ohmic contacts. SIS applies the time-of-flight by energy technique to measure the composition of
ions. EPT relies on the magnet/foil technique in order to separate electrons from protons and heavier
ions. HET is based on the multiple dE/dx vs. total E technique.
4.1.1 EPD Observables
The relevant SEP characteristics that should be taken under consideration are the following:
1. Energetic Particle Energy
2. Particle Mass
3. Particle Flux Angular information
4. Energetic particle flux temporal information.
These, in terms of instrument parameters, can be written as follows:
a) Energy Range
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b) Mass resolution (Ion composition)
c) Fields of View and Angular Resolution
d) Geometric Factor
e) Time Resolution
4.1.1.1 Energy range
The Energy Range is the lowest and highest particle energy that the instrument is capable to
measure. It is defined for the different particle species:
• Electrons (e)
• Protons (p)
• Heavy ions from Helium to Iron
The EPD sensors performance regarding particle energy is given in the following table (Table
4.1.1.1.1) and its energy coverage for the different species in Figure 4.1.1.1.1.
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4.1.1.2 Mass resolution
The mass resolution defines both the closest elements, and isotopes, that the instrument is capable to
resolve.
The EPD performance regarding particle mass is given in Table 4.1.1.2.1.
4.1.1.3 Fields of view and angular resolution
The Field of View (FoV) is that portion of the sky that is covered by the cross-sectional area
spanned by the triggering detector and the front defining detector or aperture of the various EPD
sensors. Depending on the EPD sensor, its cross section can be either circular or rectangular.
The angular resolution of a telescope is defined as its ability to distinguish different directions of the
incoming particles.
Here, we will provide information about the number of apertures, both in and out of the ecliptic
plane, the “FoV” sizes and the instrument angular resolution.
The directional information performance of EPD is given in Table 4.1.1.3.1 Figure 4.1.1.3.1 shows
the FoV of EPD referred to the S/C reference frame. Note the clustering of FoVs around the mean
Parker spiral direction.
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4.1.1.4 Geometric factor
The geometric factor is defined as the integral over the telescope FoV of the quantity dSd (where
S is the telescope/detector surface and is the solid angle that it subtends). The geometric factor is
the parameter that allows the computation of differential intensities from measured quantities. In
practice, the geometric factor is the factor of proportionality relating the counting rates to the
intensity in the field of view. Geometric factors for each sensor have been optimized after a careful
analysis of the expected fluxes close to the Sun over the very wide energy range covered by EPD.
Table 4.1.1.4.1 shows the geometric factor of the different EPD sensors.
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4.1.2 EPD modes
4.1.2.1 EPD Normal mode
TBW
4.1.2.2 EPD Burst mode
TBW
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4.2 EUI
The EUI instrument suite is composed of two high resolution imagers (HRI), one at Lyman-α and
one at 174Å, respectively named “HRI_Lyalpha” and “HRI_EUV”, and one dual band full-sun
imager (FSI) working alternatively at the 174 and 304 Å EUV passbands, named “FSI174/304”.
4.2.1 EUI observables
4.2.2 EUI modes and telemetry
TM figures
Allocated TM 20.5 Kbits/s
Download capacity per orbit 6.642 GB = 53.136Gbits
Subtelescopes/units that can be commanded independently
wavelength detector size FOV TM/raw image (15bits) compression rates
FSI 174nm or 304nm 3072 px * 3072 px 3.8ºx3.8º 17.7 MB 6-500
HRI_Lya 121.6nm 2048 px * 2048 px 16.6'x16.6' 7.9 MB 6-50
HRI_EUV 174nm 2048 px * 2048 px 16.6'x16.6' 7.9 MB 6-50
Low Latency programs
Observing mode Instrument Rebin/subfield #Images/h Data Rate Data Rate Max time
(Gbits/h) (Kbits/s) (h)
EUI/FSI Beacon mode (B) FSI 174 768x768 4
0,0005 0,146 unlimited
FSI 304 768x768 4
EUI/HRI Beacon modes (HB) HRI_EUV 2048x2048 0,0417
0,0001 0,029 unlimited
HRI_Lyalpha 2048x2048 0,0417
(1 set/day)
Note that the Beacon mode will run in parallel to most of the FSI programs defined below.
Main science programs
Observing mode Instrument Rebin/subfield #Images/h Data
Rate
Data
Rate
Max time
***
(Gbits/h) (Kbits/s) (h)
FSI Synoptic mode (S) *
FSI 174 1536x1536 4
0,0056 1,56
unlimited
FSI 304 1536x1536 4
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FSI Reference Synoptic mode (R)
* FSI 174 3072x3072 0,0417
0,00022 0,061 unlimited
FSI 304 3072x3072 0,0417
(1 set/day)
FSI Global eruptive event mode
(G) **
FSI 174 or
304 3072x3072 360 3,8 1054,7 12 h
FSI Faint High Corona (FHC) ** FSI 174 or
304 3072x3072 4 0,021 5,86 unlimited
EUV & LYA Coronal hole modes
(C)
HRI_EUV 2048x2048 120 1,5 416,7
30 h
HRI_Lyalpha 2048x2048 120
EUV & LYA Quiet Sun modes (Q) HRI_EUV 2048x2048 450 14,1 3906,3
3 h
HRI_Lyalpha 2048x2048 3600
EUV & LYA Active Region modes
(A)
HRI_EUV 2048x2048 1800 16,9 4687,5
3 h
HRI_Lyalpha 2048x2048 3600
* Note that all FSI modes will typically be combined with the Beacon mode (B) defined above to
assure regular EUI LL data. The data rates above do not yet include this extra TM, total data rate is
given in the mode-specific pages (see e.g. FSI Synoptic mode (S)).
** Some more specific FSI modes will have both Beacon mode (B) and Synoptic mode (S) running
in parallel. The data rates above do not yet include this extra TM, total data rate is given in the
mode-specific pages (see e.g. FSI Global eruptive event mode (G)).
*** Maximum time for this mode to run per orbit is based on the average EID-A rate over the 30
days RS windows. There will be times though that more downlink is available, and at other times
more stringent restrictions will apply.
Additional science programs
Observing mode Instrument Rebin/subfield #Images/h Data Rate Data Rate Max time
(Gbits/h) (Kbits/s) (h)
EUV & LYA Eruptive Event modes
(E)
HRI_EUV 2048x2048 3600 22,5 6250
2 h
HRI_Lyalpha 2048x2048 3600
EUV & LYA Discovery modes (D) HRI_EUV 2048x2048 3600 22,4 6225
2 h
HRI_Lyalpha 645x645 36000
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In the following sections, we present in detail the EUI modes.
4.2.2.1 FSI Beacon mode (B)
This mode will be used with FSI only to produce (part of) the EUI Low Latency data. It generates
highly compressed synoptic data, with extremely restricted TM.
• Parameters
TBC
• Resource usage
Observing
mode Target Instrument Channel Cadence Rebinning Compression
Calculated
TM #Images/h Data Rate
Data
Rate
Max
time
(nm) (s) or subfield rate (bits/image) (Gbits/h) (bits/s) per orbit
FSI
Beacon
mode (B)
= highly
compressed FSI 17,4 900 4 125 70778.88 4 0.00053 157.3 unlimited
synoptic
data FSI 30,4 900 4 125 70778.88 4
• Maximum time in this mode
unlimited
whole orbit would take ~4% of allocation
30 days takes 0.8%
4.2.2.2 FSI Synoptic mode (S)
This mode is using the FSI full disk telescope and provides synoptic full disk images (full sun up to
4x4 solar radii).
Depending on the distance from the sun, images are either rebinned 2x2 (close to Sun) or a quarter
of the detector is read out (subfield, far from Sun). Both choices result in same data rate.
In practice, this EUI mode will run in parallel to FSI Beacon mode (B) to ensure availability of
regular EUI/FSI Low Latency data.
• Parameters
Cadence and compression rate (TBC)
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• Resource usage
Observing mode Target Instrument Channel Cadence Rebinning Compression
rate #Images/h
Data
Rate
Data
Rate
Max time per
orbit
(nm) (s) or
subfield (Gbits/h) (bits/s)
Synoptic mode (S) Full sun up
to 4x4R
FSI 17,4 900 2x2 46.88 4 0.0056 1677.7
unlimited
FSI 30,4 900 2x2 46.88 4
FSI Synoptic + FSI
Beacon (S+B) 0.0062 1835.0 unlimited
• Maximum time in this mode
unlimited
3 RSWs = 10% allocation
full orbit (168 days) = 56% allocation
4.2.2.3 FSI Reference Synoptic mode (R)
This mode is using the FSI full disk telescope and provides synoptic full disk images (full sun up to
4x4 solar radii).
Full sun images are obtained with better resolution and less compressed than in FSI Synoptic mode
(S). Typically, only 1 set of those is planned per day (cadence 1/24hrs).
• Parameters
Cadence and compression rate (TBC!!)
• Resource usage
Observing
mode Target Instrument Channel Cadence Rebinning Compression
Calculated
TM #Images/h Data Rate
Data
Rate Max time
(nm) (s) or subfield rate (bits/image) (Gbits/h) (bits/s) per orbit
FSI
Reference
synoptic (R)
4x4Rsun
window
FSI 17,4 86400 2x2 12.5 30670848 0,0417
0.00022 65.5
unlimited
FSI 30,4 86400 2x2 12.5 30670848 0,0417
FSI
Reference
Synoptic
(R) + FSI
Beacon (B)
0.00075 222.8 unlimited
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• Maximum time in this mode (R)
unlimited
30 image duos (1 per each RSW day) = 0.01% allocation
full orbit (168 days) = 0.07% allocation
4.2.2.4 FSI Global eruptive event mode (G)
This mode is using the FSI full disk telescope and aims at capturing a global event.
• Parameters
Channel, cadence and compression rate (TBC!!)
• Resource usage
Observing
mode Target Instrument Channel Cadence Rebinning Compression
Calculated
TM #Images/h Data Rate
Data
Rate
Max
time
(nm) (s) or subfield rate (bits/image) (Gbits/h) (bits/s) per
orbit
Global eruptive
event 174
(G174)
Full
event FSI 17.4 10 1 10 11324621 360 3.80 1132462 12 h
Global
eruptive event 304
(G304)
Full event
FSI 30.4 10 1 10 11324621 360 3.80 1132462 12 h
Global
eruptive
event 174
(G174)
+ FSI
Synoptic
(S) + FSI
Beacon (B)
3.803 1134140 12 h
Global
eruptive
event 304
(G304)
+ FSI
Synoptic
(S) + FSI
Beacon
(B)
3.803 1134140 12 h
• Maximum time in this mode
restricted to 12 hours per orbit = 92% of allocation
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4.2.2.5 FSI Find Event mode (FE)
This mode is using the FSI full disk telescope to detect events onboard. The FSI thumbnails used for
event detection will not be brought to ground so the only TM resource usage in this mode is the
generation of parallel synoptic imagery and beacon (LL) imagery.
• Parameters
TBW
• Resource usage
Observing mode Target Instrument Channel Cadence Rebinning Compression #Images/h Data
Rate
Data
Rate
Max
time
(nm) (s) or
subfield rate (Gbits/h) (bits/s) per orbit
FSI Find Event (FE174) Event
detection FSI 0 0 unlimited
FSI Find Event (FE174)
+ FSI Synoptic (S) + FSI
Beacon (B)
0.006 1835 unlimited
• Maximum time in this mode
unlimited
4.2.2.6 FSI Faint High Corona mode (FHC)
This mode is using the FSI full disk telescope, with the disk occulter in the FSI filter wheel. It aims
at capturing the faint high corona off-limb.
• Parameters
Channel, cadence and compression rate (TBC)
• Resource usage
Observing mode Target Instrument Channel Cadence Rebinning Compression Calculated
TM #Images/h
Data
Rate
Data
Rate
Max
time
(nm) (s) or
subfield rate (bits/image) (Gbits/h) (bits/s) per orbit
Faint High Corona
174 (FHC174)
Full
event FSI 17.4 3600 1 6.25 22649241.6 1 0.0211 6291.5 unlimited
Faint High Corona
304 (FHC304)
Full
event FSI 30.4 3600 1 6.25 22649241.6 1 0.0211 6291.5 unlimited
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Faint High
Corona 174
(FHC174)
+ FSI Synoptic
(S) + FSI Beacon
(B)
0.0267 7969.0 unlimited
Faint High
Corona 304
(FHC304)
+ FSI Synoptic
(S) + FSI
Beacon (B)
0.0267 7969.0 unlimited
• Maximum time in this mode
unlimited
the combined mode FHC+S+B, running through all RS windows (30 days) would consume ~40% of
the orbital allocation
4.2.2.7 EUV & LYA Beacon modes (HB)
This is HRI's Low Latency programme that produces one set of small FOV images per day.
The mode can be commanded on HRI_EUV and HRI_Lya telescopes independently.
• Parameters
TBW
• Resource usage
Observing
mode Target Instrument Channel Cadence Rebinning Compression #Images/h
Data
Rate
Data
Rate
Max
time
(nm) (s) or
subfield rate (Gbits/h) (bits/s) per orbit
EUV Beacon
mode
Can be run on any target
HRI_EUV 17,4 86400 1 47 1/day 0.00005 15.5 unlimited
LYA Beacon
mode HRI_Lyalpha Ly alpha 86400 1 47 1/day
• Maximum time in this mode
unlimited
4.2.2.8 EUV & LYA Coronal hole modes (C)
This mode uses HRI telescopes targeting a full coronal hole, a CH boundary or polar plumes.
The mode can be commanded on HRI_EUV and HRI_Lya telescopes independently.
• Parameters
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cadence, compression TBC
• Resource usage
Observing
mode Target Instrument Channel Cadence Rebinning Compression
Calculated
TM #Images/h
Data
Rate
Data
Rate
Max
time
(nm) (s) or
subfield rate (bits/image) (Gbits/h) (bits/s)
per
orbit
EUV Coronal
hole (C)
Full CH +
boundary HRI_EUV 17,4 30 1 6.25 10066329.6 120 1,5 447392.4 30 h
LYA Coronal
hole (C) + plumes HRI_Lyalpha Ly alpha 30 1 18.75 3355443.2 120
• Maximum time in this mode
restricted to 30 hours per orbit = 90% of allocation (if both telescopes are used), if all resulting TM
gets downlinked
4.2.2.9 EUV & LYA Quiet Sun modes (Q)
This mode uses HRI telescopes targeting a region of quiet sun. It is most appropriate to be used
close to perihelion, near co-rotation.
The mode can be commanded on HRI_EUV and HRI_Lya telescopes independently.
• Parameters
TBW
• Resource usage
Observing
mode Target Instrument Channel Cadence Rebinning Compression #Images/h
Data
Rate
Data
Rate
Max
time
(nm) (s) or
subfield rate (Gbits/h) (bits/s)
per
orbit
EUV Quiet Sun (Q)
Full HRI FOV
of quiet sun
HRI_EUV 17,4 8 1 9.4 450
14.06 4194304.0 3 h LYA Quiet Sun
(Q) HRI_Lyalpha Ly alpha 1 1 18.75 3600
• Maximum time in this mode
TM-hungry mode, limited to 3 hours per orbit = 85% of EID-A allocation (if both telescopes are
used), if all resulting TM gets downlinked!
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4.2.2.10 EUV & LYA Active Region modes (A)
This mode uses HRI telescopes targeting an active region on the sun. It is most appropriate to be
used close to perihelion, near co-rotation.
The mode can be commanded on HRI_EUV and HRI_Lya telescopes independently.
• Parameters
TBW
• Resource usage
Observing mode Target Instrument Channel Cadence Rebinning Compression #Images/h Data
Rate
Data
Rate
Max
time
(nm) (s) or
subfield rate (Gbits/h) (bits/s)
per
orbit
EUV Active
Region (A) Full HRI FOV of active region
HRI_EUV 17,4 2 1 18.75 1800
16.88 5033165 3 h LYA Active
Region (A) HRI_Lyalpha Ly alpha 1 1 18.75 3600
• Maximum time in this mode
TM-hungry mode, limited to 3 hours per orbit (if both telescopes are used), if all resulting TM
gets downlinked!
4.2.2.11 EUV & LYA Eruptive Event modes (E)
This mode uses both HRI telescopes targeting an eruptive event on the sun. It is most appropriate to
be used close to perihelion, near co-rotation.
The mode can be commanded on HRI_EUV and HRI_Lya telescopes independently.
• Parameters
• Resource usage
Observing mode Target Instrument Channel Cadence Rebinning Compression #Images/h Data
Rate
Data
Rate
Max
time
(nm) (s) or
subfield rate (Gbits/h) (bits/s)
per
orbit
EUV Eruptive event
(E)
Full FOV on event
HRI_EUV 17,4 1 1 18.75 3600
22.5
6710886
2 h LYA Eruptive event
(E) HRI_Lyalpha Ly alpha 1 1 18.75 3600
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• Maximum time in this mode
TM-hungry mode, limited to 2 hours per orbit = 98.5% of allocation, if all resulting TM gets
downlinked!
4.2.2.12 EUV & LYA Discovery modes (D)
This mode uses HRI telescopes, with the Ly Alpha channel in its highest cadence 0.1s, to capture
high cadence dynamics outside of flaring times. HRI EUV is used with full FOV, in Lyalpha only
a subfield is read out. This mode is most appropriate to be used close to perihelion, near co-rotation.
The mode can be commanded on HRI_EUV and HRI_Lya telescopes independently.
• Parameters
• Resource usage
Observing mode Target Instrument Channel Cadence Rebinning Compression #Images/h Data
Rate
Data
Rate
Max
time
(nm) (s) or
subfield rate (Gbits/h) (bits/s)
per
orbit
EUV Discovery
(D)
High cadence
dynamics
HRI_EUV 17,4 1 1 18.75 3600
22.5 6683643 2 h LYA Discovery (D)
HRI_Lyalpha Ly alpha 0,1 3,175 18.75 36000
• Maximum time in this mode
TM-hungry mode, limited to 2 hours per orbit = 90% of allocation (for both telescopes together), if
all resulting TM gets downlinked!
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4.3 MAG
The measurement of magnetic fields in space science is constrained by environmental
factors such as temperature and radiation, and also by restrictions on resources available for the
instrument such as weight and power. The fluxgate magnetometer principle has emerged as an
optimum compromise method for magnetic field measurement in space, since such instruments are
rugged, low in power and mass and offer high precision.
The basic operating principle of the fluxgate magnetometer is well known and documented.
A soft magnetic core, usually toroidal in shape, is wound with a coil and driven into saturation with
an AC excitation current. The external magnetic field (which is to be measured) distorts the
symmetry of the magnetic flux in the core, which generates a signal at even harmonics of the
excitation frequency. This signal is detected by a sense coil wound around the core. By feeding back
a current into the sense coil proportional to the measured signal, the ambient field is backed-off and
the sensor operates in null-mode, thereby improving linearity.
Typically, the second harmonic of the drive frequency is notch-filtered, amplified,
synchronously detected, integrated, and used to drive the feedback current. It is this current which is
proportional to the ambient magnetic field. The output of a traditional analogue magnetometer is the
voltage used to drive this feedback current and is usually digitized by an analogue-to-digital
converter. In such a design, the signal processing to extract the second-harmonic signal is performed
by an analogue circuit. This is an efficient and well-developed method, however the filtering and
synchronous detection stages all show temperature dependencies, and the analogue component
count and mass are not negligible. By replacing the analogue processing with a digital system, this
signal processing is performed as a software function in the digital domain. As well as eliminating
some of the temperature effects, and reducing mass, this provides an inherent flexibility in the
design, since the processing parameters can be modified by software.
The MAG instrument performance is summarized in the following table:
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4.3.1 MAG observables
4.3.2 MAG modes
4.3.2.1 MAG Normal mode
Burst modes run in parallel with normal modes. Here, HK and LL TM are included in the normal
mode data rate.
Mode IBS Cadence (vps) OBS Cadence (vps) Data Rate (bps) Note
Normal 1 16 1000 Default Option
Equal8 8 8 950
Low 1 1 250
4.3.2.2 MAG burst mode
Burst modes run in parallel with normal modes. Here, HK and LL TM are included in the normal
mode data rate. Burst mode rates here represent the additional TM from burst data. In principle, any
normal mode can be combined with any burst mode, although not all combinations make scientific
sense unless selective downlink is implemented for MAG.
Mode IBS Cadence (vps) OBS Cadence (vps) Data Rate (bps) Note
Burst 8 128 6892 Normal + Burst = 7892
Burst1 1 128 6542 Normal + Burst1 = 7542
Equal16 16 16 1692 Normal + Equal16 = 2692
Equal128 128 128 12924 Normal + Equal128 = 13924
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4.4 METIS
METIS is an inverted-occultation coronagraph for the simultaneous observation of the whole
Sun corona in the visible light and around the UV wavelength of the H I Lyman- spectral line.
METIS will be capable of obtaining for the first time simultaneous imaging of the full corona in
polarized visible light (580-640 nm) and narrow-band ultraviolet H I Ly-α (121.6±10 nm). These
measurements will allow a complete characterization of the most important plasma components of
the corona and the solar wind, i.e. electrons and protons.
METIS can simultaneously image the visible and ultraviolet emission of the solar corona
and diagnose, with unprecedented temporal coverage and spatial resolution (down to about 4000
km), the structure and dynamics of the full corona in the range from 1.6 to 3.0 solar radii (R☉) at
minimum perihelion (0.28 AU), and from 2.8 to 5.5 R☉ at 0.5 AU. This region is crucial in linking
the solar atmosphere phenomena to their evolution in the inner heliosphere, and the study of its
properties is very important in meeting the Solar Orbiter fundamental science goals.
The instrument has been conceived to perform off-limb and near-Sun coronagraphy, with the
aim of addressing the three key scientific issues concerning:
• the origin and acceleration of the fast and slow solar wind streams;
• the origin, acceleration and transport of the solar energetic particles;
• the transient ejection of coronal mass (coronal mass ejections, CMEs) and its evolution in the
inner
heliosphere.
In addition, METIS can contribute to the study of the properties and evolution of Sun-
grazing comets. It is also able to characterize the properties of the coronal regions crossed by the
Solar Probe Plus spacecraft, during
its transit close to the Sun.
4.4.1 METIS observables
The expected science products of the off-limb and near-Sun coronagraphy performed by METIS
are:
• global maps of VL coronal emission in the range 580-640 nm for each of the 4 polarization states
of the
Liquid Crystal Variable Retarder
• global maps of the H I Ly-α (121.6 nm) UV coronal emission.
The above measurements can be obtained simultaneously on board.
On ground, the VL channel measurements can be combined to obtain maps of either the total and/or
the
polarized brightness (tB/pB) of the corona.
In addition, 8 VL light curves, 1 for each sector of the field of view, can be produced (at the
moment, only as
Low Latency data during Coronal Mass Ejections observations).
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Further processing of calibrated Ly-α and pB images can provide:
• global maps of electron densities, from the pB maps;
• global maps of the outflow velocity of the H component of the solar wind (by Doppler dimming
technique.
4.4.2 METIS modes
Note that all data volumes & rates are based on powers of 10, i.e. 1Gbit = 109 bits.
TM figures
Allocated TM 10.5 Kbits/s
Download capacity per orbit 3.4 GB = 27.2 Gbits
Subtelescopes/units that can be commanded independently
wavelength detector size FOV TM/raw image compression rates
METIS_VL polarized VL in 580-640nm (4 images) 2048 px * 2048 px annular FOV 1.5º-2.9º 7.34 MB (14bit/px) 13-50
METIS_UV 121.6nm (Ly_beta) 1024 px * 1024 px annular FOV 1.5º-2.9º 1.84 MB (14bit/px) 13-50
Full set of raw METIS images (4 VLs+1UV) accounts for 31.20MB.
If partially* rebinned and compressed (factor 2.5), a full METIS dataset accounts for 4.6MB to
16MB depending on the science goal.
* METIS images are rebinned differently in the inner band of the FOV as in the outer band:
typically, with binning factor 2 < 2.5º and 4 >2.5º for global corona and binning factors 2 < 2.5º
and 4 >2.5º for smaller scale features like CMEs and oscillations
Observational modes
Main science programs
Observing mode Use case Telescope(s) Cadence Data Vol /
set
Data
Rate
Max
Duration
(min) (Mbits) (Kbits/s)
METIS standard
modes
WIND
electron density measurement and solar wind
velocity
<0.5AU (1.6 - 5 solar radii)
METIS_VL+UV 5-30min
14.12Mb
(4VL+1UV)
8 - 47 days
MAGTOP
slow wind source regions (magnetic topology)
near perihelion (1.6 - 5 solar radii)
METIS_VL+UV 5-20min 23.34Mb
(4VL+1UV) 20 - 78 days
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GLOBAL
Global corona configuration/evolution during
CME events
close or far from sun, binning dependent on
distance
METIS_VL+UV 5-30min 12Mb - 8Mb
(4VL+1UV) 4.5 - 40
days
LT-CONFIG
Long-term evolution of coronal configuration
close or far from sun, binning dependent on
distance
METIS_VL+UV 20-30min 4.99Mb
(4VL+1UV) 2.8 - 4.2 days
METIS special
modes
FLUCTS+TBF
Brightness fluctuations spectra
at perihelion only (1.6 to 3 solar radii)
METIS_VL only
3 step
acquisition,
variable cadence:
1s - 20s -
10min
4.67Mb
(1 VL image)
311
(avg/hr)
hours
CMEOBS
CME driven shocks and SEP (filamentary
structure)
triggered by METIS-internal flag (at any
distance)
METIS_VL+UV 1-5min 23.34Mb
(4VL+1UV) 77.8 - 389 hours
COMET
Mapping emission of sungrazing comets
(at any distance)
METIS_VL+UV 5-20min 23.34Mb
(4VL+1UV) 20 - 78 days
PROBE
Coordinated observations with Solar Probe Plus
(at any distance)
METIS_VL+UV 5-30min 19.85Mb
(4VL+1UV) 11 - 66 days
METIS observing mode characteristics
In the following definitions of METIS observing modes, the following naming conventions are
applied:
1. Size of raw frames (sets of samples representing the digitized signal given by each pixel of
the relevant detector):
1. VLD: (2048x2048) @ 14 bits (analogue mode; spatial scale per pixel: 10”);
2. UVD: (1024x1024) @ 14 bits (analogue mode; spatial scale per pixel: 20”).
The UVD can be operated in photon counting mode also; the photon counting mode
is mainly required to overcome possible degradation of the UV detector during a
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long-lived mission such as Solar Orbiter and to ensure the capability of observing
also in such a case.
2. DIT: Detector Integration Time, i.e., the time interval among which all detector pixels are
collecting photons. It represents the actual exposure of a single read out of the detector
(frame).
3. NDIT: number of detector integrations, i.e., the number of frames to be averaged in order to
obtain an “Acquisition”; this is the image obtained as result of averaging, pixel by pixel,
NDIT frames.
4. TACQ: Acquisition Time, the overall integration time corresponding to a single acquisition.
The values are set in order to get a limited number of events (cosmic rays) to be removed.
5. Cosmic ray (CR) removal: software procedure to clean the acquired images from spurious
signal given by the cosmic rays and SEP.
6. CME flag: software procedure to automatically trigger the rising of a CME event.
7. NACQ: number of acquisitions to be averaged in order to obtain an “Exposure”, i.e., the
image obtained as result of averaging, pixel by pixel, NACQ acquisitions.
8. TEXP: Exposure Time, i.e., the overall time corresponding to a single Exposure.
Single acquisitions will be averaged over the exposure time interval, unless a CME flag
occurs. Exposure times are determined considering the count rates estimated for the Sun at
its minimum of activity (Tables 6-8), as expected for a launch date of July 2017. The values
are set:
1. in order to get the best performance of the detector for the brightest features in the
expected image and, contemporarily,
2. get still a significant signal-to-noise ratio for the faintest features in the same
expected image;
3. on the basis of the typical life-times of the coronal structures considered;
4. taking into account the need of limiting the data volume of the final science
images downloaded to the ground.
9. NPOL: number of polarization (LCVR) angles used (cycled) during the scientific
measurement.
10. TOH: Overhead Time, given by the instrument in order to perform the scientific measurement
(i.e., synchronization, LCVR commanding, etc.).
11. TW: Waiting Time, defined by the operator within the observation timeline.
12. TCAD: Cadence Time, needed to get NPOL VL images or 1 UV image (within the same
scientific measurement, the time between two consecutive Exposures).
13. The counts per pixel in the final science images (to be downloaded by TM) are average
values calculated over the summed-up acquisitions and the pixels binned together. In this
case the depth of each final science image is 14 bits.
The acquisition time relationships are
VLD UVD
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TACQ = NDIT * DIT = DIT , NDIT = 1 TACQ = NDIT * DIT
TEXP = NACQ * TACQ TEXP = NACQ * TACQ
TCAD = NPOL * TEXP + TOH + TW TCAD = TEXP + TOH + TW
The characteristic requirements of each observing mode are given in the two following
sections METIS standard modes and METIS special modes.
METIS standard modes
source: Metis User Manual Iss3 (May 2016) - table 4.1.3.2
Specific instrument observing modes have been defined in order to address the scientific objectives
of the METIS investigation. In general, METIS observations consist of global maps of the coronal
emission in UV H I Ly-α and VL (580-640 nm range), obtained with different spatial resolution and
detector exposure time, depending on the science goal and the instantaneous field of view (FoV).
More in detail, the observing modes defined for METIS are:
WIND – Measurement of the electron density and the solar wind outward expansion velocity
Fast and slow solar wind streams are identified in the global maps according to the values of the
outflow velocity of the H component.
MAGTOP – Wind outflow velocity measurements and relationship with the magnetic topology
Maps of the outflow velocity of the H component along streamer/coronal hole interfaces, above
streamers cusps and inside streamers. Relationship with the coronal magnetic configuration.
GLOBAL – Global corona configuration/evolution measurements before, during and after CME
events These measurements provide the geometry of the neutral H and e- corona and its evolution in
time, giving information on the timing, mass content and overall dynamics of coronal mass
ejections. They are also crucial to measure the directionality of the plasma erupted from the Sun, in
order to infer its geo- effectiveness and predict the impact on the near-Earth environment. The
evolution related to the CME transit can be followed out to the orbit of the Solar Probe Plus
spacecraft (~9 R! at 0.8 AU).
LT-CONFIG – Long-term evolution of the coronal configuration
These measurements are used to monitor the evolution of the large-scale corona and, during out-of-
ecliptic observing windows, to determine the longitudinal distribution and evolution of the electron
density in the solar corona, as well as of the mass and energy flux carried away by the solar wind.
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METIS special modes
source: Science Performance Document Iss 4.0 (CDR version - March 2016) & Metis User Manual
Iss3 (May 2016) - table 4.1.3.2
Specific instrument observing modes have been defined in order to address the scientific objectives
of the METIS investigation. In general, METIS observations consist of global maps of the coronal
emission in UV H I Ly-α and VL (580-640 nm range), obtained with different spatial resolution and
detector exposure time, depending on the science goal and the instantaneous field of view (FoV).
The special observing modes defined for METIS are:
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FLUCTS and TBF – Brightness fluctuations spectra (best near perihelion) - acquisition sequence
details below
High spatial resolution and cadence time series of VL coronal brightness over a range of distances
covering from 1.6 R! out to 5 R!, to constrain the amplitude of the density fluctuations spectrum.
High frequency: TCAD=1 s (FLUCTS), =20 s (TBF).
The typical duration of this program is approximately 1 hour
CMEOBS – CME propagation, related driven shocks and filamentary structures, SEP accelerated
by CMEs Measurement of the electron density and outward expansion velocity gradient at high
spatial resolution and temporal cadence. This mode is activated by a proper CME event flag.
Such measurements can be used to identify the path of the shock front where particles can be
accelerated in the outer solar corona. Moreover, combined with radio observations, they can help to
distinguish flare- accelerated SEPs from those associated with CMEs.
COMET – Mapping the emission of Sungrazing comets
These measurements are used to monitor the evolution of the cometary emission along its trajectory
close to the Sun.
PROBE – Coordinated observations with Solar Probe Plus (SPP)
Characterizing the properties of the coronal regions crossed by the SPP spacecraft, during its transit
close to the Sun.
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FLUCTS + TBF mode: Coronal brightness fluctuations at high frequency: acquisition details
Source: Science Performance Report 4.0
The program relevant to the study of the high frequency brightness fluctuations spectra (observing
modes FLUCTS and TBF, see Table 12) will be carried out in about 1 hr, as shown in Fig. 2:
The program relevant to the study of the high frequency brightness fluctuations spectra (observing
modes FLUCTS and TBF, see Table 12) will be carried out in about 1 hr, as shown in Fig. 2:
- Run mode FLUCTS twice as follows:
• step 1: acquisition of ~60 consecutive raw VL images with DIT=1 s, NPOL=1, temporarily
stored in the instrument memory
• step 2: wait for on-board processing (masking, binning, compression, see Sect. 5.1.1) of the
raw images (TW ~10 s per image, in total ~600 s) and transfer of the reduced size images to
the S/C SSMM (280.2 Mb) to empty the instrument memory
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- Change from mode FLUCTS to mode TBF
- Run mode TBF:
• acquisition of ~120 VL images with DIT=20 s, NPOL=2, changing the polarization angle by
90° after 10 s during each acquisition; on-board processing of image i and transfer to the S/C
SSMM is performed contemporarily to the acquisition of image i+1 (data volume: 560.4
Mb)
The expected data volume produced by the instrument in 1 hr is equal to 1120.8 Mb.
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4.5 PHI
PHI is based on two telescopes, a High Resolution Telescope (HRT) and a Full Disk Telescope
(FDT) which define the targeted FOV. The polarimetric modulation is carried out by two
polarization modulation packages PMP, each one based on two temperature stabilized liquid crystal
variable retarders (LCVRs) located close to an intermediate
focus of each telescope. In order to allow for high-resolution polarimetric observations the HRT is
equipped with
an Image Stabilization System (ISS) based on a closed loop correlation tracker. This system is
designed to reduce the image motion on the science detector by moving a tip/tilt mirror. The
correction signals are obtained from cross-correlating images of the observation target obtained with
a fast active pixel sensor (APS) camera. The two telescopes sequentially – adjustable with a Feed
Select Mechanism (FSM) – feed the FG which is based on a tunable solid state LiNbO3 Fabry-Perot
etalon and an order sorting prefilter (OSPF). Tuning the transmission bands of the etalon is carried
out by changing the refractive index of the LiNbO3 wafer by means of adjusting the applied high
voltage which is produced by a High Voltage Power Supply (HVPS). Finally, the solar light will be
focused on two science detectors, each one based on 2kx2k APS sensors. In order to reduce the
power and radiation entry into the instrument both telescopes are equipped with Heat Rejection
Entrance Windows (HREWs) mounted at the heat shield of the S/C. For calibration purposes both
the HRT and the FDT as well as the ISS are equipped with re-focusing mechanisms (RFMs).
Instrument control, data pre-processing (calibration and polarimetric demodulation) as well as
radiative transfer equation (RTE) inversion will be carried out in a Digital Processing Unit (DPU)
which is based of two re-configurable and one fixed field programmable gate arrays FPGA. The
DPU is equipped with a large mass memory (4TBytes) which allows for storing calibration data and
science data which cannot be processed in real time.
4.5.1 PHI observables
PHI is an imaging spectro-polarimeter in order to measure the photospheric continuum intensity, Ic,
the full magnetic field vector, B = [|B|, γ, φ] and the line of sight (LOS) flow velocity, v_LOS, in
each point of dedicated field of view (FOV) on the solar disk.
4.5.2 PHI modes
TM figures
Allocated TM 20.5 Kbits/s
Download capacity per orbit 6.642 GB = 53.136Gbits
Subtelescopes
PHI has 2 telescopes but these use the same detector so at each moment in time only 1 of the
telescopes can observe and generate data!
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wavelength detector size FOV Vol/raw
dataset
TM/processed
image
FDT 617.3nm line (measured at 5 wavelengths +
continuum)
2048 px * 2048
px 2ºx2º 302 MB 12 MB
HRT 617.3nm line (measured at 5 wavelengths +
continuum)
2048 px * 2048
px 16.8'x16.8' 302 MB 12 MB
Observational modes
For PHI, there is only 1 data acquisition mode for each telescope. However, the data processing
onboard differs in the following 'observing modes':
Out of date!
To be updated! PHI observational modes have changed recently
Observing mode (=processing
options) Use Case Telescope Cadence Data Rate Max time / orbit
(s) (Kbits/s) (restricted by TM
alloc.)
PHI science mode 0
nominal mode FDT/HRT 60 1607,7824 9 h
e.g. magnetoconvection
PHI science mode 1 local helioseismology FDT/HRT 60 21,81 28 days
Low resol, high cadence target: 100days
PHI science mode 2 magnetic field evolution FDT/HRT 300-3600 3,79-
52,43 ~24 days - orbit
med resol, med cadence
PHI science mode 3 subfield FDT/HRT 60-3600
(network) 1-262 2 days - unlimited
high resol, high/med cadence 300-3600 (AR)
PHI science mode 4 photospheric context FDT/HRT 300 182-210 2-3 days
PHI science mode 5
global helioseismic/synoptic observations
FDT 60-3600 1,5-5,5 100 days or more
PHI science mode 6 daily context FDT only 1-4 per day 0,24-1 unlimited
( SYNOPTIC MODE)
PHI science mode 7 burst/flare mode, triggered by STIX HRT 60 350 ~40 hours
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Science modes specific for Helioseismology
TBW
Low Latency or Precursor programs
Observing mode Use Case Telescope Cadence Data Volume Max time / orbit
PHI Full-disk Precursors Full sun continuum + magnetogram FDT or HRT 1set/day
3,2 Mibits/set
= 0,2 MB / set
PHI Calibration set full calibration set FDT and HRT n/a
8,63 Mibits/set
= 1,1 MB / set
Special datasets
Observing mode Use Case Telescope Cadence Data Volume Max time / orbit
Raw data downlink full, raw image set FDT or HRT n/a
2304Mibits / set
= 302 MB / set
1 set = 5% orbital allocation
PHI Full-disk Precursors
PHI can take full disk precursor images in the days before a RS window. This 'mode' is configured
to generate those, once a day.
• Parameters
TBD
• Resource usage
Data
processi
ng
modes
Use case
Dataset
after
inversi
on
Rebinni
ng
digital depth/phys. quantity
(bit/px)
Min
Caden
ce
Max
Caden
ce
Compressi
on
Data
Vol/set
Min
TM
Rate
Max
TM
Rate
Rate Rate
HRT/FD
T Mibits
or
subfield
I_
c
B_LO
S
v_LO
S
gam
ma
ph
i (s) (s) Rate (*) Mibits
(kibits/
s)
(kibits/
s)
(kibits/
s)
(kbp
s)
Full-disk
Precurso
rs
full sun
continuum 640,00 2 8 8 0 0 0 86400 86400 10 1,60 0,037 0,037 0,04
+magnetogr
am 640,00 2 8 8 0 0 0 86400 86400 10 1,60
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FDT or HRT
1set/day
=0.2MB/set
PHI science mode 0
PHI science mode 0 is the nominal mode of PHI. However, it generates that much TM that it cannot
all be downloaded. Only 9 hours of Mode 0 fit in PHI's orbital allocation.
This mode can be run either with the Full Disk Telescope (FDT) or with the High Resolution
Telescope (HRT), each projecting onto the same camera plane.
• Parameters
TBC if any
• Resource usage
Data processing
modes
Use case Dataset after
inversion
Rebinning digital depth/phys. quantity (bit/px)
Min Cadence
Max Cadence
Compression Data Vol/set
Min TM Rate
Max TM
Rate
Rate Rate
HRT/FDT Mibits or
subfield I_c B_LOS v_LOS gamma phi (s) (s) Rate (*) Mibits (kibits/s) (kibits/s) (kibits/s) (kbps)
PHI
science
mode 0
nominal mode 640,00 1 10 10 10 8 8 60 60 2 92,00 1570,1 1570,1 1570,1 1607,7824
e.g. magnetoconvection
PHI science mode 1
PHI science mode 1 is used for local helioseismology. It has lower resolution than PHI science
mode 0 (rebinned or sub fielded to 512x512) but still high cadence of 1 processed dataset per
minute (consists of 5 physical value images processed from 24 raw).
PHI mode 1 is designed to run both on HRT and FDT.
• Parameters
TBC if any
• Resource usage
Data
processing
modes
Use case
Dataset
after
inversion
Rebinning digital depth/phys. quantity
(bit/px)
Min
Cadence
Max
Cadence Compression
Data
Vol/set
Min TM
Rate
Max TM
Rate Rate Rate
HRT/FDT Mibits or
subfield I_c B_LOS v_LOS gamma phi (s) (s) Rate (*) Mibits (kibits/s) (kibits/s) (kibits/s) (kbps)
PHI science
mode 1
local
helioseismology 640,00 4 0 0 10 0 0 60 60 2 1,25 21,3 21,3 21,3 21,81
Low resol, high cadence
(512x512)
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PHI science mode 2
PHI science mode 2 has lower resolution than PHI science mode 0 (rebinned or sub-fielded to
1024x1024) and a medium cadence that can vary from 1 processed dataset per 5 minutes to 1
dataset per hour.
It can be used with both FDT and HRT telescopes.
• Parameters
Cadence (TBC)
• Resource usage
Data
processing
modes
Use case
Dataset
after
inversion
Rebinning digital depth/phys. quantity
(bit/px)
Min
Cadence
Max
Cadence Compression
Data
Vol/set
Min TM
Rate
Max
TM
Rate
Rate Rate
HRT/FDT Mibits or subfield I_c B_LOS v_LOS gamma phi (s) (s) Rate (*) Mibits (kibits/s) (kibits/s) (kibits/s) (kbps)
PHI
science
mode 2
magnetic
field
evolution
640,00 2 10 10 10 0 0 3600 300 2 15,00 4,3 51,2 1->51.2 52,43
(max)
med
resol,
med
cadence
640,00 2 0 10 0 8 8 3600 300 2 13,00 3,7 44,4 3,79
(min)
(1024x1024)
PHI science mode 3
PHI science mode 3 is a subfield mode with high resolution (full resolution in subfield of
1024x1024 or 512x512) and a medium to high cadence that can vary from 1 processed dataset per
minute (for network studies and ARs) to 1 dataset per hour.
It can be used with both FDT and HRT telescopes.
• Parameters
Subfield size, cadence, physical quantities (TBC)
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• Resource usage
Data processing
modes
Use case Dataset after
inversion
Rebinning digital depth/phys. quantity
(bit/px)
Min
Cadence
Max
Cadence Compression
Data
Vol/set
Min TM
Rate
Max TM
Rate Rate Rate
HRT/FDT Mibits or
subfield I_c B_LOS v_LOS gamma phi (s) (s) Rate (*) Mibits (kibits/s) (kibits/s) (kibits/s) (kbps)
network/AR
PHI science
mode 3 subfield 640,00 2 10 10 10 0 0 3600 60 2 15,00 4,3 256,0 1->256
262,14
(max)
high resol, high/med
cadence
640,00 4 10 10 10 0 0 3600 300 2 3,75 1,1 12,8 1,09 (min)
(subfield) AR
extra
possibility: 640,00 0 10 0 8 8
PHI science mode 4
PHI science mode 4 is used for photospheric context. It outputs a subset of physical quantities (3
output images out of 5) at a cadence of 5 minutes and maximal resolution.
It can be used both with HRT and FDT.
• Parameters
physical quantities (TBC)
• Resource usage
Data
processing
modes
Use case
Dataset
after
inversion
Rebinning digital depth/phys. quantity
(bit/px)
Min
Cadence
Max
Cadence Compression
Data
Vol/set
Min TM
Rate
Max TM
Rate Rate Rate
HRT/FDT Mibits or
subfield I_c B_LOS v_LOS gamma phi (s) (s) Rate (*) Mibits (kibits/s) (kibits/s) (kibits/s) (kbps)
upper quantities for dynamical
studies
PHI science
mode 4
photospheric
context 640,00 1 10 10 10 0 0 300 300 2 60,00 204,8 204,8 204,8 209,72
both FDT &
HRT 640,00 1 0 10 0 8 8 300 300 2 52,00 177,5 177,5 181,75
lower quantities for vector
magnetometry
PHI science mode 5
PHI science mode 5 is used for global helioseismology and synoptic observations. It outputs only 2
physical quantities (continuum intensity and LOS velocity) at a varying cadence highly reduced
resolution.
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Designed for Full Disk Telescope only.
Parameters
cadence, resolution (rebinning)
Resource usage
Data
processing
modes
Use case
Dataset
after
inversion
Rebinning digital depth/phys. quantity
(bit/px)
Min
Cadence
Max
Cadence Compression
Data
Vol/set
Min TM
Rate
Max TM
Rate Rate Rate
FDT Mibits or
subfield I_c B_LOS v_LOS gamma phi (s) (s) Rate (*) Mibits (kibits/s) (kibits/s) (kibits/s) (kbps)
(256x256)
PHI science
mode 5 global
helioseismic/ synoptic
observations
640,00 8 10 0 10 0 0 3600 60 2 0,63 0,2 10,7 1.4-5.4 5,53
(max)
640,00 16 10 0 10 0 0 3600 60 2 0,16 0,0 2,7 1,43
(min)
(128x128) for I_c for
v_LOS
PHI science mode 6
PHI science mode 6 is PHI's default synoptic mode for daily context. It is configured to generate 1
to 4 times per day a set of all 5 or only 2 physical quantities (continuum intensity and LOS velocity)
at slightly lower resolution.
Designed for full disk telescope only
Parameters
cadence, physical quantities
Resource usage
Data
processing
modes
Use
case
Dataset
after
inversion
Rebinning digital depth/phys. quantity
(bit/px)
Min
Cadence
Max
Cadence Compression
Data
Vol/set
Min TM
Rate
Max TM
Rate Rate Rate
FDT Mibits or subfield I_
c B_LOS v_LOS gamma
ph
i (s) (s) Rate (*) Mibits (kibits/s) (kibits/s) (kibits/s) (kbps)
PHI science
mode 6
daily
context 640,00 2 4 4 4 4 4 86400 21600 1 20,00 0,24 0,95 0.24-0.95
0,97
(max)
( SYNOPTIC
MODE)
only
FDT
(subfield or rebin
1024x1024)
1/day 4/day ? 0,24
(min)
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OR
(see RS WG)
640,00 2 4 4 0 0 0 21600 1 8,00 0,4
PHI science mode 7
PHI science mode 7 is defined to be PHI's response to a STIX flare trigger. It is a burst or flare
mode that generates fast cadence continuum images in order to observe white-light flares.
It is designed for HRT only.
Parameters
TBD
Resource usage
Data
processing
modes
Use case
Dataset
after
inversion
Rebinning digital depth/phys. quantity (bit/px) Min
Cadence
Max
Cadence Compression
Data
Vol/set
Min TM
Rate
Max TM
Rate Rate Rate
HRT Mibits or subfield I_c B_LOS v_LOS gamma phi (s) (s) Rate (*) Mibits (kibits/s) (kibits/s) (kibits/s) (kbps)
burst/flare mode,
triggered by STIX
PHI science
mode 7 640,00 1 10 0 0 0 0 60 60 2 20,00 341,33 341,33 TBD 349,53
TBD, most likely fast
cadence continuum images
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4.6 RPW
RPW is a plasma/radio wave receiver system, including high sensitivity electric and
magnetic sensors. Since the receiver system covers a very wide frequency range (near-DC to 16
MHz for electric, and 0.1 Hz to 500 kHz for magnetic), different kinds of sensors are used for the
measurements.
The RPW instrument comprises:
A Thermal Noise and High Frequency receiver (TNR-HFR) for electron measurements at the local
plasma frequency and remote detection of radio emissions. TNR-HFR will provide electric power
spectral densities from 4 kHz up to 16 MHz and magnetic power spectral densities from 10 kHz up
to 500 kHz.
A Time Domain Sampler (TDS) for waveform capture up to 500 kSPS.
A Low Frequency Receiver (LFR), covering both in-situ electric and magnetic measurements from
DC to about 10 kHz. LFR will provide both waveform and power spectral densities in this
frequency range.
These three sub-systems are connected to two different sensor units: An electric antenna unit (ANT)
and a magnetic search coil unit (SCM), both of which will be optimized to perform correctly for
near-DC as well as high frequency measurements. In particular, the antenna sensor design will be
optimized to satisfy the goal of measuring both DC/low frequency electric fields and higher
frequency radio and thermal noise emissions. A biasing unit (BIAS) will allow DC electric
measurements. The three TDS, LFR and TNR-HFR sub-systems will have a common Data
Processing Unit (DPU) that will handle commands, data and communication with S/C.
ANT: Each ANT monopole serves as a simple voltage sensor. At low frequencies, an antenna is
coupled to the local plasma potential through a photoelectron sheath. Successful measurement of
DC/low frequency plasma electric fields requires that the antenna be biased (as described below). At
sufficiently high (radio) frequencies, an antenna behaves as if in a vacuum.
Preamplifiers: Each monopole is connected to the inputs of both: A low frequency and high
frequency preamplifiers.
The LF preamplifiers will measure voltage and provide bias current, using a high input impedance
follower with a bootstrapped bias resistor and voltage source. The input stage needs to handle a high
source impedance of R ≈ 50 Mohms, and C ≈ 40 pF. Thus, it must have a low leakage current, <
10pA, for the whole mission, and a low input capacitance, < 4pF, to minimize attenuation of the
input signal. This can only be achieved by proper bootstrapping of the current generator, and by
using FET operational amplifiers.
The HF preamplifiers will provide a low noise and flat frequency response from 100Hz to more than
16 MHz. As well as the LF preamplifier, theirs input impedance will be as high as possible.
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BIAS Unit: The BIAS drives a constant current to the electric antennas allowing reliable DC/LF
electric field and satellite potential measurements. The floating potential V of a conductive surface
in plasma (either an individual probe or the entire spacecraft) is determined by current balance
between the out flowing photoelectron current and the inflowing plasma electron current. In the
expected solar wind plasma, a conductive surface will charge up positively and the value of the
floating potential will depend on the density and (to a lesser extent) the temperature of the
surrounding plasma. In contrast to the spacecraft potential, the antenna potential should be anchored
tightly to the local plasma potential, and should not depend on the plasma parameters and also
should be affected as little as possible by the nearby satellite. To achieve this, we draw a constant
bias current from the antennas to the spacecraft. If the bias current is suitably chosen, then the
antennas float close to the local plasma potential and their potential is not a function of the plasma
parameters. Then measurements of the potential difference between two opposite antennas which
are anchored to the local plasma potential gives reliable measurement of DC/LF electric field
component along that direction. In addition, the potential difference between antennas and
spacecraft give an estimate of the spacecraft potential that can be used to determine the plasma
density.
SCM: The Search Coil Magnetometer SCM is an inductive magnetic sensor. It is made of a core in
a high permeability material (ferrite or perm-alloy) on which are wound a main coil with several
thousand turns and a secondary coil with a few turns. The secondary coil is used to create a flux
feedback in order to have a flat frequency response on a bandwidth centered on the resonance
frequency of the main coil. The induced voltage is raised to a proper level by a preamplifier to allow
its transportation to the analyze system.
TNR-HFR: The proposed TNR-HFR instrument is a contribution to the RPW-E experiment
consisting of a double channel radio and plasma wave spectrometer. The TNR (Thermal Noise
Receiver) is a direct conversion receiver, providing quasi-instantaneous spectra, for the electrostatic
thermal noise and/or magnetic field, while the HFR (High Frequency Receiver) is a sweeping
receiver, for the survey of high frequency radio emissions. Its analogue front end is interfaced with
three sensors, two E-field inputs (4 kHz - 16MHz) and one B-field component (10kHz to 500kHz).
TDS: The Time Domain Sampler module (TDS) will provide in-situ waveform measurements of
plasma waves around local electron plasma frequency, notably Langmuir waves found in the source
region of type II and type III solar bursts and the associated electromagnetic waves. The TDS will
perform digitization of the electric and magnetic field waveforms in the frequency range from 100
Hz to 250 kHz, their pre-processing and selection of potentially interesting events to be stored in
internal memory and later transmitted to the ground. It is known from previous observations, which
impacts of dust particles on the spacecraft body show up in the electric field data as short large
amplitude spikes of a characteristic shape. An algorithm for detection of these spikes will be
implemented in the on-board software in order to collect statistics of dust impacts as a function of
time.
LFR: The Low Frequency Receiver (LFR) performs onboard digital processing of the electric and
magnetic field data (2E, 3B, 3V) and covers a frequency range from a fraction of a Hertz to 10 kHz.
It is designed to analyze the in-situ measurements of the low frequency (f <fce) electro-magnetic
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waves in the solar wind and in the extended corona. Given the limitations in the telemetry, it is
necessary to implement specific techniques to take the maximum advantage of the data. The LFR is
tailored to optimize the scientific return of the data. The LFR design gives the possibility of mixing
different types of output data, from low-level processed data (waveform data) to high-level
processed data (averaged Hermitian spectral matrices and their derived parameters), with various
data rate possibilities (continuous or cyclic transmission, adaptable frequency bandwidth as well as
adaptable frequency and time resolutions). The scientific added value stems from the choice of the
most relevant combination of the different data to be transmitted.
4.6.1 RPW observables
4.6.2 RPW modes
RPW Normal mode
Analyzer Product Data Rate
LFR Waveform 1391
BP/ASM 1469
TDS
Waveform 851
Histograms 65
Statistics 25
TNR-HFR
TNR Products 532
HFR Products 63
Spectral Power 3
RPW burst mode
Analyzer Product Data Rate (bps)
LFR Waveform 12528
BP 7056
TDS
Waveform 864
Histograms 43
Statistics 29
TNR-HFR
TNR Products 5328
HFR Products 144
Spectral Power 29
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RPW SBM1
This is a triggered mode involving the dump of a rolling buffer to the SSMM on detection of an
interplanetary shock. The detailed data rates calculated here don't completely match those in the
Plasson presentation used to calculate the summary rates, although the discrepancy isn't as big as for
SBM2. This is likely based on assumptions about trigger duration used in the volume estimates in
the Plasson presentation. For both SBM1 and SBM2 the summary rates are higher so will be used
wherever possible.
Analyzer Product Data Rate (bps)
LFR Waveform 201046
BP 8861.54
TDS Waveform 33784.0
RPW SBM2
SBM2 is a triggered mode used on detection of an in situ type III radio burst. These are expected to
occur roughly once every 40 days. Note that the detailed data rates listed here (source: TM Report,
CDR issue) don't match up with the figures from the Plasson presentation used to calculate the
summary rates. This could be rounding errors or different assumptions about duration of triggers
(this is configurable).
Analyzer Product Type Data Rate (bps)
LFR Waveform 12293
BP 6849
TDS Waveform 87804
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4.7 SoloHI
4.7.1 SoloHI observables
SoloHI takes visible light images of the extended corona/solar wind. The signal is made up of 4
sources: (1) photospheric light scattered by the free electrons expelled by the sun, (2) photospheric
light scattered by dust including comets and asteroids, (3) stellar, galactic and planetary sources, and
(4) instrumental stray light. Light is converted into a digital number in each pixel of the detector. An
array of pixels comprises an image. Images are compressed onboard and together with an image
header which describes the parameters of the image form a file. Each file is converted into packets
which are sent down the science channel to the SSMM. The degree of compression, including
whether there is any compression, is selected as part of the instrument schedule.
4.7.2 SoloHI modes
TM figures
Allocated TM 20.5 Kbits/s
Download capacity per orbit 6.642 GB = 53.136Gbits
Subtelescopes/units that can be commanded independently
SoloHI consists of 1 white-light telescope with wide FOV. The image is captured on a mosaic of
four 2048x1920 APS detectors that are read out independently. This gives flexibility for image
operations: independent exposures, cadences, etc.). Data can be read out either from the whole
detector or from selected subfields.
Typically, SoloHI images will have one of 3 typical FOVs defined below. These are used in the
observing programs/modes currently defined but could be changed in-flight if necessary.
FOV split (radial x
transverse) Downlinked pixels Typical cadence Comments
SoloHI Full frame image
(40ºx40º)
5º to 45º x 40º
***
1960 px * 1960 px
(incl. 2x2 bin)
24-36 min (inner
FOV)
30-72 min (outer FOV)
***split in 2 or 3 readout frames depending on solar distance,
e.g. 5º to 25º x 40º (inner) + 25º to 45º x 40º (outer) at
perihelion,
each with different cadence. Details in table below.
SoloHI inner FOV
subframe images
(3 images of 1.88ºx5º)
5.8º to 7.68º x 5º 13.5º to 15.38º x 5º
18.5º to 20.38º x 5º
192 px * 512 px (not binned)
96 px * 256 px
(binned 2x2)
96 px * 256 px (binned 2x2)
18-36 sec 36-72 sec
1.5-3.0 min
Subframe images typically only used at and near
perihelion (up to 0.36AU)
Radial Swath subframe
image
(40ºx5º)
5º to 25º x 5º
25º to 45º x 5º
1960 px * 256 px
(incl. 2x2 bin)
6 min (inner
FOV)
12 min (outer FOV)
Radial swath images typically only used at and near
perihelion (up to 0.36AU)
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Main science programs
source: SoloHI UM - SSD-DOC-SOLOHI-013 Rev. B draft 5
The SoloHI baseline observing program will be defined to repeat for each occurrence of the same
unique orbit (i.e. each orbit of the trajectory within the same resonance with Venus). Therefore, a
SoloHI baseline observing program will be defined for each orbit phase in NMP/EMP and will be
executed for all orbits within that resonance.
Example of such an orbit plan is below:
Observing modes - Example plan Use case #images /
day
Science data volume /
day
SoloHI data
rate
Observing duration /
orbit
(Gbits) estimate (Kbits/s) (days) example
Perihelion programs: 0.28-0.29 AU
SoloHI Solar Wind Turbulence @perih 1296 2.22 26.5 3
SoloHI Shock Formation @perih 468 2.54 30.3 3
Near-Perihelion programs: 0.29-0.36 AU
SoloHI Near-perihelion Synoptic Program 348 1.69 20.3 5
SoloHI Solar Wind Turbulence Near-
perihelion 750 1.94 23.2 2
SoloHI Shock Formation Near-perihelion 516 2.45 29.3 2
Far-Perihelion programs: 0.36-0.42 AU
SoloHI Far-Perihelion Synoptic Program 132 1.64 19.7 7
Southern Out-of-ecliptic programs:
0.5-0.7 AU
SoloHI Southern Synoptic Program 104 0.84 10.3 8
Dependent on the trajectory
Examples of more-detailed observing program for 1 type of orbit during the mission (source:
04_130904_SoloHI_CDR_ObsProg.ppt):
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Based on table above:
• a typical perihelion programme would produce ~25kbps (during 4 days), -> see modelled
observations HI_SHOCK_PER (DATARATE=30300 [bits/sec]), HI_TURB_PER
(DATARATE=26500 [bits/sec]),
• near-perihelion SoloHI would produce ~20kbps (during 8 days) , -> see modelled
observation HI_SYN_PER (DATARATE=20300 [bits/sec])
• ~18.5kbps even further out (during 12 days) and -> see modelled observations
HI_SYN_NEAR (DATARATE=19700 [bits/sec])
• in the far-out RS window, a data rate around 10 kbps would be reached. -> see modelled
observations HI_SYN_FAR (DATARATE=10300 [bits/sec])
(see also SoloHI concept study report Dec 2011)
Update needed:
How to organize SoloHI observations in coordination with the other instruments, i.e. does SoloHI
have 'observing modes' to choose from for each solar distance?, is still to be discussed in more
detail. Also, while the schema above may be optimal from a science perspective, the varying
downlink rate & SSMM storage limits may impose limitations on when which data rate can be used.
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4.8 SPICE
4.8.1 SPICE observables
SPICE is a spectrometer that will measure plasma density and temperature, flow velocities, presence
of plasma turbulence and composition of source region plasma on the solar disk.
4.8.2 SPICE modes
TM figures
Allocated TM 17.5 Kbits/s
Download capacity per orbit 5.670 GB = 45.36Gbits (over 30 days)
Subtelescopes/detectors
Source: SPICE User Manual Iss 5.0 (Nov 2014)
SPICE consists of only 1 spectral telescope with slit and scanning mirror. The 2D image (1
wavelength dimension, 1 spatial dimension along the slit) is projected on 2 detectors, 1 covering
the long wavelengths, the other covering the short wavelengths.
wavelength range detector (readout) size FOV Volume/Detector readout
SW array (approx.) 70.1nm - 79.3nm
968px @0.0095nm/pix
x 800 px (along slit)
11' (14' for 30" slit) x 16' 1.55MB
LW array (approx.) 97.1nm - 105.1nm
968px @0.0083nm/pix
x 800 px (along slit)
11' (14' for 30" slit) x 16' 1.55MB
Observational modes
Source: SPICE User Manual Iss 5.0 (Nov 2014)
SPICE's primary mode of observations consist of rapid on-disk scans that characterize plasma
dynamics (SPICE Dynamics mode), alternated with slower composition scans that map the source
regions of solar wind streams (SPICE Composition Mapping). However, SPICE also has alternative
modes or 'studies' listed below. These studies will be preset on board and can be called via TC,
providing some of the study parameters (see mode pages).
Observing mode Use Case/Target Line List #
repeats
Total
Duration
Data
Rate
Max time /
orbit
(min) (Kbits/s)
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SPICE Composition
Mapping
= NOMINAL MODE 1
Centre, Poles, limb,
AR
15 lines
(2 profiles, 13
intensities)
1 180 0.45
SPICE Dynamics mode
=NOMINAL MODE 2
Centre, AR, CH
H I, C III, O VI,
NeVIII profiles
+ 6 intensities
10 110 20.46
SPICE Spectral Atlas
Sun center, limb,
AR, CH Full spectrum 2 22 40.34
SPICE Limb mode
Low corona above
limb
C III, O VI, Ne VIII
profiles
+ 3 intensities
1 240 2.5
SPICE CME Watch
AR, prominence,
limb
5 spectral profiles, 10
intensities 30 22.66 hr 4.46
SPICE 30"-wide Movie
(sit&stare) center, AR 1-2 line profiles 1 10 34.44
SPICE 90"-wide Movie center, AR 1-2 line profiles 40 16 20.28
SPICE Waves mode
(sit&stare) QS, CH, AR
C III, O VI, NeVIII
profiles 5 300 50.72
SPICE Two-Exposure
mode
combi of 6 bright and
faint lines 5 300 2.86
Low Latency or Precursor programs
SPICE plans to run a mini-study before the start of each new study, that is configured exactly like
the science study but has lower resolution/cadence/slit width/... and produces ~0.1MB.
Special datasets
Observing mode Use Case/Target Line List Duration/dataset Data Rate Max time / orbit
SPICE Full Raster Scan few strong lines 32 15.1
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Selection of SPICE spectral lines
Source: SPICE User Manual Iss 5.0 (Nov 2014)
Ion Wavelength (Å) Log T (K) FIP (eV) M/q
H I 1025 4.0 13.6 --
C II 1036 4.3 11.3 12.0
C III 977 4.5 11.3 6.0
O IV 787.7 5.2 13.6 5.3
O V 760 5.4 13.6 4.0
O VI 1032 5.5 13.6 3.2
S V 1037 5.5 13.6 3.2
Ne VI 786.5 5.2 10.36 8.0
Ne VII 1005 5.5 21.6 4.0
Ne VIII 973 5.6 21.6 3.3
Ne VIII 770 5.8 21.6 2.8
Mg VIII 772 5.9 7.7 3.4
Mg IX 706 6.0 7.7 3.0
Mg XI 997 6.2 7.7 2.4
Si VII 1049 5.6 8.1 4.8
Si XII 521 (2nd) 6.5 8.1 2.6
Fe X 1028 6.0 7.9 6.2
Fe XVIII 975 6.9 7.9 3.3
Fe XX 721 7.0 7.9 2.9
Auxiliary lines:
Ion Wavelength (Å) Log T (K) FIP (eV) M/q
Ne VIII 780 5.8 21.6 2.8
Si XII 499 (2nd) 6.5 8.1 2.6
SPICE Composition Mapping
This is the SPICE nominal observation mode I
Target: Centre, Poles, E-W, N-S, limb, AR
Line List: 15 lines (2 profiles, 13 intensities)
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Slit: 4"
Data rate calculation
Parameters are highlighted in blue.
Observing
mode
#Line
Profiles
#Px/Profi
le
#
Intensiti
es
Slit
(")
Exp.ti
me
#
Mirror
Pos
Step size Tilt
time FOV (") Data cube (px)
Durati
on
Repeat
s
Cal
c.
Dat
a
Vol
Total Data
Vol
Total
Durati
on
Data
rate
= study (max=32)
2/4/6/30
(s) =#exposures
(arcsecs) (secs) X Y X Y Z (min) (MB)
Compressed (MB)
(min) (Kbits/s)
Compositio
n Mapping 2 32 13 4" 180 60 4" 0.14
240
"
660
" 60 800 90 180 1 8.64 0.606 180.1 0.45
NeVIII,
MgIX 4' 11'
sca
n
dir
alon
g
slit
spectr
al dim
Spectral-line performance
Line-SNR and spatial resolution combinations:
line-SNR>10 at 4"x2" in C-III, O-VI, Ne-VII, and Mg-IX-AR
Selection of SPICE spectral lines
Source: SPICE User Manual Iss 5.0 (Nov 2014)
Ion Wavelength (Å) Log T (K) FIP (eV) M/q
H I 1025 4.0 13.6 --
C II 1036 4.3 11.3 12.0
C III 977 4.5 11.3 6.0
O IV 787.7 5.2 13.6 5.3
O V 760 5.4 13.6 4.0
O VI 1032 5.5 13.6 3.2
S V 1037 5.5 13.6 3.2
Ne VI 786.5 5.2 10.36 8.0
Ne VII 1005 5.5 21.6 4.0
Ne VIII 973 5.6 21.6 3.3
Ne VIII 770 5.8 21.6 2.8
Mg VIII 772 5.9 7.7 3.4
Mg IX 706 6.0 7.7 3.0
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Ion Wavelength (Å) Log T (K) FIP (eV) M/q
Mg XI 997 6.2 7.7 2.4
Si VII 1049 5.6 8.1 4.8
Si XII 521 (2nd) 6.5 8.1 2.6
Fe X 1028 6.0 7.9 6.2
Fe XVIII 975 6.9 7.9 3.3
Fe XX 721 7.0 7.9 2.9
Auxiliary lines:
Ion Wavelength (Å) Log T (K) FIP (eV) M/q
Ne VIII 780 5.8 21.6 2.8
Si XII 499 (2nd) 6.5 8.1 2.6
SPICE Dynamics
This is the SPICE nominal observation mode II
Target: Centre, Poles, E-W, N-S, limb, AR
Line List: 15 lines (2 profiles, 13 intensities)
Slit: 2" Step size: 2"
Data rate calculation
Observing mode
#Line Profiles
#Px/Profile
Intensities
Slit (")
Exp.time
# Mirror Pos
Step size
Tilt time
FOV (")
Data cube (px) Duration
Repeats
Calc.
Data Vol
Total Data Vol
Total
Duration
Data rate
= study (max=3
2)
2/4/6/3
0 (s)
=#exposur
es
(arcsec
s)
(secs
) X Y X Y Z (min) (MB)
Compress
ed (MB) (min)
(kbits/
s)
Dynamic
s 4 32 6 2 5 120 2 0.07
24
0
84
0 120 800 140 10.26 10
268.8
0 15.756 101.4 20.46
4 14
scan
direction
alon
g slit
spectral
dimension
for 10
repeats
for 10
repeats
Spectral-line performance
Line-SNR and spatial resolution combinations:
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line-SNR > 10 at 2"x2" in C-III, O-VI, Ne-VIII (in ARs),
and at 5"x5" in Ne-VIII (in CHs and QS)
Selection of SPICE spectral lines
Source: SPICE User Manual Iss 5.0 (Nov 2014)
Ion Wavelength (Å) Log T (K) FIP (eV) M/q
H I 1025 4.0 13.6 --
C II 1036 4.3 11.3 12.0
C III 977 4.5 11.3 6.0
O IV 787.7 5.2 13.6 5.3
O V 760 5.4 13.6 4.0
O VI 1032 5.5 13.6 3.2
S V 1037 5.5 13.6 3.2
Ne VI 786.5 5.2 10.36 8.0
Ne VII 1005 5.5 21.6 4.0
Ne VIII 973 5.6 21.6 3.3
Ne VIII 770 5.8 21.6 2.8
Mg VIII 772 5.9 7.7 3.4
Mg IX 706 6.0 7.7 3.0
Mg XI 997 6.2 7.7 2.4
Si VII 1049 5.6 8.1 4.8
Si XII 521 (2nd) 6.5 8.1 2.6
Fe X 1028 6.0 7.9 6.2
Fe XVIII 975 6.9 7.9 3.3
Fe XX 721 7.0 7.9 2.9
Auxiliary lines:
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Ion Wavelength (Å) Log T (K) FIP (eV) M/q
Ne VIII 780 5.8 21.6 2.8
Si XII 499 (2nd) 6.5 8.1 2.6
SPICE Spectral Atlas
Observation configuration and data rate calculation
Observing
mode Target
Line
List
#Line
Profil
es
#Px/Pr
ofile
#Int
ens
Slit
(")
Exp.ti
me
#
Mirror
Pos
Step
size
Tilt
time
FOV
(")
Data cube
(px)
Durat
ion
Repe
ats
Cal
c.
Dat
a
Vol
Total
Compre
ssed
Data
Vol
Total
Durat
ion
Data
rate
Comm
ents
= study (max=32)
2/4/6/30
(s) =#exposures
(arcsecs)
(secs)
X Y X Y Z (min) (MB)
(MB) (min) (kbits/s)
Spectral
Atlas (A
tomic
Physics)
Sun
centre,
limb,AR
,CH
Full
spectr
um
32 64 0 4 60 10 4 5.1 40 840 10 800 2048 11.1 2 65.
54 6.666 21.5 40.34
Excepti
onal
campai
gn, but
also
used for
calibration
5s
read
out
+0.1
s
step
sc
an
dir
alo
ng
slit
sc
an
dir
alo
ng
slit
=M
AX
30'/2r
ep
Spectral-line performance
Line-SNR and spatial resolution combinations:
line-SNR>10 at 4"x2" in C-III, O-VI, Ne-VII, and Mg-IX-AR
SPICE Limb mode
Observation configuration and data rate calculation
Observ
ing
mode
Tar
get
Line
List
#Line
Profil
es
#Px/Pr
ofile
#
Intensi
ties
Slit
(")
Exp.ti
me
#
Mirror
Pos
Step
size
Tilt
tim
e
FOV
(") Data cube (px)
Durat
ion
Repe
ats
Cal
c.
Dat
a
Vol
Total
Data
Vol
Total
Durat
ion
Data
rate
Comm
ents
= study (max=
32)
2/4/6/
30 (s)
=#expos
ures
(arcse
cs)
(sec
s) X Y X Y Z (min)
(M
B)
Compre
ssed
(MB)
(min) (kbit
s/s)
Limb
Low
coro
na
abov
e
limb
C III, O VI, Ne
VIII
profiles
+
3intensi
ties
3 32 3 4 60 240 4 0.1
4
96
0
84
0 240 800 102 240.6 1
39.
17 4.545 240.6 2.52
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16
'
14
'
scan
direction
alo
ng slit
spectral
dimension
Spectral-line performance
Line-SNR and spatial resolution combinations:
line-SNR >10 at 4"x2" in C-III, O-VI, Ne-VIII and Mg-IX (in ARs)
SPICE CME Watch
Observation configuration and data rate calculation
Obser
ving
mode
Target Line
List
#Line
Profil
es
#Px/Pr
ofile
#
Intensi
ties
Slit
(")
Exp.ti
me
#
Mirror
Pos
Step
size
Tilt
tim
e
FOV
(") Data cube (px)
Durat
ion
Repe
ats
Cal
c.
Dat
a
Vol
Total
Data
Vol
Total
Durat
ion
Data
rate
Comm
ents
= study
(max=32)
2/4/6/30
(s) =#exposures
(arcsecs)
(secs)
X Y X Y Z (min) (MB)
Compre
ssed
(MB)
(min) (kbits/s)
CME
Watch
AR,
promine
nce,
limb
5
spectra
l
profile
s, 10
intensi
ties
5 32 10 4 30 90 4 0.1
4
36
0
84
0 90 800 180 45.32 30
777.
60 45.45
1356.
2 4.46
scan
direct
ion
alo
ng
slit
spectra
l
dimens
ion
Spectral-line performance
Line-SNR and spatial resolution combinations:
line-SNR >10 at 4"x2" in C-III, O-VI, Ne-VIII and
at 4"x4" in Mg-IX (in ARs)
SPICE 30"-wide Movie
Observation configuration and data rate calculation
Observing
mode
Tar
get
Line
List
#Line
Profil
es
#Px/Pr
ofile
#
Intensi
ties
Slit
(")
Exp.ti
me
#
Mirror
Pos
Step
size
Tilt
tim
e
FOV
(") Data cube (px)
Durat
ion
Repe
ats
Cal
c.
Dat
a Vol
Total
Data
Vol
Total
Durat
ion
Data
rate
Comm
ents
= study (max=
32)
2/4/6
/30 (s)
=#expos
ures
(arcse
cs)
(sec
s) X Y X Y Z (min)
(M
B)
Compre
ssed
(MB)
(min) (kbit
s/s)
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30"-wide
Movie
(sit and
stare)
centr
e,
AR
2
line
profi
les
2 32 0 30 5 120 0 0 3
0
84
0 120 800 64 10.2 1
12.
29 2.626 10.0 34.44
To be
run very
rarely,
for
high-
cadenc
e
variabil
ity
OR
1
profi
le
extract
full slit
width
scan
direct
ion
alo
ng
slit
spectra
l
dimens
ion
Spectral-line performance
N/A
SPICE 90"-wide Movie
Observation configuration and data rate calculation
Observing mode
Target
Line List
#Line Profile
s
#Px/Profile
# Intensi
ties
Slit (")
Exp.time
# Mirror
Pos
Step size
Tilt tim
e
FOV (")
Data cube (px) Duration
Repeats
Cal
c. Dat
a
Vol
Total Data
Vol
Total Durat
ion
Data rate
Comments
= study (max=
32)
2/4/6/
30 (s)
=#expos
ures
(arcse
cs)
(sec
s) X Y X Y Z (min)
(M
B)
Compre
ssed
(MB)
(min) (kbits
/s)
90"-
wide
Movie
centr
e,
AR
2
line
profi
les
2 32 0 30 5 3 28 0.9
8
8
6
84
0 3 800 64 0.435 40
12.
29 2.626 11.3 20.28
To be
run very
rarely,
for
high-
cadence
variabil
ity
OR 1 profi
le
extract full slit
width
scan direct
ion
along
slit
spectral dimensi
on
Spectral-line performance
N/A
SPICE Waves mode
Observation configuration and data rate calculation
Observ
ing
mode
Tar
get
Line
List
#Line
Profile
s
#Px/Pr
ofile
#
Intensi
ties
Slit
(")
Exp.ti
me
#
Mirror
Pos
Step
size
Tilt
tim
e
FOV
(") Data cube (px)
Durat
ion
Repe
ats
Calc
.
Dat
a
Vol
Total
Data
Vol
Total
Durat
ion
Data
rate
Commen
ts
= study (max= 2/4/6/ (s) =#expos (arcse (sec X Y X Y Z (min) (MB Compre (min) (kbits
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Page 143/312 Date 10 Jul 2017- Issue 0 Rev 1
32) 30 ures cs) s) ) ssed
(MB)
/s)
Waves
(sit
and
stare)
QS,
CH,
AR
C III,
O
VI,
NeV
III
profi
les
3 64 0 4 5 720 0 0 4 84
0 720 800 192 60 5
109.
23 114.13 300.0 50.72
For
oscillatio
ns
measure
ments
TBC
scan direct
ion
along
slit
spectral dimens
ion
Spectral-line performance
line-SNR > 10 at 4"x2" in C-III, O-VI, Ne-VIII (in ARs) and
at <4"x6" in Ne-VII (CHs and QS)
SPICE Two-Exposure mode
Observation configuration and data rate calculation
Observ
ing
mode
Tar
get
Lin
e
List
#Line
Profile
s
#Px/Pr
ofile
#
Intensi
ties
Slit
(")
Exp.ti
me
# Mirror
Pos
Step
size
Tilt
tim
e
FOV
(") Data cube (px)
Durat
ion
Repe
ats
Cal
c.
Dat
a
Vol
Total
Data Vol
Total
Durat
ion
Data
rate
Comm
ents
= study (max=
32)
2/4/6/
30 (s)
=#expos
ures
(arcse
cs)
(sec
s) X Y X Y Z (min)
(M
B)
Compre
ssed
(MB)
(min) (kbits
/s)
Two-
exposu
re
com
bi of
6
brig
ht
and
faint lines
3
brig
ht
line
s
3 32 0 4 5 60 4 0.1
4
24
0
84
0 60 800 96 60 5
46.
08 6.464 275.7 2.86
Profiles
are used
to
monitor
saturati
on
3
fain
t
line
s
3 32 0 4 55 60 4 0.1
4
24
0
84
0 60 800 96 5 46.08 SUM!
TBC
max
max
scan directi
on
along
slit
spectral dimensi
on
simultan
eous exposure
s!
Spectral-line performance
line-SNR > 10 at 4"x2" in C-III, O-VI, Ne-VIII and Mg-IX (in ARs)
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SPICE Full Raster Scan
Observation configuration and data rate calculation
Observi
ng
mode
Targ
et
Line
List
#Line
Profile
s
#Px/Pro
file
#
Intensit
ies
Slit
(")
Exp.ti
me
# Mirror
Pos
Step
size
Tilt
tim
e
FOV
(")
Data cube
(px)
Durati
on
Repe
ats
Cal
c.
Dat
a
Vol
Total
Data Vol
Total
Durati
on
Data
rate
Comme
nts
= study (max=
32)
2/4/6/
30 (s)
=#exposu
res
(arcse
cs)
(sec
s) X Y X Y Z (min)
(M
B)
Compres
sed (MB) (min)
(kbits
/s)
Full
raster
scan
few
stro
ng lines
3 32 0 2 4 480 2 0.07 96
0
84
0
48
0
80
0
9
6 32 1
73.7
3 32.6
Campai
gn not
listed in
EID-B
but in
RS WG doc for
first
orbits
TBC
16
ma
x
14
ma
x
(RS
WG
doc)
Spectral-line performance
line-SNR > 10 at 4"x2" in C-III, O-VI, Ne-VIII (in ARs) and
at <4"x6" in Ne-VII (CHs and QS)
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4.9 STIX
The Spectrometer Telescope for Imaging X rays (STIX) provides imaging spectroscopy of solar
thermal and non-thermal X-ray emissions from ~4 to 150 keV, with unprecedented sensitivity and
spatial resolution (near perihelion), and good spectral resolution.
Observationally, STIX determines the location, spectrum and timing of transient X-ray emission on
the Sun at energy ranges that encompass emission from both hot thermal plasmas and
bremsstrahlung from energetic electrons. The properties of the electrons that
generated the X-rays can be inferred from their X-ray spectrum. The distinction between a thermal
plasma and non-thermal electron population is based on the shape of the X-ray spectrum with the
latter having a characteristic power law (or broken power law) profile
and the former providing a black body spectrum (corresponding to 106 to 108 K). The spectra are
very steep and so good spectral resolution is required for their interpretation.
There is also an Iron line complex at 6.7 keV which, if isolated, can be interpreted in terms of the
thermal electron population. Since a typical flare typically generates both thermal and non-thermal
emission, which often are not co-located (for example at the top
and footpoints of magnetic loops respectively), both good spatial and spectral resolution are
required.
The observational objectives are achieved by imaging the Sun as a function of time and energy with
enough spatial, spectral and temporal resolution to match the sources of interest. Comparing the
resulting images at different energies yields the X-ray spectra of
individual features (e.g. footpoints or flaring loops). Comparing the images as a function of time
discloses the temporal behavior of the hot plasma and accelerated electrons. The data can also be
combined to yield spatially-integrated light curves and spectra. In all
cases, the basic observational datum is a single, photometrically-accurate image corresponding to a
well-defined time and energy interval.
Focusing optics are not a feasible option for arcsecond-class hard X-ray imaging within Solar
Orbiter constraints. As a result, STIX uses an indirect Fourier imaging technique based on X-ray
collimation. Conceptually, the instrument is made up of three
mechanically separate modules: X-ray transparent windows; a passive imager containing front and
rear grids; and a Detector/Electronics Module (DEM) containing passively cooled X-ray detectors
and electronics.
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4.9.1 STIX observables
4.9.2 STIX modes
TM figures
Allocated TM 0.7 kbits/s
Download capacity per orbit
226.8 MB
OR 1270 MB
= 1.8144 Gbits (if only acquiring during 30 days)
= 10.16 Gbits (if extrapolated to 168 days orbit)
Subtelescopes
non
Observational modes
source:
For STIX there is only 1 data acquisition mode. Acquired science data are selected on-board
autonomously (~500bps) or selected on-ground and retrieved via TC (~100bps).
Mode Data Rate (bps) Orbital Volume (Gbits)
STIX Normal Mode - automatic onboard selection of data ~500 7.3 (if acquiring full orbit)
STIX Normal Mode - TC-requested data ~100 1.45 (if acquiring full orbit)
STIX LL ~50 0.7 (if acquiring full orbit)
STIX HK ~50 0.7 (if acquiring full orbit)
STIX has only 1 science acquisition mode, which is independent from solar activity, campaigns, etc.
Data acquisition operations onboard do change based on the incoming flux. The attenuator is
automatically used during high solar activity to suppress the count rates, cadence can be changed
too.
STIX does change the data processing and packaging onboard, based on solar activity and
automatic event detection:
During non-flaring times: STIX FPGA accumulates background counts for energy calibration.
During flares:
• STIX may autonomously change attenuator and/or enable selective pixel suppression using
predefined count rate criteria.
• STIX transmits flare flag message to s/c.
• In real time: STIX FPGA converts fine native A/D channels into detector-matched ‘science
energy channels’.
• In real time: data are accumulated as a function of science energy channel, pixel and time
bins (0.1 second or larger depending on statistics).
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• In background (using pre-parameterized algorithms or in response to TC): STIX selects and
compresses data in archive memory into TM-ready packets in the 'to be transmitted’ buffer.
When requested through TC: extra TM-ready packets can be selected and transferred to the SSMM
for downlink.
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4.10 SWA
SWA consists of a suite of sensors which are able to measure the three-dimensional velocity
distribution functions of the major solar wind constituents: protons, alpha particles and electrons.
The basic moments of the distributions, such as density, velocity, temperature tensor, and heat flux
vector need to be obtained under all solar wind conditions and must be sampled sufficiently rapidly
to characterize fully the fluid and kinetic state of the wind. In addition, measurements of
representative high-FIP elements (the C, N, O group) and of low-FIP elements (such as Fe, Si or
Mg) are required. These measurement challenges require an instrument suite comprising 3 distinct
sensors:
• The Electron Analyser System (EAS) to make high temporal resolution measurements of the
core, halo and ‘strahl’ electron VDFs and their moments;
• The Proton-alpha sensor (PAS) to measure the VDFs of major ion species at high time
resolution and determine their moments;
• The Heavy Ion Sensor (HIS) to measure the 3-D VDFs and determine the abundance and
charge states of prominent minor ion species.
The SWA-EAS will resolve the full 3-D velocity space distributions of solar wind electrons with
high cadence (<10 s time resolution). Since the electron thermal velocity is much higher than their
bulk velocity, even in the solar wind, this can only be achieved by a sensor, or sensors, that have a
combined field-of-view (FoV) covering a large fraction of the 4π steradians of the full sky. In
addition, the sensor is required to provide pitch angle distributions with high time resolution (ideally
at a cadence of 0.125 s). Magnetic field data from MAG is used onboard to produce the PAD’s and
is thus required at this high cadence. Note this cadence, which is slightly slower than that detailed in
the PDD, has been selected to allow synchronisation with anticipated magnetometer data rates.
The SWA-PAS sensor is capable of resolving the full 3-D VDFs of solar wind protons and alpha
particles with high cadence (<10 s time resolution), as well as measuring the bulk plasma
parameters at ultra-high time resolution of ~0.1 s to characterize the global structure and dynamics
of the 3-D inner heliosphere and improve our basic understanding of the kinetic processes and
microstate of the evolving solar wind from 0.28–1.3 AU. The PAS energy coverage and resolution,
field-of-view, angular coverage and resolution, geometric factor and time resolution are such that it
can measure solar wind protons and alpha particles for more than 99% of the time during the SO
mission profile.
The Heavy Ion Sensor (HIS) must address two fundamentally different sets of measurement
objectives. First, HIS will measure the composition and 3-D VDFs of heavy ions in the bulk solar
wind. Second, HIS will measure the composition and 3-D VDFs of the major constituents in the
suprathermal energy range. The sensor is able to resolve the full 3-D velocity distribution functions
of the prominent heavy ions at a resolution of 5 minutes in normal mode and 30 s in burst mode.
Additionally, HIS will measure 3-D VDFs of alpha particles at 4 second resolution in burst mode.
Measurements will be made up to 60 keV/e, with 64 energy steps (6-10% resolution) and 6º x 6º
resolution (15 azimuthal sectors and 6 elevation steps). The mass resolution (m/Δm) is ~5. These
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will be the first such measurements in the inner heliosphere and hence any new result from HIS will
be breakthrough science.
4.10.1 SWA observables
4.10.2 SWA modes
Note that all data volumes & rates are based on powers of 10, i.e. 1Gbit = 109 bits. HK and LL are
included in normal mode.
TM Figures
Allocated TM 14.5 kbits/s
Download capacity per orbit (168 Days) 26.309 GB = 210.470Gbits
SSMM allocation
Operational Modes
Sources: SWA Budget Report Issue 1 (scaled to reach 14500 not 14848), SWA input to SAP
Planning
Mode Data Rate (bps) Duty Cycle (%) Duty Cycle (hrs / orbit) Orbital Volume (Gbits)
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SWA Normal Mode 11551 99.653 4018 167.083
SWA Burst Mode (Scheduled) 413350 0.347 (min) 14 20.833
SWA Triggered Mode 447504 0.347 (max) 14 22.554
SWA's Book Keeping Algorithm trades off triggered burst modes against scheduled burst modes
and it isn't entirely clear from available documentation what the native data rates of these burst
modes are, only what proportion of the SWA TM budget is available after normal mode, LL and
HK are accounted for. Discussions with C. Owen and ISWG imply that EAS can burst for 12
minutes per day total in scheduled burst mode, and that one triggered burst mode (dump of 5 minute
rolling buffer) is equivalent to 7 minutes of scheduled burst mode. For simplicity's sake we assume
the same is true for PAS (which has many different burst mode options in any case so this could
probably reflect some form of reality). HIS does not respond to triggers.
For the purposes of calculating rates and duty cycles, then, we assume the following:
1. Scheduled burst mode involves EAS, HIS and PAS.
2. Scheduled burst mode replaces normal mode.
3. EAS and PAS respond to triggers; HIS does not.
4. Rate calculations are based on 5 minutes scheduled burst mode per day and 1 trigger
response per day. A trigger produces 7/12 of the daily burst mode volume for EAS and PAS.
Thus scheduled bursts produce 5/12 of the available daily burst mode volume for EAS and
PAS and 100% of the available burst mode volume for HIS.
SWA Normal Mode
SWA Normal mode is expected to operate for the majority of the mission, except during scheduled
bursts (~5 minutes per day). Normal Mode Provides the following data products (note that LL
products are listed here for convenience). Housekeeping represents 300bps which is included
normal mode on the summary page but not listed explicitly here
Product Cadence (s) Data Rate (bps)
EAS 3239
Electron Moments 4
3D VDFs 100
Strahl Energy Shell (LL) 100
PAS 2864
Proton moments (LL) 4
3D proton & alpha VDFs 4
8s duration proton reduced VDF snapshots 300
HIS 5148
Heavy ion rates and VDFs 300
Alpha rates and VDFs 30
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Total 11551
SWA Burst Mode (Scheduled)
SWA burst mode is scheduled and coordinated with the rest of the in situ payload. Current thinking
is that there will be five minutes of 'protected' burst mode per day (on average - this will likely be
weighted towards perihelion) and a further five minutes that can be traded off against triggers (for
PAS and EAS). Scheduled burst mode replaces normal mode. PAS Burst mode is highly
configurable, trading off cadence and angular/energy coverage. Duty Cycles on the previous page
are calculated by assuming triggers occur with sufficient frequency that the second five minute
scheduled burst is always replaced.
Product Cadence (s) Data Rate (bps)
EAS 130388
2D Electron Pitch Angle Distribution 0.125
PAS 189257
High time resolution 1D & 2D VDFs < 1
HIS 93704
Heavy ion rates and VDFs 30
Alpha rates and VDFs 3
Total 413350
SWA Triggered Mode
SWA triggered mode will definitely involve EAS. PAS involvement is subject to power constraints.
Here we assume PAS is involved. In case that it isn't there will be a proportional increase in PAS
scheduled burst mode data volume. Triggers modes do not interrupt normal mode operations, and is
not technically a mode but rather represents a dump of a 5 minute duration rolling buffer to the
SSMM.
Product Cadence (s) Data Rate (bps)
EAS 182544
3D Electron VDFs 1
PAS 264961
3D proton & alpha VDFs 1
Total 447504
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5 SCIENCE ACTIVITIES
5.1 Introduction and SOOPs
TBW
5.2 List of SOOPs
This section describes all SOOPs that have been defined up to now and are meant to cover all
science objectives. The current list is given below with all details per SOOP given in the following
sections.
• I_DEFAULT
• L_BOTH_HRES_LCAD_CH_Boundary_Expansion
• L_BOTH_LRES_MCAD_Pole-to-Pole
• L_BOTH_MRES_MCAD_Farside_Connection
• L_BOTH_MRES_MCAD_Flare_SEPs
• L_FULL_HRES_HCAD_Coronal_Dynamics
• L_FULL_HRES_HCAD_Eruption_Watch
• L_FULL_HRES_LCAD_MagnFieldConfig
• L_FULL_HRES_MCAD_Coronal_He_Abundance
• L_FULL_LRES_MCAD_Coronal_Synoptic
• L_FULL_LRES_MCAD_ProbeQuadrature
• L_FULL_MRES_MCAD_CME_SEPs
• L_IS_SoloHI_STIX
• L_IS_STIX
• L_SMALL_HRES_HCAD_Fast_Wind
• L_SMALL_HRES_HCAD_SlowWindConnection
• L_SMALL_MRES_MCAD_Ballistic-connection
• L_SMALL_MRES_MCAD_Connection_Mosaic
• R_FULL_HRES_HCAD_Density_Fluctuations
• R_FULL_LRES_HCAD_GlobalHelioseismology
• R_SMALL_HRES_HCAD_AR_Dynamics
• R_SMALL_HRES_HCAD_Ephemeral
• R_SMALL_HRES_HCAD_PDF_Mosaic
• R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure
• R_SMALL_HRES_HCAD_RSburst
• R_SMALL_HRES_HCAD_WaveStereoscopy
• R_SMALL_HRES_LCAD_Composition_vs_Height
• R_SMALL_HRES_LCAD_FineScaleStructure
• R_SMALL_HRES_MCAD_PolarObservations
• R_SMALL_MRES_MCAD_AR_LongTerm
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5.2.1 I_DEFAULT
Standard In situ operations out of remote sensing windows. Will contribute to all in situ only
objectives. Coordinated bursts schedules to obtain equal coverage in heliocentric distance rather
than time.
Instrument Mode Comment
EPD Normal + Burst
MAG Normal + Burst
RPW Normal + Burst Triggers Active
SWA Normal + Burst
SAP
objective Target Duration
Opportunity(e.g.,
orbital
requirements,
solar cycle phase,
quadrature ...)
Operationa
l
constraints
Additional comments
2.3.2.2
Identify
interplanetary
shocks and
characterise
their spatial
and temporal
evolution
In Situ
Sufficient
Coverage
for good
statistics
Multi
spacecraftstudy(SPP
) in multiple
orientations (radial,
spiral, quadrature
alignments) –
should come for
free.
For quadrature, one
s/c would need RS,
and the other in-situ,
so again for free if
WISPR R< 0.5 AU
EMC Quiet
Burst modes
scheduled for flat
radial coverage
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2.3.3 Resolve
the
interplanetary
shock field
and plasma
structure
down to the
spatial and
temporal
scales
comparable
and smaller
than the
typical ion
scales.
In Situ
Sufficient
Coverage
for good
statistics
EMC Quiet
Burst
Modes,mostlytriggere
d.
SPP Good to have.
2.3.4 Shock-
surfing
acceleration
mechanism
In Situ
Heliosphere
Sufficient
Coverage
for good
statistics
Radial dependence EMC Quiet
Normal + Burst modes
Triggered by RPW
(this is important to
know for low
telemetry periods
when the triggering
would be de-
activated).
SPP Good to have.
2.3.5
Understand
the radio
emissions
from the
ICME driven
shocks
In SItu
Sufficient
Coverage
for good
statistics
Good coverage of
different radii EMC Quiet
RPW Triggers
important.
1.2.2.7 Study
the
correlation
degree
between
velocity and
magnetic
field
fluctuations
in the
interplanetary
space
In Situ
Sufficient
Coverage
for good
statistics
Good coverage of
different radii and
latitudes
EMC Quiet Normal Mode
Sufficient
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1.1.2.6
Disentangle
the spatial
and temporal
variability of
the slow wind
In Situ
1.1.2.8
Determine
the velocity,
acceleration
profile and
the mass of
the transient
slow wind
flows
In Situ
1.1.4.1.3
Identify
reconnection
exhausts in
the solar wind
In Situ
1.1.4.1.5
Identify and
characterise
the solar wind
reconnection
physics in
current sheets
with
thickness
down to the
ion scales and
smaller
1.2.2.1
Determine
where energy
is deposited
in the solar
wind
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1.2.2.2 What
drives the
evolution of
the solar wind
distribution
functions in
situ?
EMC Quiet
• Long-term
observations in
normal mode.
• Good radial
spread of short
duration bursts
over the
mission.
• Radial
coverage
including
perihelion.
• SO-SPP(-
Earth) radial
alignments.
1.2.2.4
Identify and
characterise
the solar wind
reconnection
physics in
current sheets
with
thickness
down to the
ion scales and
smaller.
3.1.1.2.2
Composition
variations
In Situ
Seed
population
specification
from the
heavy ion
composition
of solar wind
and
suprathermal
s in the inner
heliosphere
timing
Statistics
Normal mode
sufficient.
SWA/HIS
EPD/SIS
RPW/MAG not
needed.
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3.1.1.3 How
are superhalo
particles
accelerated
continuously
in the corona
and solar
wind?
In Situ
Particles
acceleration
Statistics Good radial spread.
However, SPICE
could be used to
observe non-thermal
electron distributions
at the limb, e.g in
combination with SPP
in quadrature.
However, if SPICE
cannot observe the
superhalo energies,
then it is not needed
here (TBC).
3.1.1.5 Do
proton-
amplified
Alfvén waves
play a role in
accelerating
particles at
shocks?
In Situ
Alfvén
waves role in
SEP
acceleration
Statistics
Perihelion SWA Burst mode
3.1.2.0.1
Measure the
enhancements
of trans-iron
elements in
impulsive
SEPs (to be
deleted -
trans-iron
SEPs cannot
be detected
by EPD)
In Situ Statistics
EPD (SIS) Normal
mode sufficient
SWA/MAG/RPW not
needed
RS window not
necessarily needed (for
this specific
objective), however it
would be great if we
can identify the solar
source as for all
Objective 3 (SPICE
would be nice to have)
It seems to be feasible
with SIS.
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3.1.3.1 How
shocks can
accelerate
electrons to
relativistic
energies
(never
observed for
shocks near 1
AU)?
In Situ Statistics
Perihelion Normal + Burst modes
3.2.4 How do
large and
small-scale
structures
modulate
particle
fluxes?
In Situ
Corona +
Heliosphere
Long term
observation
s
Radial dependence
RS context is good
to have (not
necessary).
Good to
have Metis
compatible
May need full-disk RS
+ Metis for context.
SWA regular
scheduled and
triggered burst by
RPW (Multiple bursts
per day only needed
for the small-scale
structures but we
cannot predict when
this will happen in
order to plan for
multiple scheduled
bursts per day, except
if we know we are in a
particularly active
period. Large scale
CIRs/ICMEs may
rather need triggered
burst)
3.2.0 What
controls the
escape of the
particles to
the
heliosphere?
In Situ
Heliosphere Statistics Perihelion
EPD
MAG
RPW
1.1.3.2.3
How does the
heliospheric
magnetic
field
disconnect
from the Sun?
In Situ
Heliosphere
Radial dependence
& Perihelion
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3.2.1.2
Measurement
s of SEP
events time
profiles and
anisotropy in
order to probe
solar wind
turbulence
In Situ
Heliosphere Statistics
Close to perihelion
and different range
of distances
Needs to be done for a
range of distances to
monitor the SEP
contributions from
solar source and IP
turbulence.
Also partly addressed
by L_IS_STIX and
better by
L_IS_SoloHI_STIX.
3.2.1.3
Identify
dropouts and
measure
scattering of
SEPs by
turbulence.
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5.2.2 L_IS_STIX
This SOOP is aimed at understanding X-Ray emission from energetic particles. STIX is leading this
SOOP, while IS payload provides continuous observations.
Default SOOP duration: currently modelled to run during all RS windows
Pointing requirements: disk centre preferred but not mandatory
Triggers: enabled
Instrument Mode Comment
STIX (leads) STIX Normal Mode
SoloHI N/A
Metis N/A
EUI N/A
PHI N/A
SPICE N/A
MAG MAG Normal Mode + MAG Burst Mode
EPD EPD Normal Mode + EPD Burst Mode
RPW RPW Normal Mode + Burst modes Triggers on
SWA SWA Normal Mode and SWA Burst Mode (Scheduled)
SAP objective Target Duration
Opportunity
(e.g., orbital
requirements,
solar cycle
phase,
quadrature ...)
Operational
constraints Additional comments
3.1.1.4 How
can SEPs be
accelerated to
high energies so
rapidly?
Corona +
Heliosphere
Particles
acceleration
Statistics
Perihelion
SWA Burst mode
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3.1.2.3 Flare
seed particles.
Corona &
Heliosphere
SEP events
with large
initial Fe/O
ratios
Particles
accelerated
by flares
Statistics
Close to
perihelion
3.1.2.7 How are
so many
electrons
accelerated on
such short time
scales to
explain the
observed hard
X-ray fluxes?
Corona &
Heliosphere
X-ray
signatures of
energetic
particle
interactions
at loop
footpoints,
or on loops
themselves
Statistics
Perihelion (if
possible, but to
be connected to
earth is more
important)
Best when the
solar limb from
SO is connected
to Earth, or
other s/c (not
necessarily the
limb, but behind
the limb, up to
20 degrees)
To use RS context from
Earth
STIX: High-cadence
energy resolved imaging
3.2.1.2
Measurements
of SEP events
time profiles
and anisotropy
in order to
probe solar
wind turbulence
Heliosphere Statistics
Close to
perihelion and
different range
of distances
Needs to be done for a
range of distances to
monitor the SEP
contributions from solar
source and IP turbulence
Also partly addressed
by I_DEFAULT or better
with L_IS_SoloHI_STIX.
3.2.1.3 Identify
dropouts and
measure
scattering of
SEPs by
turbulence.
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5.2.3 L_IS_SoloHI_STIX
This SOOP is aimed at measuring variability in SEP events trough corona and heliosphere. IS
instruments provide continuous observations while SoloHI gives spatial context.
Default SOOP duration: 10 days
Pointing requirements: disk centre
Triggers: enabled
Instrument Mode Comment
SoloHI
SoloHI Shock or
Turbulence and
Synoptic modes
The distance to the sun will define the SoloHI mode: shock
and turbulence mode, in combi with HI_SYN_PER to be
scheduled only within 0.4AU, further out only synoptic modes
(HI_SYN_NEAR or HI_SYN_FAR).
EUI N/A
PHI N/A
STIX STIX Normal Mode Needed only for 3.2.1.2
SPICE N/A
Metis N/A
EPD Triggered Burst Mode
MAG MAG Normal Mode
and MAG Burst Mode
RPW RPW Normal Mode
and RPW Burst Mode
SWA
SWA Triggered Mode
or SWA Burst Mode
(Scheduled)
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SAP objective Target Duratio
n
Opportunity
(e.g., orbital
requirement
s, solar cycle
phase,
quadrature
...)
Operationa
l
constraints
Additional comments
3.1.1.2.3
Warped shock
fronts
Corona +
Heliosphere
Gradual SEP
events
Statistics
Good radial
spread.
Multispacecraft
Multiview point data preferred
to address this objective
IS Triggered burst
3.1.1.2.4
Turbulence and
inhomogeneitie
s
Magnetic
field, plasma
wave and
solar wind
measurement
s to
determine
turbulence
levels and
identify
shock
passages
Statistics
Good radial
spread.
Near
Perihelion
(EPD)
EMC
Quiet
SoloHI is used with
Solarprobe (turbulence mode).
Burst Modes: scheduled or
triggered.
MAG: high-cadence magnetic
field
RPW: high-cadence electric
and magnetic field, power
spectral densities
3.2.1.2
Measurements
of SEP events
time profiles
and anisotropy
in order to
probe solar
wind
turbulence
Heliosphere
Statistics
Close to
perihelion
and different
range of
distances
Needs to be done for a range
of distances to monitor the
SEP contributions from solar
source and IP turbulence
SoloHI Turbulence mode
STIX
Also partly addressed
by I_DEFAULT or L_IS_STI
X.
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3.2.3 Properties
of near-Sun
shocks, their
fluctuations
and particle
acceleration
Corona +
Heliosphere Statistics Perihelion EMC Quiet
MAG burst mode (good
trigger needed)
SWA burst mode (triggered)
RPW
SoloHI Shock mode
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5.2.4 L_FULL_LRES_MCAD_Coronal_Synoptic
Synoptic SOOP designed for CME science and global coronal structure. Typically runs for a whole
RS window.
Default SOOP duration: 10 days
Pointing requirements: disk centre
Triggers: IS + Metis triggers active
Instrument Mode Comment
EUI FSI Synoptic mode (S) 10 min cadence
Metis
METIS standard modes: GLOBAL mode,
METIS special modes CMEOBS whenever triggered
(model as max 3 CMEs per RSW)
CME Watch On
PHI FDT Synoptics: PHI science mode 6 with 6hr cadence 6 hour cadence
SoloHI Normal Operations: SYN + SHOCK (at perihelion)
(modelled as HI_SYN_NEAR for now)
SPICE N/A
STIX STIX Normal Mode
EPD Normal + Burst Mode
MAG Normal + Burst Mode
RPW Detection Mode Burst Triggers Active
SWA Normal + Burst Mode
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SAP objective Target Duration
Opportunity
(e.g., orbital
requirements,
solar cycle phase,
quadrature ...)
Operational
constraints Additional comments
2.1.1.2 CME
structure
Full Disk Long
Duration
Quadrature ideal,
perihelion would
be best, but useful
at all distances.
No Offpoints
For EUI could do
onboard prioritisation
to increase cadence -
discard 9/10 images
unless there's a trigger.
2.1.1.3 CME
evolution
Full Disk Long
Duration
Higher latitude
orbits particularly
interesting.
Angular Separation
from other
spacecraft is a
bonus
No Offpoints Temperature evolution
from Metis in UV
1.1.3.2 How
does the Sun's
magnetic field
link into space?
Corona +
Heliosphere
Metis
compatible
Coordination with SPP
is a bonus
1.3.3 Plasma
turbulence
variability
Full Disk Long
Duration
Perihelion good for
SoloHI
contribution
Several Latitudes
No Offpoints
EMC Quiet
Radial alignment with
SPP useful
1.3.4 Plasma
turbulence
anisotropy
Full Disk Long
Duration
Radial
Dependence,
EMC Quiet
No Offpoints
SPP important to have
1.2.1.10
Heating in
flaring loops vs
heating in
active regions
Full Disk Long
Duration
More statistical
study: having
STIX on,
observing all the
flares and using
EUI synoptics to
find out about the
source region
N/A
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3.1.2.2
Evaluate how
significantly
large flares
contribute
directly to
gradual SEP
events
Corona &
Heliosphere Statistics
Near perihelion in
order to get 'as-
pure-as -possible'
signal
Metis
compatible
FSI 10 min cadence
(default)
SoloHI, Metis, STIX
Low resolution does
not seem to be a
problem
3.1.4.1 Hard X-
ray emission of
escaping
electron beams
(thin-target
emission)
Corona
Statistics
Near perihelion to
get best spatial
resolution with
STIX
RPW
STIX
3.1.4.2 X-ray
emission from
electrons
accelerated at
CME shocks
Corona &
Heliosphere
Statistics
NOT at Perihelion Metis
compatible
RPW
STIX
Metis
We cannot hunt for
this event (too specific
requirements: CME
lifting off from behind
the limb to shield the
X-rays from AR
footpoints, to allow
STIX to observe hard
X-rays from shock)
and we hope we get it
for free through this
SOOP
4.2 What are
the properties
of the magnetic
field at high
solar latitudes?
[source: SOL-
PHI-MPS-
MN5100-TN-2]
Full disk
synoptic
during
solar
minimum
opposition with
NEO
Co-
observations
with NEO
global field
maps
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5.2.5 L_FULL_LRES_MCAD_ProbeQuadrature
Default SOOP duration: 3 days
Pointing requirements: Disk Centre
Triggers: only IS triggers enabled
SOOP designed to study the corona while Solar Probe Plus is in quadrature with Solar Orbtier, PHI
& EUI provide context, Metis and SoloHI provide imagery of solar wind that will be encountered
by solar Probe Plus.
Instrument Mode Comment
PHI PHI science mode 6 (FDT) Full FoV, 6 hours cadence, 5 quantities
EUI FSI Synoptic mode (S)
Metis METIS special modes (PROBE mode) 1-10 minute cadence depending on
heliocentric distace
SoloHI SoloHI mode depending on heliocentric
distance
STIX STIX Normal Mode
EPD Normal Mode
MAG Normal Mode
RPW Detection Mode Burst Triggers Active
SWA Normal Mode
SAP objective Target Duration
Opportunity
(e.g., orbital
requirements, solar
cycle phase,
quadrature ...)
Operational
constraints
Additional
comments
1.2.2 What
mechanisms heat
and accelerate the
solar wind?
helmet
streamers
several hours
METIS at
high cadence
SPP in quadrature
with Solar Orbiter
METIS co-
observations
required
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5.2.6 L_FULL_MRES_MCAD_CME_SEPs
This SOOP is aimed at measuring in-situ SEPs properties and kinetics and linking them to the
observing SEPs CME acceleration. Location, timing, and motion of CMEs and shocks.
SoloHI and Metis are leading this SOOP, while IS payload provides continuous observations. Disk
centre pointing needed.
Default SOOP duration: 10 days
Pointing requirements: disk centre
Triggers: IS+Metis triggers active
Instrument Mode Comment
SoloHI
(leads) SoloHI Shock and Synoptic modes
Metis (leads) METIS standard modes GLOBAL, METIS special modes
CMEOBS
Disk centre pointing
needed
EUI FSI Synoptic mode (S)
FSI 10 min cadence
(default)
PHI N/A
SPICE N/A
STIX STIX Normal Mode
EPD EPD Normal Mode
MAG MAG Normal Mode
RPW RPW Normal Mode
SWA SWA Triggered Mode or SWA Burst Mode (Scheduled)
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SAP objective Target Duration
Opportunity
(e.g., orbital
requirements, solar
cycle phase,
quadrature ...)
Operational
constraints
Additional
comments
3.1.1.1 Where
and when are
shocks most
efficient in
accelerating
particles?
Corona +
Heliosphere
Particles
acceleration
Statistics
Good radial spread.
Good to have
Metis
compatible
SWA 5 minutes
captures
SoloHI needed for
estimating the total
CME energy
FSI not needed
3.1.1.2.1
Intensity
variability
Heliosphere
Gradual
SEP events
Statistics
Preferable with Earth
images and SO +/- 70
degrees from Earth-
Sun line (probably
because gradual
events extend up to
100 degrees in
longitude)
Metis
compatible
MAG/RPW/SWA
not needed
FSI 5 min cadence
for source
identification (FSI is
10 min cadence by
default)
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3.1.1.7 Are there favourable environments for
particle acceleration?
Corona +
Heliosphere Statistics
FSI 10 min cadence
(default)
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5.2.7 L_FULL_HRES_LCAD_MagnFieldConfig
High spatial resolution synoptic operations with full disk RS instruments, Metis and insitu
instruments. Low cadence (1/hour to 1/day). Most relevant for windows or parts of windows with
very restricted TM and for out of window synoptic campaigns if these are possible.
Default SOOP duration: 1 day
Pointing requirements: disk centre
Triggers: only IS triggers active
Instrument Mode Comment
PHI PHI science mode 6: FDT at highest spatial resolution
(2Kx2K), cadence at least 1 per day
model cadence as twice per
day
EUI FSI Reference Synoptic mode (R), highest spatial
resolution (3Kx3K), low cadence (hrs to 1/day):
Configure instrument to
enhance off-limb structures
model cadence as 1/day
Metis
synoptic programme for magnetic field structure
(METIS standard modes: GLOBAL and/or LT-
CONFIG)
no reaction to CMEs needed
EPD normal
MAG normal
SWA normal
RPW normal
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SAP objective Target Duration
Opportunity
(e.g., orbital
requirements, solar
cycle phase,
quadrature ...)
Operational
constraints
Additional
comments
1.1.3.2.1 How
does the Sun's
magnetic field
change over
time?
Full disk,
photosphere
and corona
long-term
observations
(RS synoptics
in between
RSW?)
• far side
• from Earth
alignment to
far side
(different
separation
angles with
Earth for
projection
effects)
• range of
latitudes,
definitely
repeat at
highest
latitudes late in
mission
RSW
placement if
no synoptics
allowed
EMC quiet for
linkage
science
1.1.3.2.2 How is
the heliospheric
current sheet
(HCS) related to
coronal
structure?
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5.2.8 L_FULL_HRES_MCAD_Coronal_He_Abundance
Simultaneous observations of the resonantly scattered component of He+ emission by EUI/FSI 30.4
nm, and neutral hydrogen by Metis Lyα (121.6 nm) to examine helium abundances in the corona /
inner solar wind. These can be compared with SWA.
Compared to its value in the solar convective envelope, the helium abundance in the in-situ
measurements of the fast and slow solar wind has long been known to be depleted relative to
hydrogen, with occasional transient exceptions (Bochsler 1998, SSRv 85, 291). In the slow solar
wind, the degree of depletion has more recently been shown to depend upon the wind speed and the
level of solar activity (Aellig et al. 2001, GeoRL 28, 2767). Measurements of the helium abundance
in the corona, associated to measurements of the coronal outflow velocity, will provide evidence for
the degree of correlation between wind speed and helium abundance and allow identification of the
source regions of the slow wind streams with different helium abundance. During quasi-corotation
the intrinsic evolution of magnetic topology will be observed and thus its influence on the wind
parameters (such wind outflow velocity and helium abundance) will be directly assessed. The
abundance can be derived from simultaneous observations of the resonantly scattered component of
singly ionized helium by EUI/FSI in its 30.4 nm channel and of that of neutral hydrogen by Metis in
Lyα (121.6 nm).
Useful contributions can be given by SPICE, mapping the near-surface elemental abundances,
including that of helium (TBC), which constitutes a reference for establishing abundance variations
in the wind. PHI can also contribute, providing data suitable for coronal magnetic field
extrapolations.
Note: this is essentially an EUI/Metis sub-objective, but in-situ may be interested also. For instance,
SWA/HIS will measure the α/p density ratio.
Instrument Mode Comment
EUI FSI Synoptic mode (S), 20-min
cadence
Occultor likely to be used in conjunction with this
mode
Metis
MAGTOP, 20-min cadence,
duration ≥ 2 hours
WIND, 20-min cadence,
duration ≥ 2 hours
Global maps of:
• neutral hydrogen Lyα intensity
• electron density
• outflow velocity in corona
PHI PHI science mode 6: FDT, 6-
hour cadence Data suitable for coronal magnetic field extrapolation
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SoloHI Measure solar wind speed above potential source
region in co-ordination with Metis
SPICE SPICE Composition Mapping
STIX
EPD Normal Mode
MAG Normal Mode
RPW Detection Mode Burst Triggers Active
SWA Normal Mode
SAP objective Target Duration
Opportunity
(e.g., orbital
requirements, solar
cycle phase, quadrature
...)
Operational
constraints
Additional
comments
1.1.2 Source
regions of the
slow solar
wind
1.1.2.12
Abundance of
helium as a
function of
height and
latitude in the
corona as a
tracer of the
source regions
of the slow
solar wind
Inner
corona
within
EUI/FSI
and
Metis
FOVs
few hours
per day
• Inside 0.45–0.5
AU (optimised
distance for
EUI/FSI's
occulter)
• or perihelion, for
quasi-corotation
measurements.
Disc-centre
pointing
Earth view
beneficial before
RSW to estimate the
global solar
magnetic field. Is
this because another
view than PHI/FDT
is needed, e.g., to get
front- and back-side
measurements?
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5.2.9 L_FULL_HRES_HCAD_Eruption_Watch
Full-Disk, high resolution SOOP designed to catch eruptive events.
Default SOOP duration: 1 day
Pointing requirements: disk centre
Triggers: enabled for both IS and RS
Instrument Mode Comment
EUI
Global Eruptive Event Mode:
FSI Global eruptive event
mode (G)
Triggered:
EUI will be in global mode all the time (i.e. full SOOP
length) but it will only prioritize the data of 1-2 events.
Global mode generates 4,42 Gbit/hr. Let's say for now
that we flush 2 hrs of data = ~8Gbits.
Metis
METIS standard modes:
GLOBAL + CMEOBS on
trigger
(modelled as 2 CME events of
1 hr in the 1-day-SOOP)
CME Trigger on
PHI
FDT, 2-5 minute cadence,
highest spatial resolution
(PHI science mode 4 with
FDT, default is 5 mins
cadence)
Selection of data so will check LLD
SoloHI Combination of shock and
synoptics at perihelion
Combine HI_SHOCK_PER + HI_SYN_PER (each 50%
of time)
SPICE SPICE Waves mode
Could be SPICE Waves mode and SPICE Composition
Mapping interleaved if off-pointing possible.
In this case, we would use observation called
SPICE_WAVES_COMP in modelling. But
SPICE_WAVES will do for SAP v.0
STIX Standard Operations: STIX
Normal Mode
MAG Normal Mode
EPD Normal Mode
RPW Detection Mode Detection algorithms active
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SWA Normal Mode
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SAP
objective Target Duration
Opportunity
(e.g., orbital
requirements,
solar cycle phase,
quadrature ...)
Operational
constraints Additional comments
2.1.1.1
CME
initiation
Full
Disk
1 Day
(limited by
EUI/PHI
internal
memory)
Perihelion
preferred.
Quadrature with
Earth preferred.
Interesting
throughout the solar
cycle.
No
offpointing
beyond Metis
limit.
EUI/PHI at highest cadence
appropriate to spatial
resolution (could be slower if
further away). SPICE Sit &
Stare in waves mode to try and
catch EUV waves.
Metis Modes:
• GLOBAL (before the
event, if possible), min.
obs time 2 hr, data
volume ≤ 300 Mb.
• CMEOBS, starts after
CME flag rise, min. obs
time 1 hr (high
cadence, 1 min), data
volume ~ 2.137 Gb.
• GLOBAL (after the
event), min. obs time 2
hr, data volume ≤ 300
Mb.
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5.2.10 L_FULL_HRES_HCAD_Coronal_Dynamics
This SOOP is aimed at observing structures in the outer corona and linking them to the heliosphere
observed in-situ. Metis and SoloHI are leading this SOOP, while IS payload provides continuous
observations. Synoptic support from other full disk RS instruments. Disk centre pointing preferred.
Default SOOP duration: 1 day
Pointing requirements: disk centre
Triggers: only IS triggers active
See SOOP 5 defined for SOWG8 planning exercise.
(SOOP will be modelled with 1 day duration, can be repeated as many times as needed)
Instrument Mode Comment
SoloHI
(leads)
SoloHI combination of high cadence TURB and
synoptic mode (model as HI_SYN_NEAR for now)
Metis
(leads)
Generic program like WIND (METIS standard
modes) interleaved with FLUCTS (METIS special
modes) (FLUCTS runs 1 hr/day)
disk centre pointing preferred
PHI FDT synoptic: PHI science mode 6
PHI may be processing in
between observations. all data
gets downloaded
EUI FSI synoptic: FSI Synoptic mode (S) all generated data gets
downloaded
EPD Normal Mode
MAG Normal Mode
RPW Detection Mode Burst Triggers Active
SWA Normal Mode
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SPICE
(optional)
SPICE currently proposes to use either SPICE Limb
mode or SPICE Dynamics, depending on
requirement for Metis-compatibility
Limb target (not Metis-Compatible; SPICE Limb
mode) :
• Slit: 4"
• Exposure time: 5 s
• X positions: 224
• FoV: 15' x 11'
• Nº repeats: 10
• Observation time:
o 18 mins per study
o 3.2 hours total.
Active region, if before or after Metis observatiosn
(SPICE Dynamics) :
• Slit: 4"
• Exposure time: 60 s
• X positions: 128
• FoV: 8.5' x 11'
• Nº repeats: 1?
• Observation time: 2.1 hours
• Limb active region target
best if present. Cannot
participate if Metis is
observing
Lines for SPICE Limb mode:
• H I 1025 Å,
• C III 977 Å,
• O VI 1032 Å,
• Ne VIII 770 Å,
• Mg IX 706 Å,
• Si XII 520 Å (2nd order)
– 3 profiles and 3 intensities.
Lines for SPICE Dynamics:
• 4 profiles and 6
intensities.
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SAP objective Target Duration
Opportunity
(e.g., orbital
requirements, solar
cycle phase,
quadrature ...)
Operational
constraints
Additional
comments
1.1.2.7 Trace streamer
blobs and other
structures through the
outer corona and the
heliosphere.
Corona +
Heliosphere
A few
Days
Quadrature with
Earth for combining
SO RS with L1 IS
and Earth
Coronagraphs with
SO IS
Metis
compatible
Quadrature or
radial
alignment with
SPP would be a
bonus
1.3.2 How is turbulent
energy dissipated and
how does turbulence
evolve within the
heliosphere?
Corona +
Heliosphere Statistics
Radial Alignment
with anything else.
Metis
compatible
Radial
alignment with
SPP would be a
bonus
1.1.2.5 Structure and
evolution of streamers
Corona +
Heliosphere Statistics
Quadrature with
Earth for combining
SO RS with L1 IS
and Earth
Coronagraphs with
SO IS
Metis
compatible
SPICE
participates if
possible.
2.3.1 Coronal shocks
Corona +
Heliosphere Statistics
Near perihelion for
highest spatial
resolution and best
spatial coverage in
the corona
Earth side for radio
obs from ground and
magnetic field
models for help with
post facto analysis
Metis
compatible
Quadrature
with SPP would
be a bonus
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2.3.2.1 Understand
coronal conditions
under which the
shocks form and
determine the
interplanetary
conditions where they
evolve
Corona +
Heliosphere Statistics
Near perihelion for
highest spatial
resolution and best
spatial coverage in
the corona
Earth side for radio
obs from ground and
magnetic field
models for help with
post facto analysis
Metis
compatible
Quadrature
with SPP would
be a bonus
2.3.2.3 Study heating
and dissipation
mechanisms at shocks
with radial distance
Corona +
Heliosphere Statistics Close to perihelion
Metis
compatible
EMC quiet
Burst modes
most important
here
Metis needs to
see within 5RS
for Lyα
2.3.2.4 Identify
mechanisms that heat
the thermal solar wind
particle populations
near shocks and
determine their
energy partition
Corona +
Heliosphere Statistics
Good alignment
with SPP (Radial or
Quadrature) would
be beneficial
EMC quiet
Metis not
strictly needed
for this one but
could still
provide useful
context.
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5.2.11 L_SMALL_MRES_MCAD_Ballistic-connection
In this SOOP, the spacecraft points at the modelled ballistic connection point, so this involves
tracking.
Default SOOP duration: 3 days
Pointing requirements: offpointing
Triggers: only IS triggers enabled
Based on SOOP 4 b as planned during PlanningExercise2016_withResults.pptx
Instrument Mode Comment
EUI
FSI Synoptic mode (S) with FSI
HRI in EUV & LYA Coronal hole
modes (C)
FSI in synoptic
HRI in CH mode at cadence ~ 900s;
PHI
PHI/HRT synoptic program:
HRT in PHI science mode 2 at
900-s cadence
ideally regular flushes;
SoloHI SoloHI: Nominal synoptic
perihelion program Model as HI_SYN_NEAR
SPICE SPICE composition & dynamics
interleaved
Use observation called SPICE_WIND_CONNECT
in modelling.
STIX STIX Normal Mode
EPD Normal Mode + regular burst
mode Close mode until at least the end of RSW
MAG Normal Mode + regular burst
mode
RPW Normal Mode + regular burst
mode Burst Triggers Active, selective OK.
SWA Normal Mode + regular burst
mode
SAP
objective Target Duration
Opportunity
(e.g., orbital
requirements,
solar cycle
phase,
Operational
constraints Additional comments
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quadrature ...)
1.1.2.7 Trace
streamer
blobs and
other
structures
through the
outer corona
and the
heliosphere.
Near-
quadrature, so
that SoloHI
can image
Earth-directed
blobs
objective also addressed by
L_FULL_HRES_HCAD_Coronal_Dynamics
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5.2.12 L_SMALL_MRES_MCAD_Connection_Mosaic
This SOOP is designed to let the high-resolution RS observations cover a wider area than normal, in
particular the SPICE FOV. This SOOP is to be used in particular with a mosaic of pointings, though
it does not necessarily need to be. Alternatively it could be used when we point for a limited amount
of time (few hours typically) to the most likely connectivity point, within a mainly sun-disk-centred
observation.
Default SOOP duration: 3 hours
Pointing requirements: mosaic of offpointing
Triggers: only IS triggers enabled
See SOOP 4a defined for SOWG8 planning exercise.
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Instrument Mode Comment
SPICE
(leads)
SPICE Composition Mapping (30mins
per map)
• Slit: 4’’or 6” • Exposure time: 180 s
• number of X positions = 10
(inferred, to maintain 30 minute
cadence)
• Field of View: 1.2’ x 11’
o depending on step size
chosen – i.e., not the
same as the slit width –
this could cover a
larger region.
• Number of repetitions of the
study: 6 (one at each position)
• Observation time: 6.4 hours
Mosaic of 6 positions would take bit more than 3 hours
Lines for SPICE Composition Mapping (from 1.1.2.2
Does slow and intermediate solar wind originate from
coronal loops outside of coronal holes?) :
– 15 lines (2 profiles+ 13 intensities)
• Ne VIII 770 Å,
• Ne VIII 780 Å,
• Mg IX 706 Å,
• O II 718 Å,
• O IV 787 Å,
• O V 760 Å,
• O V 761 Å,
• O VI 1032 Å,
• Ne VI 999 Å,
• Ne VI 1010 Å,
• Mg VIII 772 Å,
• Mg VIII 782 Å,
• C III 977 Å,
• Fe III 1017 Å
• Si II 992 Å
N.B. This line list for SPICE Composition Mapping is
similar to that of
L_BOTH_HRES_LCAD_CH_Boundary_Expansion,
except for O VI 1037 (removed) and Si II 992 (added).
PHI PHI science mode 2 with HRT, 15 mins
cadence 2 PHI/HRT datasets per SPICE map
EUI
EUV & LYA Coronal hole modes (C)
with HRI, 10-15 mins cadence
FSI Synoptic mode (S) with FSI
2-3 HRI images per SPICE map
SoloHI
SoloHI Synoptic program
modelled as HI_SYN_NEAR
cadences adapted to applicable sun distance and TM
corridor
STIX STIX Normal Mode
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Metis
will generally be SAFE+door closed,
due to off-pointing mosaic.
IF far enough from Sun: METIS standard modes for context for
connection
EPD Normal Mode
MAG Normal Mode
RPW Detection Mode Burst Triggers Active
SWA Normal Mode
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SAP
objective Target Duration
Opportunity
(e.g., orbital
requirements, solar
cycle phase,
quadrature ...)
Operational
constraints
Additional
comments
1.1.2.7 Trace
streamer blobs
and other
structures
through the
outer corona
and the
heliosphere.
N-S mosaic
centred on
most likely
connection
point
3 hrs +
slews for
mosaic of
6 positions
(30mins
dwell
time)
• RS window -
quadrature
with Earth:
SoloHI
images Earth-
directed
blobs, Earth
images blobs
that will hit
Solar Orbiter
• Daily N-S
mosaic driven
by SPICE
composition
maps: 6
positions, 30
mins dwell time
at each position
• Pre-window
observations
with EUI and
Metis desired
for coronal
context
radial or
quadrature
alignment
with SPP is
a plus
1.1.2.2 Does
slow and
intermediate
solar wind
originate from
coronal loops
outside of
coronal
holes?
coronal
loops
outside of
coronal
holes
few days
• near
perihelion for
resolution &
better linkage
conditions
• different
phases of
solar cycle
• Mosaic to map
larger region
(e.g. around
AR)
• Modelling to
find best
candidate
source regions
radial
alignment
with SPP is
a plus
1.1.3.2.3 How
does the
heliospheric
magnetic field
disconnect
from the Sun?
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5.2.13 L_SMALL_HRES_HCAD_Fast_Wind
Relating observed coronal holes and boundaries to the fast solar wind measured in situ.
Default SOOP duration: 3 days
Pointing requirements:
Triggers: only IS triggers active (TBC)
Instrument Mode Comment
PHI PHI_HRT_NOM_0: HRT nominal
mode at 1-min cadence
Flush 1-2 hr period (725MB per hour),
coordinated with EUI.
Note that FDT, highest spatial resolution
(2Kx2K), cadence 6 hours preferred, but
certainly ≥ 1 per day may be useful for context
beforehand (Earth source as well?).
Model PHI flush with 1550MB after SOOP has
completed.
EUI
HRI: 1 – 2 hours, at 1-min cadence,
e.g. EUV & LYA Coronal hole
modes (C) (reduced cadence)
FSI: FSI Synoptic mode (S) (chosen
for model) or FSI Reference Synoptic
mode (R), at several-hour cadence,
both 174 and 304.
≥ 12 hours of much lower-cadence HRI data for
context
We downlink 1-2 hours long period at 1min
cadence (220MB per hour) based on modelling
& LL data from EUI. Coordinated with PHI.
Model EUI flush with 440MB for HRI. FSI
volume based on 2.04kbits/s during whole
SOOP, 3 days =>66Mbytes.
Metis
If polar CH is target, modes are (see
METIS standard modes):
WIND: ≥ 2 hours
MAGTOP: ≥ min. 2 hours; repeated
each day in available observing
windows
FLUCTS: ≥ 1 hour; at perihelion (see
METIS special modes)
Rest of the time Metis will need to
switch off for off-pointing (TBC)
For SAP v0, we can model the Metis
contribution in this SOOP as:
• 2 hours WIND per day
• 2 hours MAGTOP per day
• 1 hour FLUCTS, once in the SOOP (will
be scheduled at closest point)
• Rest of the time Metis off
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STIX STIX Normal Mode
SPICE
SPICE Composition Mapping: 1
repeat (3.2 hours) (low-latitude CH
only);
SPICE Dynamics: 14 x 4" slit + 8 x
2" slit = 3.72 hours per day
throughout RSW; FOV 8' x 11'
(maybe 4' x 11' at poles)
Possible 4-pixel binning in Y
Default FOV in X for both studies is 4'. Will
assume that these are just placed side-by-side
but with half the time on each position if the CH
is large enough to warrant this. Duration is
therefore unaffected.
Use observation called SPICE_FAST_WIND in
modelling.
SoloHI
Mode: synoptic + turbulence
Currently modelled as
HI_SYN_NEAR +
HI_TURB_PER split to get EID-A
rate
EPD normal+burst
MAG normal+burst
SWA normal+burst
RPW Detection Mode Triggers on
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SAP
objective Target
Duration
(milliseconds
)
Opportunity
(e.g., orbital
requirements, solar
cycle phase,
quadrature ...)
Operational constraints Additional
comments
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1.1.1.1
Low FIP
fast wind
origins
Sufficientl
y wide
spatial area
across a
CH to
cover
connection
to
spacecraft.
Several hours
of RS
observations
of coronal
source,
considering
time taken for
wind to arrive
at s/c
• Pre-
perihelion
(so, south?)
and
perihelion
parts of RSW
• low-latitude
(to compare
with the intermediate
speed wind observed at
Earth) • any phase of
the cycle, but
more likely
during the
rising phase
• Performed on
several orbits
(preferably;
PHI request
10 orbits)
• EMC quiet for in situ
observations 12 hours
after the RS
observations of the
source region
o can be EMC-
noisy during
RS
observations
Likely to
involve
pointing
away from
disc center
("DC"):
difficult to
see how
Metis can
participate at
those times.
However,
could involve
DC pointing,
especially
between
observations
of the polar
CHs, to look
at
propagating
fluctuations
in/near plane
of sky (but
not
connected to
s/c) at 1.5–
2.9º from
DC.
Doppler
dimming
measurement
s will not be
so useful in these cases,
since the fast
wind will not
have made it
out into the
Metis FOV
(need to check exact
distances)
from the
observed
"source
region".
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1.1.1.2
Origin of
the small-
scale X-ray
and UV
jets in polar
coronal
holes
Polar CH
with
sufficiently
wide
extension
to catch
multiple
jets.
as above
• High-latitude
perihelion
(preferred, for fastest
solar wind) • solar
minimum or
declining
phase
as above
Metis will
require
repoints to
DC after off-
center
observations.
1.1.3.2.3
How does
the
heliospheri
c magnetic
field
disconnect
from the
Sun?
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5.2.14 L_SMALL_HRES_HCAD_SlowWindConnection
Try to catch with Remote Sensing instruments the dynamics at an open-closed field boundary which
will then be crossed in situ. High resolution RS observations required to catch dynamics. This
SOOP will in general need specific target pointing or target tracking unless the SC is far enough
from the Sun to catch.
Default SOOP duration: 3 days
Pointing requirements: Off-pointing: specific target pointing or target tracking
Triggers: only IS triggers enabled TBC
See SOOP 2 defined for SOWG8 planning exercise.
Instrument Mode Comment
EPD close mode +
scheduled/triggered burst
MAG normal +
scheduled/triggered burst
RPW normal +
scheduled/triggered burst selective downlink useful
SWA normal +
scheduled/triggered burst
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SPICE
combination SPICE
Dynamics & SPICE
Composition Mapping
rasters
• Slit: 4" or 6"
• Exposure time:
180 s
• X positions: 128
• FoV: (8' / 12') x
11'
• Observation time:
6.4 hours
• Nº repeats: 1
Raster area should be optimized to make sure open-closed
field boundary is captured:
• At perihelion if possible, since highest possible
resolution is preferred.
Lines for SPICE Composition Mapping (from 1.1.2.2 Does
slow and intermediate solar wind originate from coronal
loops outside of coronal holes?) :
– 15 lines (2 profiles+ 13 intensities)
• Ne VIII 770 Å,
• Ne VIII 780 Å,
• Mg IX 706 Å,
• O II 718 Å,
• O IV 787 Å,
• O V 760 Å,
• O V 761 Å,
• O VI 1032 Å,
• Ne VI 999 Å,
• Ne VI 1010 Å,
• Mg VIII 772 Å,
• Mg VIII 782 Å,
• C III 977 Å,
• Fe III 1017 Å
• Si II 992 Å
N.B. This line list for SPICE Composition Mapping is
similar to that of
L_BOTH_HRES_LCAD_CH_Boundary_Expansion, except
for O VI 1037 (removed) and Si II 992 (added).
Lines for SPICE Dynamics not specified
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EUI
EUV & LYA Coronal
hole modes (C) (uses
HRI) at 1min cadence
FSI Synoptic mode (S)
(FSI) throughout
If TM is limited, varying priority schema can be used to
keep HRI data manageable
PHI
Regularly spaced HRT
data at medium to high
resolution
e.g. PHI_HRT_MODE_2
(600s cadence default)
PHI LL magnetograms needed throughout.
Note that 1 hour cadence can suffice for interchange
reconnection at high resolution.
Processing after the RSW: PHI could focus downlink on the
most interesting periods, as inferred from other LL data
Metis
MAGTOP or GLOBAL
(see METIS standard
modes) for context
and linkage of solar wind
source regions to SC
modelled as GLOBAL
with default settings
Only applicable if beyond ~0.5AU during target tracking
SoloHI Context via SoloHI
synoptic modes Model for now as HI_SYN_NEAR
STIX STX_NORMAL
not strictly needed for SOOP although context is
appreciated
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SAP
objective Target Duration
Opportunity
(e.g., orbital
requirements,
solar cycle phase,
quadrature ...)
Operational
constraints
Additiona
l
comments
1.1.4.1.1
Interchange
reconnection
between open
and closed
field lines and
its role in
slow wind
generation
(see Planning
exercise Jan
2016 -
SOOP2)
Open-Closed field line
boundaries (near
ballistic connection
point):
• CH boundaries
• AR edges close
to low-latitude
open field
• Intermediate
areas of quiet
Sun
Target tracking
~1 RSW
(10 days)
• to be
studied for
CHs in
different
locations
(high vs
low
latitudes)
• different
opportuniti
es along
the orbit:
high-
latitude
windows +
perihelion
• to be
studied in
different
solar cycle
phases
• Earth view
before the
observation
s would be
asset to use
modelling
to define
best target
• During
RSW
• pre-
window
synoptics
needed
for target
choice
• VSTP
updates
needed
for target
tracking
• EMC
quiet for
connectiv
ity
radial
alignment
with SPP
is a plus
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1.1.2.2 Does
slow and
intermediate
solar wind
originate
from coronal
loops outside
of coronal
holes?
coronal loops outside
of coronal holes few days
• near
perihelion
for
resolution
& better
linkage
conditions
• different
phases of
solar cycle
• it may be
worthwhi
le to map
around
the whole
AR to
have
higher
chance of
being
connected
• EMC
quiet for
connectiv
ity
• Raster
area
should be
optimized
to make
sure
open-
closed
field
boundar
y is
captured
radial
alignment
with SPP
is a plus
1.2.2.6 Study
fast plasma
flows from
the edges of
solar active
regions
discovered
with
Hinode/EIS
edges of solar active
regions - at most likely
ballistic connection
point
few
hours/days
• fast flows
require
high
cadence
observati
ons
(mainly
SPICE
and
HRI?)
radial
alignment
with SPP
is a plus
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1.1.2.7 Trace
streamer
blobs and
other
structures
through the
outer corona
and the
heliosphere.
• Near-
quadrature,
so that
SoloHI can
image
Earth-
directed
blobs
objective
also
addressed
by SOOP
y: Link-
outer-
corona-to-
heliospher
e
1.2.2.5
Magnetic
reconnection
in the
chromosphere
, the
transition
region and the
corona
4.4 Are there
separate
dynamo
processes
acting in the
Sun?
in particular:
4.3.1.
Compare the
distribution of
small-scale
fields at low
and high
latitudes
(source:
[SOL-PHI-
MPS-
MN5100-TN-
2])
Quiet Sun at various
latitudes.
several
days
low + high
latitude, to be
repeated along the
cycle
• during
RSW
• feature
tracking
Only PHI,
EUI/HRI
and
SPICE
really
necessary
for this
goal.
PHI/HRT
mode 2
with 5-10
min
cadence, 5
quantities
and no
binning
EUI/HRT
follows
cadence of
PHI
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1.1.3.2.3 How
does the
heliospheric
magnetic
field
disconnect
from the Sun?
To be discussed whether any of the following can be linked:
1.1.3.1 Full characterization of photospheric magnetic fields and find structures
1.1.3.3 What is the distribution of the open magnetic flux?
1.2.1.7 Detect and characterize waves in closed and open structures
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5.2.15 L_BOTH_LRES_MCAD_Pole-to-Pole
This SOOP is designed to be used as a whole- or half-orbit Synoptic campaign that scans the Sun
from one high latitude to the other (therefore mainly to be used later in the NMP at min inclination
of 15º, i.e. orbit 3 or later). Close to the higher latitude windows, we get a radial sweep at nearly
constant latitude for in-situ.
This SOOP resembles very much the L_FULL_LRES_MCAD_Coronal_Synoptic but this time
SPICE is necessary as well.
We could possibly schedule this SOOP after every Venus GAM.
Default SOOP duration: 10 days
Pointing requirements: mainly disk center
Triggers: only IS triggers active
Instrument Mode Comment
EUI FSI Synoptic mode (S) with FSI 10 min cadence
Metis
One of METIS standard modes to observe
large scale coronal structures: GLOBAL or
LT-CONFIG
+ CMEOBS whenever a CME is triggered
Model: LTCONFIG (cadence 20mins) and 3 x
1 hour CMEOBS
CME Watch On
could do a slight offpoint without
switching metis off, e.g. to poles
during RSW1+3
PHI
FDT in general though HRT could be
scheduled too depending on solar distance, and
e.g. for pole observations
PHI_FDT_MODE_3 / PHI_HRT_MODE_3
(TBC)
at higher latitudes, point to the poles
for polar magnetic field observations.
Can be modelled as
PHI_FDT_MODE_3 everywhere as
resources are same.
SoloHI SoloHI synoptic Operations
Currently modelled as HI_SYN_NEAR
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SPICE SPICE Composition Mapping raster followed
by multiple instances of SPICE CME Watch
SPICE scans many latitudes during the
SOOP: large swath around the sun
magnetic field spiral less curved at
higher latitudes:
SPICE naturally closer to connection
point
Use observation called
SPICE_CME_COMP in modelling.
STIX N/A
EPD Normal + Burst Mode
MAG Normal + Burst Mode
RPW Detection Mode Burst Triggers Active
SWA Normal + Burst Mode
SAP objective Target Duration
Opportunity
(e.g., orbital
requirements, solar
cycle phase,
quadrature ...)
Operational
constraints
Additional
comments
2.2.1 How do CMEs
contribute to the global
evolution of magnetic flux
in the heliosphere?
2.2.2 What is the role of
ICMEs in the Sun’s
magnetic cycle?
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5.2.16 L_BOTH_MRES_MCAD_Farside_Connection
In this SOOP, the spacecraft points at the modelled ballistic connection point, so this involves
tracking. This is different from slow connection because it is intended to be used when PHI full disk
imagery is needed to see behind the limb as viewed from Earth.
Default SOOP duration: 1 day
Pointing requirements: Off-pointing
Triggers: only IS triggers enabled
(Will be modelled with duration of 1 day so that it can be repeated as often as needed in a RSW.)
Instrument Mode Comment
EUI
FSI Synoptic mode (S) with FSI
HRI in EUV & LYA Coronal hole
modes (C)
HRI in CH mode at cadence ~ 800s;
FSI continuously synoptic mode
Metis N/A Door closed
PHI
PHI/FDT synoptic program:
FDT in PHI science mode 6 at low
cadence
at EID-A rate, ideally regular flushes;
Maybe higher spatial resolution than true
synoptics.
SoloHI
SoloHI: Nominal synoptic program
(currently modelled as HI_SYN_NEAR)
SPICE SPICE Composition Mapping & SPICE
Dynamics interleaved
Use observation called
SPICE_WIND_CONNECT in modelling.
(See SPICE Pseudo-observations for SOOPs)
STIX STIX Normal Mode
EPD Normal Mode + regular burst mode Close mode until at least the end of RSW
MAG Normal Mode + regular burst mode
RPW Normal Mode + regular burst mode Burst Triggers Active, selective OK.
SWA Normal Mode + regular burst mode
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SAP
objective Target Duration
Opportunity
(e.g., orbital
requirements,
solar cycle
phase,
quadrature
...)
Operational
constraints Additional comments
1.1.2.7 Trace
streamer
blobs and
other
structures
through the
outer corona
and the
heliosphere.
At least
for 8 hrs,
best
several
days.
Near-
quadrature, so
that SoloHI
can image Earth-directed
blobs
objective also addressed by
L_FULL_HRES_HCAD_Coronal_Dynamics
Add
connectivity
objectives
here
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5.2.17 L_BOTH_MRES_MCAD_Flare_SEPs
This SOOP is aimed at understanding SEP properties and dynamics in relation to Flare events. EUI
and STIX are leading this SOOP, while IS payload provides continuous observations. Synoptic
support from other full disk RS instruments. Disk center pointing preferred.
For most of the science objectives SPICE is good to have, but as the SEP events are rather rare, it is
not feasible to go hunting for these events and to point SPICE to any particular region. However,
since all other instruments are involved, SPICE will probably be observing anyway (SPICE best
mode TBC).
Default SOOP duration: 1 day
Pointing requirements: disk center preferred
Triggers: IS + metis trigger active
Instrument Mode Comment
EUI (leads)
FSI Synoptic mode (S),
EUV & LYA Active Region modes
(A) (triggered) with 1min cadence
(default = 1or2s)
Most of the objectives in this SOOP are too
exceptional to be hunted for with high-res high-
cadence observations, except when they can be
triggered by a STIX flag whenever the flare is in
the HRI FOV.
Trigger needed: EUI to download only 1 event
in HRI AR mode (~120kbits/s for 1 hour)
STIX
(leads) STIX Normal Mode
Metis
METIS standard
modes GLOBAL,METIS special
modes CMEOBS
CMEOBS starts after CME flag rise, min. obs
time 1 hr (high cadence, 1 min), data volume~
2.137 Gb
SoloHI SoloHI Shock and Synoptic modes
PHI PHI science mode 2 (FDT) FDT in general
SPICE SPICE Waves mode (is sit-and-stare)
SPICE is good to have, but no target hunting
will be performed in this SOOP, so SPICE will
be disk center pointed
MAG MAG Normal Mode
EPD EPD Normal Mode
RPW RPW Normal Mode and RPW Burst
Mode
SWA SWA Normal Mode
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SAP
objective Target Duration
Opportunity
(e.g., orbital
requirements,
solar cycle phase,
quadrature ...)
Operational
constraints Additional comments
3.1.1.6 What
causes SEPs'
spectral
breaks?
Corona +
Heliosphere
Statistics
Radial dependence
Multiple orbits
PHI, STIX and FSI
seem to be the most
valuable RS
instruments to address
this goal.
SoloHI + SPP would
be interesting
combination as well.
EUI/FSI 174, 304, 5
min cadence (FSI
default cadence is 10
minutes)
EUI/HRI 174 & Ly-
alpha if source is
connected to SolO
SPICE would be
helpful if by chance the
SEP source region is in
the HRI FOV AND
EUI/HRI (Lya alpha)
observed it -> not
feasible to go hunting
for these events (TBC
by SWT)
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3.1.2.4
Explore the
fact that only
some of the
hard X-ray
peaks are
related to
escaping
electrons,
while others
are not
Corona &
Heliosphere
Statistics
Perihelion Metis
compatible
EUI/HRI 1 min
cadence for 30 min
before and during X-
ray peak (EUI AR
mode triggered based
on STIX flare)
EUI/FSI 10 min
cadence
SoloHI & Metis
connectivity with
Probe.
STIX
RPW
PHI/FDT
3.1.2.5 X-ray
prompt
events
Corona &
Heliosphere
In-situ
observed
electron
spectrum and
hard X-ray
photon
spectrum
correlation
Statistics
Perihelion
TBC: EUI/FDT 10 min
cadence, EUI/HRI 1
min cadence for 30
min before and during
X-ray peak
EUI is required to see
details of the flare
region in order to
decide if electrons and
ions are accelerated by
different processes at
different times (but not
sure how this is done
in practice, we should
discuss what EUI
mode is needed,
synoptic as described
in the objective page
does not seem
adequate).
This question should
be asked to Sam
Krucker.
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3.1.2.6
Delayed
events
(between X-
ray peak and
electron
release time)
Corona &
Heliosphere
Explore the
origin of the
delay
It is about a
secondary faint
emission after
the main X-ray
burst that could
be possibly
detected for the
first time.
Statistics
Perihelion
Best when the solar
limb from SolO is
connected to Earth,
or other s/c.
Metis
compatible
TBC: EUI/FDT 2 min
cadence, EUI/HRT 1
min cadence for 30
min before and during
X-ray peak
EUI could be needed
for studying the flare
region. If there is a
CME shock (2nd case)
SoloHI & METIS
could be interesting.
3.1.2.8
Explore the
type III radio
bursts delays
Corona &
Heliosphere
1)Propagation
effects in the
interplanetary
medium.
2)Coronal
magnetic
restructuring in
the aftermath
of CMEs (PHI
needed)
Statistics
Perihelion for
Propagation effects
in the
interplanetary
medium
Track connected
region for many
days at different
distances from Sun
does not seem to be
needed, even the
value of high-res
data for this
objective seems
doubtful (or too
ambitious to plan
for). If it would
drive the SOOP,
this would mean
tracking connected
region for many
orbits, to enhance
the chance of
catching a type III
radio burst.
Metis
compatible
EUI/FDT 174, 304 5
min cadence (EUI
default cadence is 10
minutes)
EUI/HRI 174 & Ly-
alpha 2 min cadence
SPICE again, not to be
hunted for. If SPICE
sees it, that is better but
not driver for the
SOOP.
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3.1.3.2
Double-
power law
spectra
Corona +
Heliosphere Statistics
Metis
compatible
Metis & SoloHI for the
shock observations and
EUI/FSI+STIX for the
flare ones
PHI not needed.
3.2.2
Latitudinal
and
longitudinal
transport of
SEPs
Corona +
Heliosphere Statistics
Needs range of
latitudes but indeed
this SOOP needs to
be scheduled at
some high-lat
windows as well
Metis
compatible
Best addressed by
multi-viewpoint
statistical dataset, e.g.
SPP (really close to the
sun) + SolO + earth-
based RS data. -> need
many events, viewed
from different
viewpoints and
different distances
Multi-viewpoint
(stereo), thus multiple
SC or Earth assets
SolO: full disk imagery
+ Metis
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5.2.18 L_BOTH_HRES_LCAD_CH_Boundary_Expansion
TO BE REVIEWED
SOOP needs review and more elaboration + to be compared to other SOOP
Similar to L_SMALL_HRES_HCAD_SlowWindConnection, but specifically targeted at
overexpanded CH boundaries as slow wind sources, requiring a different PHI mode and SPICE
observations.
Default SOOP duration: 1 day
Pointing requirements: Off-pointing, combined with disk-center
Triggers: only IS triggers enabled
Instrument Mode Comment
EPD N/A
MAG normal + scheduled/triggered burst
RPW normal + scheduled/triggered burst
SWA normal + scheduled/triggered burst
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SPICE
SPICE Dynamics (on 1 day), 4–6 rasters,
depending on structure. (Modelling as 6
rasters = 1 hour)
3–4 subsequent days of regular
observations: SPICE Composition
Mapping interleaved with SPICE
Dynamics
Full north-south raster only needed for
extended coronal holes. Otherwise point
at part of boundary modelled to be
connection point.
• Slit: 4" or 6"
• Exposure time: 30 s
• X positions: 224 (4") / 160 (6")
• FoV: 16' x 11'
• Observation time: 1.9 / 1.3 hours
per study
• Nº repeats: depends on target
Raster area should be optimized to make sure
open-closed field boundary is captured:
• At perihelion 6 rasters. Full north-
south raster only to be done for
extended holes. Otherwise it is
enough to point at the boundaries, it
could be that only one raster is
enough.
• For stable structures: 3-4 days of
standard observations + 1 day of
mosaic.
• Highest possible resolution is
preferred.
Lines for SPICE Composition Mapping:
– 15 lines (2 profiles+ 13 intensities)
• Ne VIII 770 Å,
• Ne VIII 780 Å,
• Mg IX 706 Å,
• O II 718 Å,
• O IV 787 Å,
• O V 760 Å,
• O V 761 Å,
• O VI 1032 Å,
• O VI 1037 Å,
• Ne VI 999 Å,
• Ne VI 1010 Å,
• Mg VIII 772 Å,
• Mg VIII 782 Å,
• C III 977 Å,
• Fe III 1017 Å
Lines for SPICE Dynamics:
– 4 profiles + 6 (2?) intensities:
• H I 1025 Å,
• C III 977 Å,
• O VI 1032 Å,
• Ne VIII 770 Å,
• Mg IX 706 Å,
• Si XII 520 Å (2nd order)
Estimate for daily mosaic with 4 pointing
positions of full rasters (6" slit, 10s exptime,
few strong lines, 30mins duration): 10 x 4 x
9.84Mb = 0.4Gbits (18% of orbital vol)
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EUI
FSI Synoptic mode (S) (FSI174)
Cadence of a few hours around s/c
magnetic connection point
Synoptic observations with FSI174 (perhaps
partial frame)
PHI
Regularly spaced FDT images for
I_cont, B_LOS, gamma and phi.
Cadence of 6 hours
PHI science mode 6 (with 6hrs cadence)
Throughout RSW.
Metis
Global maps of v_outflow in corona
Global n_e maps
Brightness fluctuation spectra
• WIND ≥ 2 hr, at 5 min cadence
• MAGTOP ≥ 2 hrs, twice a day, at
5 min cadence, during perihelion
RSW
• FLUCTS ≥ 1 hr, at 10s cadence,
VL intensity only
SoloHI
Measure solar wind speed above
potential source region in co-ordination
with Metis
STIX N/A
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SAP
objective Target Duration
Opportunity
(e.g., orbital requirements, solar
cycle phase, quadrature ...)
Operational
constraints
Additional
comments
1.1.2.1 Does
slow wind
originate
from the
over-
expanded
edges of
coronal
holes?
Coronal
hole(s),
Coronal
hole
boundary.
few hours
per day
• South RSW
o High latitude
observing is ideal
for polar coronal
holes.
• Perihelion RSW
o good if coronal
hole has equatorial
extension.
• Earth view beneficial
before RS window for
magnetograms to estimate
global solar magnetic field
SPICE
mosaic for
mapping.
Best if s/c is
on same
streamline as
SPP, but not
required
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5.2.19 R_FULL_LRES_HCAD_GlobalHelioseismology
From Obj 4 meeting
This SOOP was inspired on the discussion about Global Helioseismology needs during the
Objective 4 meeting in MPS, Göttingen (Oct 2016).
Default SOOP duration: 10 days (TBC)
Pointing requirements: Disk Centre
Triggers: none
SOOP designed for global helioseismology with PHI/FDT only. The PHI configuration is currently
designed to fulfil helioseismology needs for far side imaging - to detect modes passing through the
solar core - and for deep focusing. Note that the specified PHI mode is not (yet) part of the
predefined PHI modes.
Instrument Mode Comment
PHI
PHI FDT at 1 min cadence, processing
to v_LOS only
(model as PHI_FDT_SYNOPTIC_5
with
parameters CADENCE = 60
[s] COMPR = 8 IC = 0)
Images compressed by binning to 2x2 or
cropping when farther away from the Sun.
Higher compression may possibly be
acceptable for far side imaging.
SAP objective Target Duration
Opportunity
(e.g., orbital
requirements, solar
cycle phase,
quadrature ...)
Operational
constraints
Additional
comments
4.2.3 Probe the
structure in deep
layers of the Sun -
Deep focusing
Full
Disk
several
days (e.g. 3
days)
Earth (SDO/HMI) -
SC angle between 45
and 60 degrees
10-15Mm
resolution, i.e. 2x2
binning at perihelion
or cropping further
out
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4.2.3 Probe the
structure in deep layers
of the Sun - Far side
imaging to detect
modes passing through
the solar core.
Full
Disk
As long as
possible, 60
days would be
ideal but may
not be feasible
Earth (SDO/HMI)
observations in
combination with PHI
observations at far side:
angle 150 to 210
degrees.
Low resolution, v_LOS
only but 1min cadence.
Compression can be
quite high (TBC how
high).
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R_FULL_HRES_HCAD_Density_Fluctuations
Study of density fluctuations in the extended corona as a function of the outflow velocity of the
solar wind while evolving in the heliosphere. If SPICE participates, will need to have limb pointing
for a period of time, then return to disc center for the full-Sun instruments.
Default SOOP duration: 8 hours
Pointing requirements: may require offpointing (e.g. limb)
Triggers: only IS triggers enabled (TBC)
Instrument Mode Comment
EUI
FSI Synoptic mode (S)
Deep exposures
Long exposures needed to get good SNR
where it overlaps with Metis
Metis
(lead)
FLUCTS for 1 hour:
(20 mins of: 60 x 1s DIT + 10mins
processing x2
then 40mins at 20s cadence)
then MAGTOP (5 to 20 mins cadence)
Extra processing time is only needed at
1s cadence - 60 images can be stored
and queened for processing
primary at perihelion: METIS SNR for
1s cadence + near corotation (8hours)
(latitude changes not a big problem if
~1degree or less)
+ lower cadence during the rest of the
window: several days min to observe the
lower freq
duration: preferred to have it also 8h in
the other 2 RS windows
PHI
PHI science mode 6
6-hr cadence
1024 x 1024
Ic, B, γ, φ
for context
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SoloHI
(lead)
Contributes with synoptic+shock+turbulence
modelled as 50% HI_SHOCK_PER + 50%
HI_SYN_NEAR
Seems like this could be a little better
defined.
SPICE
(optional)
SPICE currently proposes to use either
SPICE Limb mode or SPICE Dynamics,
depending on the target.
Limb:
• Slit: 4"
• Exposure time: 5 s
• X positions: 224
• FoV: 15' x 11'
• Nº repeats: 10
• Observation time:
o 18 mins per study
o 3.2 hours total.
Active region (SPICE Dynamics):
• Slit: 4"
• Exposure time: 60 s
• X positions: 128
• FoV: 8.5' x 11'
• Nº repeats: 1?
• Observation time: 2.1 hours
• Limb active region target best if
present. Cannot participate if
Metis is observing, so not
modelled for now.
• If at limb: 224 positions of 60-s
exposures, so lasting 4 hours;
otherwise, the duration is as long
as Metis requires.
Lines for SPICE Limb mode:
• H I 1025 Å,
• C III 977 Å,
• O VI 1032 Å,
• Ne VIII 770 Å,
• Mg IX 706 Å,
• Si XII 520 Å (2nd order)
– 3 profiles and 3 intensities.
Lines for SPICE Dynamics:
• 4 profiles and 6 intensities.
SAP objective Target Duration
Opportunity
(e.g., orbital
requirements,
solar cycle phase,
quadrature ...)
Operational
constraints
Additional
comments
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1.1.2.4 Study of density
fluctuations in the
extended corona as a
function of the outflow
velocity of the solar
wind while evolving in
the heliosphere
Disc-center to
observe whole
corona (except
when SPICE
observes at
limb)
8 hours x
3
1 8-hour
observation in each
RSW of an orbit
Good to have SPP
in quadrature to
observe the
fluctuations in-situ.
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R_SMALL_MRES_MCAD_AR_LongTerm
Decay of Active regions: see slides in 4-Bellot-Rubio-SAP4_magnetoconvection5_final.pptx:
Decay process of ARs is not well known:
• Slow, may last a few weeks
• ARs approach limb and suffer from projection effects
• Emerging flux starts to reconnect with preexisting flux very soon
• Appearance of filaments, flux rope eruptions and CMEs in late phases of decay
• Sunspot fragmentation by light bridges?
• Flux erosion by convective flows?
• Role of moving magnetic features?
• How is the AR flux dispersed?
• What is the fate of the flux?
Default SOOP duration: 15 days
Pointing requirements: target pointing and tracking
Triggers: disabled
Instrument Mode Comment
EUI EUV & LYA Coronal hole
modes (C) 600 s cadence
EUV + Lyman alpha, FOV and resolution matching
PHI
SPICE SPICE Dynamics (bracketed by
SPICE Composition Mapping) 2" slit, 128 slit positions, scan duration 11 minutes
PHI PHI science mode 2 HRT Full FOV, 10 min cadence, 5 quantities, 2x2 binning
Metis METIS standard modes:
MAGTOP or GLOBAL
Metis can potentially provide context data before
and after the 15 day SOOP (offpointing during the
15 days)
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SAP objective Target Duration
Opportunity
(e.g., orbital
requirements,
solar cycle
phase,
quadrature ...)
Operational
constraints Additional comments
5.5.2.4
Isolated
AR,
complex
AR on E
limb
15 days
Perihelion: to
ensure near-co-
rotation,
Stereoscopy,
which
constrains
Earth-Sun-SC
angle
Duration 15
days
Potentially coordinated with
ground based
DKIST/EST/GREGOR/NST
for short-term studies
1.1.3.3 What is
the distribution
of the open
magnetic flux?
coronal
holes,
QS, AR
1.2.1.3
Contribution of
flare-like events
on all scales
increase cadence; only
partially
1.2.1.7 Detect
and characterize
waves in closed
and open
structures
plumes high latitude
pointing
1.2.2.5 Magnetic
reconnection in
the
chromosphere,
the transition
region and the
corona
increase cadence
(2.1.1.1 CME
initiation)
probably better
FDT
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4.1.1.1 Track
granules and
magnetic features
to follow their
motion and
interactions
increase cadence
4.1.3.2 Follow
individual
magnetic features
flux from lower
to high latitudes
4.1.3.4 How
supergranular
flows
facilitate/impede
transport of
magnetic
features?
increase cadence
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R_SMALL_HRES_LCAD_Composition_vs_Height
Mapping the abundance of minor ions as a function of height in the corona to distinguish between
slow and fast wind. This will be targeted at the boundary of a streamer, or at an active region on the
limb. SPICE-led.
Default SOOP duration: 3 hrs
Pointing requirements: off-pointing close-in / disk-center farther out
Triggers: disabled
Instrument Mode Comment
EUI
HRI observations, e.g. in EUV
& LYA Active Region modes
(A). Cadence TBD (currently
modelled with 360s cadence,
i.e. 20.2 kbps)
Context and also higher (than SPICE) cadence
observations in order to interpret the SPICE
composition map
Metis
One of METIS standard modes
before main observations, get
coronal context, but will not
participate at limb pointing
(unless s/c at large heliocentric
distance): mode TBD
Currently modelled as WIND
Context especially important for streamers
PHI
FDT synoptic data:
1024 x 1024 pixels, at 6 hr
cadence
Ic, B, γ, φ
PHI science mode 6
PHI is needed for context magnetic field, but mainly
before or after the SPICE observations as the target
will be on the limb!
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SPICE
(lead)
SPICE Composition Mapping
• Slit: 6"
o At larger
distances, when
possible, use
narrow slit (4"),
to enhance
spatial
resolution.
• Exposure time: 180 s
• X positions: 60
• FoV: 6' (or 4') x 11'
• Observation time: 3
hours per repeat
• Nº repeats: depends on
target
Lines for SPICE Composition Mapping:
– 15 lines (2 profiles+ 13 intensities)
• Ne VIII 770 Å,
• Ne VIII 780 Å,
• Mg IX 706 Å,
• O II 718 Å,
• O IV 787 Å,
• O V 760 Å,
• O V 761 Å,
• O VI 1032 Å,
• O VI 1037 Å,
• Ne VI 999 Å,
• Ne VI 1010 Å,
• Mg VIII 772 Å,
• Mg VIII 782 Å,
• C III 977 Å,
• Fe III 1017 Å
( L_BOTH_HRES_LCAD_CH_Boundary_Expansion
assumed to be default.)
.
SAP objective Target Duration
Opportunity
(e.g., orbital
requirements, solar
cycle phase,
quadrature ...)
Operational
constraints
Additional
comments
1.1.2.3
Abundance of
minor ions as a
function of
height in the
corona as
indicator of slow
or fast wind
Limb active
region (incl
the edges),
or streamer
boundary at
limb
3.2 hours,
twice
Any RSW is okay, but
perihelion preferred
for the AR case, and
>0.55AU is preferred
for the streamer case
(requires Metis
compatibility)
put EUI in at least
twice the cadence
of SPICE to
interpret
composition map
(and possible
aliasing in time)
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R_SMALL_HRES_LCAD_FineScaleStructure
High resolution observations of ARs or other solar features to discover the finest scales. SOOP with
low cadence observations as we do not aim at analyzing dynamics here. Perihelion preferred.
Pointing requirements: may be run at disk center or offpointed (without metis)
Default SOOP duration: 12 hours (TBC)
Triggers: disabled
Instrument Mode Comment
EUI
HRI in highest resolution, e.g. in EUV & LYA
Quiet Sun modes (Q) or EUV & LYA Active
Region modes (A) - LOW CADENCE: 10mins
best at perihelion. Model AR mode,
10 mins cadence
PHI
HRT in highest resolution, e.g. PHI science
mode 4 with cadence 10mins
(or PHI science mode 0 with much lower
cadence)
PHI need depending on science goal
SPICE
one of the high res modes, depends on science
goal. Model as SPICE Dynamics (highest
resolution mode).
may be needed for waves and/or
temperature structure discrimination,
depending on science goal
Metis in one of METIS standard modes, e.g.
MAGTOP with 10mins cadence (=default)
may be part of this SOOP for off-limb
observations (close to perihelion
only) like plumes
SAP objective Target Duration
Opportunity
(e.g., orbital
requirements, solar cycle
phase, quadrature ...)
Operational
constraints
Additional
comments
1.2.1.6 Resolve the
geometry of fine
elemental loop
strands
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R_SMALL_HRES_MCAD_PolarObservations
From SOL-PHI-MPS-MN1500-TN-2
This SOOP is inspired by science goal 1 of 2.3 in the MPS document and is also consistent with the
second part of 2.1 (apart from the number of physical quantities)
Default SOOP duration: Several days
Pointing requirements: Poles
Triggers: disabled
SOOP design to address polar magnetic field objectives that don't necessarily rely on the highest
resolution and cadence PHI data, nor all five physical parameters that PHI can return. The rest of the
remote sensing observations provide supporting data.
Instrument Mode Comment
PHI PHI science mode 2 (HRT) Full FoV, 2-5 minute cadence, 3
quantities, no binning
EUI
FSI Synoptic mode (S)
EUV & LYA Coronal hole modes (C)
Keep defaults for now, could potentially
match PHI cadence
SPICE SPICE_fast_wind (SPICE Pseudo-
observations for SOOPs)
SoloHI Normal observing programme
Metis Off pointing so door closed
STIX Normal observations
SAP objective Target Duration
Opportunity
(e.g., orbital
requirements, solar
cycle phase,
quadrature ...)
Operational
constraints
Additional
comments
4.1 How is magnetic
flux transported to
and re-processed at
high solar latitudes?
High
Latitudes
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4.2 What are the
properties of the
magnetic field at
high solar latitudes?
Solar
Poles
repeated high
cadence bursts
of several days
duration
high latitudes,
median to high
resolutions
off pointing so
no Metis
4.4 Are there
separate dynamo
processes acting in
the Sun? (4.3.1)
R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure
SOOP designed to study fine structure in the photosphere, similar to an overall RS 'Burst mode' but
does not involve Metis, SoloHI and STIX.
Default SOOP duration: 1 hr
Offpointing requirements: may be run at disk-center
Triggers: disabled
Instrument Mode Comment
EUI
EUV & LYA Quiet Sun modes (Q) or EUV & LYA
Coronal hole modes (C)
Model as EUI_HRI_QS with flush volume 2000MB
HRI 1 - 30 s cadence, maybe
interleaved with 0.1s cadence if
TM allows. Lyα
PHI PHI science mode 0 (HRT) 1 min cadence (flush volume
725MB in 1 hr)
SPICE SPICE_WIND_CONNECT (pseudo mode including
SPICE Dynamics and SPICE Composition Mapping)
2" slit, many rasters over similar
FoV to EUI and PHI
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SAP objective Target Duration
Opportunity
(e.g., orbital
requirements,
solar cycle
phase,
quadrature
...)
Operational
constraints
Additional
comments
1.1.3.1 Full characterization of
photospheric magnetic fields and find
structures
Quiet Sun
Coronal
Holes
A few
minutes
to an
hour as often as
we can.
Perihelion for
quiet sun
Close High
Latitudes for
Coronal Holes
Reduced FoV
(1024x1024
for PHI,
equivalent for
others)
1.1.4.1.6 Photospheric reconnection
Quiet Sun
Coronal
Hole
1 hour
per target
Perihelion for
quiet sun
Close High
Latitudes for
Coronal Holes
Full FoV.
Composition
mode may be
useful for
SPICE here
too.
4.2.2 Basic properties of solar high-
latitude magnetic field structures
5.5 Additional Science Objectives of
PHI
5.5.2.5 (Study flux appearance modes
and interactions in QS)
[source: 4-Bellot-Rubio-
SAP4_magnetoconvection5_final.pptx]
Quiet Sun
(disk center,
E/W limb,
N/S polar
region)
3 h per
pointing
(to
secure
good
statistics)
Perihelion for
highest
resolution
First orbit &
later in
mission for
higher
latitudes
Coordination
with DKIST
and/or EST
[alternatively,
GREGOR]
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5.5 Additional Science Objectives of
PHI
5.5.2.1 What are the velocity and
magnetic vectors in the solar
photosphere?
[source: SOL-PHI-MPS-MN5100-TN-
2, proposed by Alex Feller]
Disk center
(DKIST/FSP
looking at
solar limb)
few
hours
Perihelion +
quadrature
with Earth
co-
observations
with
DKIST/FSP,
pointing at
limb
Coordination
with DKIST:
FSP
instrument
pointing at
limb
Quadrature
with Earth!
FSP is
designed to
carry out
highly
sensitive
Hanle
diagnostics of
solar limb
magnetic
structures.
These
measurements
can be
calibrated only
if
simultaneous
high resolution
photometric
and Zeeman
measurements
are carried out
from a much
less inclined
vantage point.
5.5 Additional Science Objectives of
PHI
5.5.3.1 Effect of granulation and
oscillations, i.e., interaction of modes
and convection (stereoscopic
helioseismology).
Sunspot few
hours
What is the
max heliocentric
distance that
will give the
required
spatial
resolution?
co-
observations
with NEO at
~30º.
Earth/near-
Earth assets
assumed to be
continuously observing
(HMI, GONG
network, etc.).
If not, then co-
ordination will
need to be
taken into
consideration.
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R_SMALL_HRES_HCAD_AR_Dynamics
Tracking a complex AR of the Sun, for AR dynamics or tracking of a region for initiation. Making
use of Solar Orbiter's higher resolution capabilities, so would be at the best resolution (no binnings).
The acceleration mechanism in solar flares, tremendously enhancing (up to factors of ten thousand)
rare elements like 3He and ultra-heavy nuclei, has been puzzling for almost 50 years (e.g., Mason
2007; Reames 2017). The goal of this SOOP is to examine in detail underlying photospheric sources
of these so-called 3He-rich solar energetic particles (SEPs).
Specifically, the SOOP addresses the following two problems (References at bottom of this page):
1) Does the magnetic flux emergence (cancelation) play a fundamental role in energetic particle
production and release from the Sun? What is the growth rate of ARs associated with 3He-rich SEPs
(e.g., rapid growth may imply high intensities/enrichments, shorter time to SEP production/release)?
Statistically, the 3He-rich SEP sources (regions to which we are connected from the Earth) are
located near the west limb (~W55). Due to a projection effects this science question cannot be
properly investigated with NEO based observations. Though STEREO-A was in the right position
having a direct view on 3He-rich SEP sources, it does not provide surface magnetic field data.
Approaching to the Sun the SO connecting point move towards the Sun-SO line, improving
magnetic field observations of the connected regions (Bucik et al. 2014; Chen et al. 2015).
2) Would we detect 3He-rich SEPs from frequent small emerging dipoles (although without
significant EUV flaring as speculated by Wang et al. 2006) at closer distances to Sun? The SO will
be able to detect the events with much smaller intensities than at 1 AU (at perihelion, a factor of ~
50 if applying a simple inverse cube scaling law) probably allowing their detection also during solar
minimum conditions. Thus, with SO we may see 3He-rich SEPs from new sources.
[source: SOL-PHI-MPS-MN5100-TN-2, description updated by Andreas Lagg (email 10 Apr '17)]
Default SOOP duration: 1 day
Pointing requirements: target tracking
Triggers: enabled
Instrument Mode Comment
EUI FSI Synoptic mode (S) and EUV & LYA
Active Region modes (A)
Will need triggers to manage TM.
Model flush volume as 2500MB (~1hr
event+FSI)
Metis May contribute if disk-center pointed
(not modelled for now)
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PHI
PHI_HRT_NOM_0
+ PHI_HRT_MODE_2 (3 quantities) for
context around 'event'
Model flush volume as 850MB (~1hr
event + rest in mode 2)
SPICE SPICE Composition Mapping & SPICE
CME Watch
Use observation called
SPICE_CME_COMP in modelling.
STIX STIX Normal Mode Triggers Active
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SAP objective Target Duration
Opportunity
(e.g., orbital
requirements,
solar cycle phase,
quadrature ...)
Operational
constraints
Additional
comments
2.1.1.1 CME
initiation
Active Region 2 days
Earth side (≤ 60°),
All phases of the
solar cycle,
perihelion
In RSW
RSW
Extension
needed for
target
selection
VSTP needed
for target
updated
Offpointing so
no Metis
2.1.1.2 CME
structure
Active Region 2 days
Earth side (≤ 60°),
All phases of the
solar cycle,
perihelion
In RSW
RSW
Extension
needed for
target
selection
VSTP needed
for target
updated
Offpointing so
no Metis
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3.1.2.1
Understand
energy release
and particle
acceleration
process
Active Region
Properties of
flare energy
release
Acceleration
relating to the
magnetic
reconnection
process
Statistics
(Limited
by EUI
internal
memory)
Perihelion
preferred.
Best when co-
temporal images
from Earth are
available
Target
tracking
EUI/HRI 1 min
cadence over
several hours for
flare with direct
connection to
SolO.
To be done
together with
I_DEFAULT
3.1 How and
where are
energetic
particles
accelerated at
the Sun?
[source: SOL-
PHI-MPS-
MN5100-TN-2]
ARs several
days perihelion
co-
observations
from Earth
5.5 Additional
Science
Objectives of
PHI:
5.5.2.5 How do
magnetic fields
emerge on the
solar surface?
[source: SOL-
PHI-MPS-
MN5100-TN-2]
emerging flux
regions
PHI cadence may
need to be higher
in QS (2-5
minutes) than for
ARs (5-10 mins)
References
Bucik R., Innes D.E., Mall, U. et al. 2014, ApJ 786, 71
Chen N.-H., Bucik R., Innes D.E., Mason, G.M. 2015, A&A 580, A16
Mason G.M. 2007, SSRv 130, 231
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Reames D.V. 2017, Solar Energetic Particles, Lecture Notes in physics 932, Springer
Wang Y.-M., Pick M., Mason G.M. 2006, ApJ 639, 495
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R_SMALL_HRES_HCAD_PDF_Mosaic
Looking at the PDF of the magnetic elements. Idea is to scan the solar radius, which is almost 1º,
with a mosaic made up of different positions (3 or 4) from equator to pole.
Inspired on presentation 3-Lagg_polar_magnetic_fields.pdf
Default SOOP duration: ~2 hours
Pointing requirements: mosaic made up of 3 or 4 positions down central meridian
Triggers: disabled
Action on SOC: Compare (and possibly merge) with
R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure.
(Has similarities with L_SMALL_MRES_MCAD_Connection_Mosaic, too.)
Instrument Mode Comment
PHI PHI science mode 0 (HRT)
1 minute cadence. Mosaic down the central
meridian, 3 dwell positions, or 4 for some overlap.
At least 10 images per dwell
Overlap in pointings between the measurements to
do.
SPICE
Quick version of SPICE
Composition Mapping (similar to
SPICE mosaic SOOP)
SPICE would require 30 mins at each dwell
position
EUI To be filled in by EUI team - HRI
STIX
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SAP objective Target Duration
Opportunity
(e.g., orbital requirements,
solar cycle phase,
quadrature ...)
Operational
constraints
Additional
comments
4.1.2 Study the subtle
cancellation effects
that lead to the
reversal of the
dominant polarity at
the poles
4.1.4 Study the
influence of
cancellations at all
heights in the
atmosphere
4.2.1 Probability
density function
(PDF) of solar high-
latitude magnetic
field structures.
March or September, to see
the pole that is seen best from
Earth. For the same reason, it
should be able to see the same
target as Earth. High-latitude,
as close as problem.
4.2.2 Basic properties
of solar high-latitude
magnetic field
structures
4.4 Are there separate
dynamo processes
acting in the Sun?
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R_SMALL_HRES_HCAD_RSburst
This SOOP describes a coordinated observation of all high resolution RS instruments, running at
highest resolution and variable but high cadence, for a short period of time.
As a planning scenario, we propose to run this SOOP at every perihelion window where we have
some extra TM to spare, or where the campaign would fit without sacrificing too much of the rest of
the orbit.
This SOOP can be run for different targets, also at plain disk center, as it is aimed to discover new
physics and compare high cadence dynamics in all kinds of solar regions.
Default SOOP duration: 10 mins (can be repeated several times when it fits)
Pointing requirements: may be run at disk-center or off-pointed
Triggers: disabled
Instrument Mode Comment
EUI
HRI high res / high cadence modes,
depending on target:
EUV & LYA Quiet Sun modes (Q),
EUI/HRI Coronal hole mode (C) and
EUI/HRI Active Region mode (A))
EUI/HRI Discovery mode (D) to
"discover" periods <10s, min
observation time 300s.
(currently modelled with AR mode with
default cadence (1&2s), 5400kbps)
QS and AR mode, as defined now, generate at
4600-5400 kbps. (i.e. about 250x EIDA rate)
In most extreme case, i.e. Discovery mode,
HRI would generate 543MB per 10 mins
SOOP duration.
(data rate = 7240 kbps, i.e. about 350x EIDA
rate)
We download all generated data.
PHI PHI science mode 0 for HRT In mode 0, PHI generates at 1607 kbps (about
80x EIDA rate)
SPICE
Depending on science goal, but SPICE
Waves mode for highest resolution sit-
and-stare,
SPICE 30"-wide Movie if you want
more spatial information
In Waves mode, SPICE generates at ~50 kbps
(about 3x EIDA rate). Model as SPICE
Waves mode for SAP v0.
STIX STIX Normal Mode
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SAP
objective Target Duration
Opportunity
(e.g., orbital
requirements,
solar cycle
phase,
quadrature
...)
Operational
constraints Additional comments
1.2.1.1 Energy
flux in the lower
atmosphere
Bright source
(could be AR but not
necessary)
20mins to 3/5
hours
perihelion for
highest resolution,
multiple orbits
good to have
Better early
in mission for Ly alpha
degradation.
Duration for EUI observations depend on
mode: discovery mode may not run long but
can be combined with other mode like A or Q.
EUI/HRI-Lya preferred over HRI-EUV.
1.2.1.3
Contribution
of flare-like
events on all
scales
Flaring region
(could be
'quiet' sun for
nanoflares)
2-3 hours near-perihelion Need EPD/SIS as well
1.2.1.4
Observe and
explore flare-
like ‘heating
events’ from
the quiet
corona
Quiet Sun ?
perihelion
preferred but
not required
multiple
orbits good
to have
1.2.1.5
Determine
whether
coronal
heating is
spatially
localized or
uniform, and
time steady or
transient or
impulsive for a wide range
of magnetic
loops with
different
spatial scales.
AR, AR moss,
QS
several
hours
multiple
orbits good
to have
1.2.1.8
Investigate the
role of small
scale magnetic
flux
emergence in
energizing the
above laying
layers
several
hours
near-perihelion
(<0.5AU)
multiple
orbits good
to have
PHI/HRT leading, EUI/HRI required as
support
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1.2.2.3 What
are the origins
of waves,
turbulence and
small scale
structures?
to be combined with
L_FULL_HRES_HCAD_Coronal_Dynamics
(RS cadence TBC!):
• Long-term observations.
• Short duration bursts near perihelion.
1.3.1 Solar
and local
origin of
Alfvénic
fluctuations
spicules above
limb,
AR loops,
CH+boundary
1-2 hours perihelion
Better early
in mission
for Ly alpha
degradation.
to be combined with I_DEFAULT
1.2.1.7 Detect
and
characterize
waves in
closed and
open
structures
spicules at
limb, AR loops
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R_SMALL_HRES_HCAD_WaveStereoscopy
Default SOOP duration: 1 day
Pointing requirements: Off-pointing
Triggers: N/A
The scientific aim is to characterize the properties of waves in the photosphere and their coupling
with the atmosphere.
Waves are one clear mechanism for transferring energy from the photosphere to the chromosphere
and corona. Measuring the properties of the waves requires, in part, a determination of the velocity
field. The line-of-sight velocity component can be determined at different heights in the atmosphere
by observing Doppler shifts in different spectral lines. From the earth’s vantage point we have high
resolution ground based, balloon borne, and satellite instruments. Determining the horizontal
velocity has previously relied on using correlation tracking of intensity variations and rely on the
questionable assumption that the changes in location of the brightness fluctuations reflect the actual
velocity. The orbit and capability to measure Doppler velocities, in conjunction with existing and
upcoming ground-based or near-earth observatories, offers the unique chance to directly measure
two components of the velocity field using the Doppler effect.
High resolution co-temporal measurements including Doppler velocity maps from SO as well as
ground and NEOs are required. In particular, the ground-based and NEOs should include high
resolution Doppler images in the same line (with a higher cadence than that of SO), as well as lines
sampling different heights of the atmosphere. Co-observation with IRIS would be desirable. During
the observing period the earth-Sun-SO angle should be between 30◦ and 60◦ – a range which
represents a compromise between determining the two components of the velocity field and
allowing magnetic features which can act as wave guides to be partially resolved.
For ease of understanding the connection between the different heights, the observations would best
be performed at the center of the disk as observed from the earth (where observations over different
wavelengths are possible). Because also the achievable cadence will be higher on ground than with
SO/PHI, it is preferable to select targets which are closer to the disk center as seen from earth and at
higher heliocentric angles as seen from SO. The highest possible cadence is desirable, and a shorter
time series (of down to 30 minutes of Solar orbiter observations) would still allow the scientific
objectives to be met. (The ground based and NEO should be made for a period of 90 minutes
centered on the 30 minute SO observations). However, in order to guarantee reliable conditions
(seeing) at the coordinating ground-based observing facility (e.g. DKIST) a continuous high-
cadence observation period of several hours is required.
High resolution context magnetic maps from solar-orbiter immediately before and after the 30
minute observing window are required to provide context and aid co-alignment.
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A second observational campaign of an area 45◦ from disk center, with an earth-Sun-SO angle of
90◦, would be desirable.
Instrument Mode Comment
PHI PHI science mode 0 (HRT) full FoV 1 min cadence, 2 quantities plus magnetic
field context
EUI EUV & LYA Quiet Sun
modes (Q)
Full FoV 1-10s cadence (Just Lyman alpha?)
SAP objective Target Duration
Opportunity
(e.g., orbital
requirements,
solar cycle
phase,
quadrature ...)
Operational
constraints
Additional
comments
1.2.1 What mechanisms
heat the corona?
see
text
above
several
hours, best
30 minutes
downlinked
What is the max
heliocentric
distance that will
give the required
spatial
resolution?
co-
observations
with NEO at
30, 60, 90
degrees.
1.2.1.1 Energy flux in
the lower atmosphere
see
text
above
several
hours, best
30 minutes
downlinked
What is the max
heliocentric
distance that will
give the required
spatial
resolution?
co-
observations
with NEO at
30, 60, 90
degrees.
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5.5 Additional Science
Objectives of PHI
5.5.3.2 Two
components of velocity.
What is the relationship
between the components
of the velocities in
granulation?
Supergranulation?
Various modes in quiet
Sun? (stereoscopic
helioseismology)
Quiet
Sun
several
hours
What is the max
heliocentric
distance that will
give the required
spatial
resolution?
co-
observations
with NEO at
~30º.
Earth/near-Earth
assets assumed to
be continuously
observing (HMI,
GONG network,
etc.). If not, then
co-ordination will
need to be taken
into consideration.
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R_SMALL_HRES_HCAD_Ephemeral
From SOL-PHI-MPS-MN1500-TN-2
This covers half of SOOP 2.1 from the MPS Document, the other half is covered by
L_SMALL_HRES_HCAD_SlowWindConnection
Default SOOP duration: Several days
Pointing requirements: Quiet Sun
Triggers: disabled
The emergence, diffusion and decay of ephemeral regions near the poles and below high- latitude
coronal holes should be studied for the aspect of how they feed the magnetic network (see e.g.
Simon et al. 2001, ApJ 561, 49 427; Gosic et al. 2014, ApJ 797,). In particular, the latitudinal
dependence of this decay process would be interesting to study.
Instrument Mode Comment
PHI PHI science mode 3 (HRT) Half FoV 1-2 minute cadence, 5 quantities, no
binning
EUI EUV & LYA Coronal hole modes
(C)
FOV and cadence matching PHI, 15-30s cadence
SPICE SPICE Composition Mapping FOV Matching PHI
SAP objective Target Duration
Opportunity
(e.g., orbital
requirements, solar
cycle phase,
quadrature ...)
Operational
constraints
Additional
comments
4.1.1 Study the
detailed solar
surface flow
patterns in the polar
regions, including
coronal hole
boundaries.
- - see Helioseismology
SOOPs -
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4.4 Are there
separate dynamo
processes acting in
the Sun? (4.3.1.2)
Ephemeral
Regions, (quiet
Sun below the
poles, above
polar coronal
holes).
Several
Days
High latitude, solar
minimum
feature
tracking
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6 PLANNING STRATEGY
In order to construct the current version v0 of the SAP, we implemented the planning
strategy described below, by following, step by step, the different presented criteria.
We have to note that the SAP v0 does not include the planning of the Objective 4, even
though all the relevant SOOPs have been defined. However, many of its sub-objectives are already
covered by some of the planned SOOPs. Especially for the strategy of the Objective 4, see 4.0
Overall remarks and feasibility concerning Objective 4 observations with Solar Orbiter. The
remarks from that section will be integrated here at the next version.
Criterion 1: Best resolution RS data at different perihelia through the mission
We schedule an RS burst SOOP (R_SMALL_HRES_HCAD_RSburst) whenever we have a
perihelion with good telemetry downlink. We could aim at different types of targets or even plain
disk center to discover new physics in Solar Orbiter's highest cadence data, also possibly in
unexpected locations.
This campaign would serve several science goals that need very high cadence, need
perihelion and are aiming at different types of regions (see SOOP page). Even if off-pointing is not
possible, this campaign could still be useful to be run on the Sun-disk center region.
The SOOP under consideration is telemetry demanding (see telemetry estimates in SOOP
page R_SMALL_HRES_HCAD_RSburst, 3 to hundreds of times the EID-A rate), so we need good
telemetry at the time of the perihelion or right afterwards.
Alternatively, for some science objectives, e.g. 3.2.6 Effects of energetic particles
propagating downward in the chromosphere, it may be beneficial to schedule a few short RS bursts
into an RS window (or a series of RS windows), to enhance the chances of catching energetic
particles (instead of dedicating all telemetry for high-resolution high-cadence observations during a
few hours of the window).
Implementation of Criterion 1 for SAP v0:1
-> perihelion windows of
• MTP06 - 2021/07/01 - 2022/01/01,
• MTP08 - 2022/07/01 - 2023/01/01,
• MTP10 - 2023/07/01 - 2024/01/01,
• MTP14 - 2025/07/01 - 2026/01/01 (schedule at real perihelion, if you can hold onto the data
until the big underrun)
• MTP20 - 2028/07/01 - 2029/01/01 perihelion though this one if farther out (0.38AU)
1 For the definition and properties of the different MTPs for the October 2018 Option E, see next chapter.
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For objectives like 3.2.6 Effects of energetic particles propagating downward in the chromosphere,
we could possibly try to reserve the best telemetry windows for this case:
• MTP08 - 2022/07/01 - 2023/01/01 has 2 concatenated windows at good telemetry, can be
improved even more by shifting both perihelion and north window later, so that perihelion
RSW starts at perihelion.
• Alternatively, MTP14 - 2025/07/01 - 2026/01/01 can be used (similar but at higher
latitudes).
End of the mission may be more restrictive for RS bursts. It is peculiar in the sense that we get
many RSW per MTP but in general somewhat lower telemetry downlink:
- average rate for period 2026-2028 is 1.3 * EIDA rate
- average rate for period 2026-2029 is 1.64 * EIDA rate
Criterion 2: Objectives requiring Metis & SoloHI to observe Earth-directed Transients
The CME structure & propagation objectives as well as the blobs objectives ideally need
Solar Orbiter and Earth in quadrature with SoloHI looking towards Earth, so Solar Orbiter at GSE -
Y.
This criterion can preferably be applied at perihelia but also during high-latitude windows.
Alternatively, instead of quadrature, it will be interesting to observe at 45 degrees separation angle.
SOOPs that are most suitable to run during these times are:
1) L_FULL_HRES_HCAD_Coronal_Dynamics: focused on the off-limb corona up to Earth
2) L_FULL_HRES_HCAD_Eruption_Watch: same as above + the solar disk signatures including
PHI observing at higher cadence to see CME initiation
The second SOOP is more telemetry demanding but helpful if the CME happens to come towards
Solar Orbiter: then it can be viewed sideways from Earth.
Windows that fall close to Equinox could also be preferred since Earth-directed CMEs (and
southward IMF SW-Magnetosphere coupling in general) are more geoeffective (Russell and
McPherron, 1973).
Other SOOPs that should at least be run few times at quadrature
is L_FULL_HRES_MCAD_CME_SEPs and L_FULL_MRES_MCAD_Flare_SEPs (this one with
SoloHI towards Earth). Though these can be run at all times, some of its sub-objectives benefit from
quadrature with Earth, so that Earth can observe the structure of the CME heading towards Solar
Orbiter.
Implementation of Criterion 2 for SAP v0:
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This happens at the following perihelion windows:
• MTP05 - 2021/01/01 - 2021/07/01
• MTP07 - 2022/01/01 - 2022/07/01
• MTP09 - 2023/01/01 - 2023/07/01 may be less favorable for SoloHI (Thomson sphere) TBC
• MTP13 - 2025/01/01 - 2025/07/01 Perihelion+North window: high telemetry downlink
rate!! + higher latitude at declining phase may be ideal to study polar reversal.
Also, the following RS windows may be interesting, not at quadrature but at 45 degrees separation
angle:
• MTP11 - 2024/01/01 - 2024/07/01 last high latitude window, looking down on North pole at
21º (distance range 0.31-0.34AU)
• MTP 5,7,9 windows happen to fall close to Equinox.
Criterion 3: Slow solar wind connection science requiring Earth context for modelling pre-
RSW
We consider 2 very different types of connection science campaigns each requiring different
contributions from Earth and modelling:
1. During solar minimum, the magnetic field configuration is supposed to be quite simple,
with slow SW coming from the streamer belt. Also during the early orbits of the mission,
Solar Orbiter will stay close to the ecliptic.
If PHI observes the far side magnetic field in good resolution, and we combine that with the
Earth side magnetic field, the full solar magnetic field configuration can be modelled
including the location of the HCS that will determine the hemisphere Solar Orbiter will be
connected to.
This model could be the ideal starting point to do a longer term connection SOOP using
synoptic data of both IS and RS payload pointed to the most likely connection point.
During this campaign, PHI keeps on taking regular full disk magnetograms to update the
magnetic field model as we go. The modelling should also improve as Earth and Solar
Orbiter see overlapping longitude ranges on the Sun.
Proposed planning strategy:
o plan during solar minimum (i.e. early in the mission)
o start with PHI magnetogram data at far side (during one of the higher latitude
windows)
o take some time to construct the model
o use perihelion extension window to update the model and choose the RS target
o keep synoptic program during 10-20 days chasing the connection point
o SOOPs: L_FULL_HRES_LCAD_MagnFieldConfig for the magnetic field modelling
(during first RSW), during the connection observations we use
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L_BOTH_MRES_MCAD_Farside_Connection, possibly combined
with L_SMALL_MRES_MCAD_Connection_Mosaic
2. During the rest of the solar cycle, as the Sun becomes more active, also the magnetic field
modelling will become more challenging. In these periods, we hope to rely more on Earth
observations to get a well-constrained model of the field that Solar Orbiter is going to fly
through. If PHI data are restricted or not available, we mainly rely on Earth to produce the
model 4 days in advance due to VSTP turn-around loop. For this to happen, we need Solar
Orbiter in the GSE sector X<1 and Y<0, i.e. similar orbits than the ones needed for Earth-
directed transients above. The further Solar Orbiter moves away from that sector, the more
we rely on PHI data to model the most likely connection point.
Proposed planning strategy:
o start early but also plan later in the mission, towards solar maximum to explore other
types of solar wind source regions
o due to the more complicated and less reliable magnetic field modelling, we may want
to use pointing mosaics to establish the most likely connection
point: L_SMALL_MRES_MCAD_Connection_Mosaic
o during later orbits, the concatenated perihelion+North windows will span a large
range of latitudes over a short period of time (we are currently not sure this is an
asset or rather a complication to connection observation planning)
o SOOPs L_SMALL_MRES_MCAD_Connection_Mosaic combined
with L_SMALL_HRES_HCAD_SlowWindConnection (higher res/cadence RS to
explore source region in more detail) or L_SMALL_MRES_MCAD_Ballistic-
connection
If we point at a coronal hole boundary, also the
SOOP L_BOTH_HRES_LCAD_CH_Boundary_Expansion fits.
Implementation of Criterion 3 for SAP v0:
Case 1 (solar minimum):
Possible timing: MTP08 - 2022/07/01 - 2023/01/01: PHI observations during South Window
(combined with other science goals), connection science during perihelion+North window based
on magnetic field model.
(Note that MTP06 - 2021/07/01 - 2022/01/01 is a similar orbit but the first RSW needs to be shifted
due to the VGAM. We also do not find more opportunities later in the mission because the Sun will
be more active already and so solar minimum conditions will not be met)
Case 2 (rest of the solar cycle):
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Possible timing:
o (MTP05 we skip because we will not offpoint during the first orbit)
o MTP07 - 2022/01/01 - 2022/07/01 perihelion window
o MTP10 - 2023/07/01 - 2024/01/01 2nd South
o MTP11 - 2024/01/01 - 2024/07/01 Last North
o MTP13 - 2025/01/01 - 2025/07/01 (EMP) Perihelion
o MTP15 - 2026/01/01 - 2026/07/01 1st South and 2nd North
o MTP17 - 2027/01/01 - 2027/07/01 could be good if perihelion window gets shifted
earlier a few days: also, good for telemetry
o MTP19 - 2028/01/01 - 2028/07/01 Perihelion
o MTP21 - 2029/01/01 - 2029/07/01 South (+perihelion)
Criterion 4: Polar objectives
The different objectives that require high latitude have to be identified (mainly from Objective 4
that is not included in the current version). They have to be planned during high-latitude windows
and split between objectives that need good telemetry (for a high telemetry window) and those that
they don't.
Partial implementation of Criterion 4 for SAP v0:
MTP20-N + MTP21-S good opportunity for detailed pole analysis, telemetry very good for big
volume of PHI polar data. SOOPs to be added.
Criterion 5: Opportunities for long-term RS observations (concatenated windows or minimal
interruption)
Some science objectives benefit from a longer period of continuous RS observations, typically in
some sort of synoptic mode. A special case of this criterion is RS observations that could run from
pole to pole with minimal interruption.
Implementation of Criterion 5 for SAP v0:
MTP7 to MTP12 have naturally concatenated RS windows, i.e. 20 days continuous RS window +
possible 4-day extension. These MTPs are favorable
for L_FULL_LRES_MCAD_Coronal_Synoptic SOOP.
Also MTP13 - 2025/01/01 - 2025/07/01 (EMP) could be used as the 4-day extension window
could link 2 RS windows together and give the potential to keep on running (minimal) synoptic
observations.
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The following MTP periods are even better in the sense that RS observations can run from pole to
pole with minimal interruption. This is particularly interesting for L_BOTH_LRES_MCAD_Pole-
to-Pole SOOP.
• MTP11 - 2024/01/01 - 2024/07/01 is particularly interesting because it has 2 sets of
concatenated windows. In the 2nd orbit it has only 7 interwindow days between the South
window and the extension start of the concatenated period, so you can observe from 9 May
to 25 June with minimal interruption. You go from -21º to +21º passing through one of the
closest perihelia.
Telemetry is low in that MTP period which is OK since this SOOP has quite moderate
telemetry needs.
• The 2nd orbit in MTP15 - 2026/01/01 - 2026/07/01 (after VGAM) offers the possibility to
observe from 5 May to 19 June with min interruption (few days), flying from pole (31º) to
pole through a perihelion windows of 0.3AU!
• MTP16 - 2026/07/01 - 2027/01/01
• MTP17 - 2027/01/01 - 2027/07/01
• MTP18 - 2027/07/01 - 2028/01/01 - 2nd orbit
Criterion 6: Fast wind connection
L_SMALL_HRES_HCAD_Fast_Wind addresses 2 main science goals that in general need coronal
holes as a target.
Science objective 1.1.1.1 Low FIP fast wind origins would benefit from a low-latitude (or extended)
coronal hole, to increase the chance of connection and to compare the composition of low and fast
solar wind streams: this is most likely to happen in the declining phase of the solar cycle
(Hathaway: DOI 10.1007/lrsp-2015-4).
We prefer orbits/MTP with a fast scan through a big range of latitudes (like the opportunities above
for pole-to-pole SOOP).
In particular, in order to address science goal 1.1.1.2 Origin of the small-scale X-ray and UV jets in
polar coronal holes, high latitude windows are preferred to ensure the presence of a well-established
polar coronal hole that can be observed in full.
Limb pointing from medium latitude is also interesting to get the Doppler velocity component from
SPICE combined with EUI for off-limb intensity. Being up close is a big asset as well!
Implementation of Criterion 6 for SAP v0:
For fast wind origins (objective 1.1.1.1)
• Best opportunity: MTP15 - 2026/01/01 - 2026/07/01 (2nd orbit, after VGAM) offers the
possibility to observe from 5 May to 19 June with min interruption (few days), flying from
pole (31º) to pole through a perihelion windows of 0.3AU. This MTP has better telemetry
than MTP11, and this is during declining phase. Also, best quadrant for context from Earth.
• MTP13 - 2025/01/01 - 2025/07/01 is also possible, but not very good telemetry
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• MTP17 - 2027/01/01 - 2027/07/01
• MTP18 - 2027/07/01 - 2028/01/01 - 1st windows are very good telemetry-wise
(we did not add MTP16 because on far side, and did not add MTP11 because of low telemetry)
For the goal 1.1.1.2:
Closest high-latitude windows are:
• Any North window from MTP11 to MTP18 gives a good range of latitudes and stays within
0.4AU.
• MTP21 - 2029/01/01 - 2029/07/01 - south (up to 31º and 0.38AU)
Criterion 7: Science objectives needing perihelia but low telemetry requirements
L_FULL_MRES_MCAD_Flare_SEPs need medium telemetry downlink. Some of its sub-
objectives require quadrature with Earth, so this SOOP is also mentioned above in Criterion 2.
L_IS_STIX needs low telemetry (in practice this SOOP is likely to run throughout all RS windows).
L_SMALL_HRES_MCAD_Suprathermal_Popul needs perihelion for most of its sub-objectives,
and off-pointing to a target. telemetry needs are low. -> SOOP still needs review and clean-up. Not
yet scheduled in timeline (only 1 of its sub-objectives that need limb pointing).
As the telemetry needs of both SOOPs are moderate to low, we rather schedule at outbound
perihelia.
Implementation of Criterion 7 for SAP v0:
• MTP11 - 2024/01/01 - 2024/07/01- 1st orbit
• MTP12 - 2024/07/01 - 2025/01/01- End of perihelion with better telemetry rate
• MTP13 - 2025/01/01 - 2025/07/01 (EMP) works as well, also good for quadrature
• MTP15 - 2026/01/01 - 2026/07/01- 1st orbit
• MTP16 - 2026/07/01 - 2027/01/01
• MTP20 - 2028/07/01 - 2029/01/01
Criterion 8: Global magnetic field reconstruction & symmetry
RS windows at the far side of the Sun should be used to have regular, low cadence imaging of
magnetic field, to allow global field reconstruction.
This goal can be addressed by SOOP L_FULL_LRES_MCAD_Coronal_Synoptic
or L_FULL_HRES_LCAD_MagnFieldConfig.
Ideally, we plan this SOOP at regular far side windows covering a wide range of phases in the solar
cycle.
In addition, the same opportunities will satisfy allow L_FULL_LRES_MCAD_Coronal_Synoptic to
address the study of symmetry of the magnetic field and active longitudes (4.2 What are the
properties of the magnetic field at high solar latitudes?).
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Implementation of Criterion 8 for SAP v0:
At first sight, about 7 windows seem to cover these criteria (examples to be added).
Criterion 9: Rest of the objectives and special circumstances
1.1.2.3 Abundance of minor ions as a function of height in the corona as an indicator of slow
or fast wind
will be addressed through SOOP R_SMALL_HRES_LCAD_Composition_vs_Height:
We need limb pointing to address this science goal: either an AR (with open field at the edges) on
the limb or the boundary of a streamer, so the target can be chosen at the time of VSTP. We prefer
but do not require perihelion. Running this SOOP a bit further out could benefit from Metis
participation and enough signal in SPICE.
Exactly the same requirements are needed for 3.3.1.3 Role of shocks in generating SEPs.
This objective is addressed by SOOP L_SMALL_HRES_MCAD_Suprathermal_Popul and also
needs limb pointing at an RS window above 0.55AU, so that Metis can contribute to the
observations.
Implementation for SAP v0:
There seem to be plenty opportunities to plan this. Currently, this SOOP has been scheduled in
MTPs:
• MTP07 - 2022/01/01 - 2022/07/01
• MTP09 - 2023/01/01 - 2023/07/01
• more to be found
1.2.1.6 Resolve the geometry of fine elemental loop strands
will be addressed by SOOP R_SMALL_HRES_LCAD_FineScaleStructure.
This needs the highest possible resolution at close perihelia, but no high cadence and thus no
particularly high telemetry needs. All close perihelia seem to be candidates for this SOOP.
Implementation for SAP v0:
• MTP11 - 2024/01/01 - 2024/07/01 and MTP12 - 2024/07/01 - 2025/01/01 at 0.28AU
• MTP 7 to 10 and MTP15 to 18 at 0.3AU
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1.1.2.4 Study of density fluctuations in the extended corona as a function of the outflow
velocity of the solar wind while evolving in the heliosphere
will be addressed by SOOP R_FULL_HRES_HCAD_Density_Fluctuations. Metis and SoloHI are
leading this SOOP.
Telemetry constraints: The SOOP is telemetry limited (less than average) except for Metis that
needs more telemetry and SoloHI that seems to need its average allocation.
This SOOP needs to be repeated at several distances, i.e. each RS window, but not too far out for
Metis to still see the density fluctuations (distance limit to be added! 0.5AU TBC). 8 hour per
window should be enough.
Implementation for SAP v0:
The ideal orbits to tackle this science objective are the ones with 3 RSW that are quite close to the
Sun, e.g. between VGAM 7 and 8 all MTPs have their RS windows within 0.5AU and perihelia at
0.3AU:
• MTP15 - 2026/01/01 - 2026/07/01 (2nd half/orbit is best telemetry-wise)
• MTP16, MTP17, and MTP18
These MTPs all happen to fall in EMP. Another option, in the nominal mission, is to schedule this
SOOP once in every window of
• MTP10 - 2023/07/01 - 2024/01/01 if the option is chosen to move the 3 windows closer
together (mainly south window should be moved closer to make it suitable).
Photospheric dynamics (1.1.3.1 and 1.1.4.1.6)
addressed by SOOP R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure.
High telemetry needs for EUI, PHI and SPICE during short time (up to 1 hour). We need either
perihelion for quiet Sun or close-in high-latitude windows for coronal holes (i.e. North windows).
For perihelion windows, we can select the same ones as for the RS burst above.
Implementation for SAP v0:
For perihelion windows, we can select the same ones as for the RS burst above (MTP6, 8, 10, 14,
20).
Closest North windows, with reasonable latitude (>20º), are:
• MTP11 - 2024/01/01 - 2024/07/01 (especially start of 2nd north window seems perfect, also
telemetry-wise)
• MTP12 - 2024/07/01 - 2025/01/01 (not good telemetry-wise, so not added to the MTP page
for now)
AR dynamics
SOOP R_SMALL_HRES_HCAD_AR_Dynamics
The best opportunities to study CME initiation and structure (close to the Sun), are to point to ARs
at perihelion. We prefer Earth context for modelling and CME context.
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Implementation for SAP v0:
Earth-sided perihelion or close-in windows (which we interpret at <0.4AU) are:
• MTP13 - 2025/01/01 - 2025/07/01 - start of perihelion
• MTP15 - 2026/01/01 - 2026/07/01 - 2nd perihelion & 2nd north RSW
• MTP18 - 2027/07/01 - 2028/01/01 - first North
• MTP19 - 2028/01/01 - 2028/07/01 - perihelion
Opportunities at the edge of 60º angle with Earth are (not added to MTP pages for now):
• MTP07 - 2022/01/01 - 2022/07/01 - start of perihelion
• MTP17 - 2027/01/01 - 2027/07/01 - perihelion
Most of these opportunities are in EMP which seems OK as we will have more chances for CMEs
and dynamic events in general.
3.2.2 Latitudinal and longitudinal transport of SEPs
SOOP L_FULL_MRES_MCAD_Flare_SEPs (important to have SPP data).
It needs many events, viewed from different viewpoints (also separated from Earth) and different
distances. It also needs a range of latitudes (some high-latitude windows as well).
Ideally, it should be scheduled as many periods as possible.
Implementation for SAP v0:
The best opportunities are the following,
• MTP11 - 2024/01/01 - 2024/07/01- Both Perihelia + NW to have different viewpoints with a
range of latitudes and long-term duration (very good telemetry)
• MTP15 - 2026/01/01 - 2026/07/01- Both Perihelia + NW to have different viewpoints with a
range of latitudes and long-term duration (good telemetry)
• MTP18 - 2027/07/01 - 2028/01/01- Both Perihelia + NW to have different viewpoints with a
range of latitudes and long-term duration (bad telemetry in the second Perihelion)
Additionally: MTP14 - 2025/07/01 - 2026/01/01- Perihelion + NW for a range of latitudes and long-
term duration (good telemetry)
Energy flux in the lower atmosphere (1.2.1 What mechanisms heat the corona?)
This requires co-observations with Earth-based (DKIST) and NEO (IRIS) facilities, with sets of
particular geometries between Solar Orbiter, the target on the Sun, and Earth.
Implementation for SAP v0:
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Not added to corresponding MTP pages yet
Subject to the resolution requirements (i.e., minimum distance from the Sun) we suggest the
following opportunities:
• MTP05 - 2021/01/01 - 2021/07/01 - South
• MTP06 - 2021/07/01 - 2022/01/01 - North
• MTP07 - 2022/01/01 - 2022/07/01 - South
• MTP08 - 2022/07/01 - 2023/01/01 - North
• MTP10 - 2023/07/01 - 2024/01/01 - Perihelion + North
• MTP13 - 2025/01/01 - 2025/07/01 (EMP) - South
• MTP14 - 2025/07/01 - 2026/01/01 - North
• MTP15 - 2026/01/01 - 2026/07/01 - Perihelion #2
• MTP18 - 2027/07/01 - 2028/01/01 - Perihelion #1 + North #1
• MTP19 - 2028/01/01 - 2028/07/01 - South
• MTP20 - 2028/07/01 - 2029/01/01 – North
Limb stereoscopy of magnetic fields (5.5.2.1 What are the velocity and magnetic vectors in the
solar photosphere?)
This requires perihelion observations at quadrature so that Earth-based/-orbiting assets (specifically
the DKIST Fast Solar Polarimeter) can measure the magnetic field from an orthogonal
view. R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure seems like the best fit for this.
Implementation for SAP v0:
Not added to corresponding MTP pages yet
Objectives that could be enhanced with observations from the Parker Solar Probe
• L_FULL_MRES_MCAD_Flare_SEPs. In particular: 3.1.1.6 What causes SEPs' spectral
breaks? and 3.2.2 Latitudinal and longitudinal transport of SEPs
• L_IS_SoloHI_STIX. In particular 3.1.1.2.3 Warped shock fronts and 3.1.1.2.4 Turbulence
and inhomogeneities
• L_FULL_LRES_MCAD_ProbeQuadrature which requires SPP in quadrature with Solar
orbiter
• This list has to be completed.
Implementation for SAP v0:
The associated opportunities have not yet been defined as they need final trajectory of both Solar
Orbiter and Parker Solar Probe. However, the first two SOOPs are scheduled at other opportunities
already (based on the other sub-objectives).
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Objective 4 issues
We have to note that during the SAP v0 the Objective 4 has not been planned, even though many of
its sub-objectives are already covered by some of the planned SOOPs. Especially for the strategy of
the Objective 4, see 4.0 Overall remarks and feasibility concerning Objective 4 observations with
Solar Orbiter. The remarks from that section will be integrated here at the next version
Decay of ARs
•Study long-term behavior of active regions (5.5.2.4)
•Capabilities: near co-rotation, i.e. close to the Sun ideally perihelion and close to ecliptic,
SolO as observatory: we would use SOOP R_SMALL_MRES_MCAD_AR_LongTerm including
PHI, SPICE and EUI,
we also need stereoscopy: combine solar orbiter data with data from ground based observatories
•Target: Isolated active region, complex active region close to E limb
Implementation for SAP v0:
Possible target in option E:
• MTP07 - 2022/01/01 - 2022/07/01 perihelion + RSW3, but better move it earlier (i.e. close
to the Earth) to allow stereoscopy and increase telemetry
• extra repetition in the declining phase: MTP17 - 2027/01/01 - 2027/07/01 first 2
windows, but not in ecliptic (TBC!)
Probability distribution functions of the magnetic elements
These observations require a combination of high solar latitude with low heliocentric distance. In
the example trajectory, this is often in the third (North) RSW of an orbit. Coordinated observations
form Earth-based assets are needed, ideally at a large B0 angle, so close to equinox, particularly
September for the North Solar pole/ March for the South Solar Pole. This is challenging.
Implementation for SAP v0:
This is challenging, although there is one ideal opportunity in Option E.
SOOP R_SMALL_HRES_HCAD_PDF_Mosaic to be run possibly at following opportunities:
• MTP14 - 2025/07/01 - 2026/01/01 - RSW3
Polar observations
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• R_SMALL_HRES_HCAD_Ephemeral would be run on high-latitude patches of quiet Sun
or underneath high-latitude coronal holes. As this is looking for variations in the lifecycle of
ephemeral regions throughout the cycle (at high latitude), this needs to happen several times
throughout the cycle. This will certainly require off-pointing early in the nominal mission,
and maybe during the extended mission.
In addition, ...
Regular polar observations above 15-20 degrees latitude, ideally covering several phases in the solar
cycle.
Minimal duration of each campaign is 1 day at cadence 60s for v_LOS at highest resolution (in
HRT). Magnetic context once per hour is fine.
Total integrated length should be at least 30 days (more is better, especially for observing cycle
variations).
-> each SOOP using PHI in mode 0 can be used to address these goals (SOOPs to be added)
Implementation for SAP v0:
Not implemented yet.
Deep focusing
To be scheduled at times when Earth (SDO/HMI) - SC angle lies between 45 and 60 degrees.
Run R_FULL_LRES_HCAD_GlobalHelioseismology SOOP during several days (e.g. 3 days).
(FDT at 1 min cadence, only v_LOS. 10-15Mm resolution, i.e. 2x2 binning at perihelion or
cropping further out)
Far side imaging for modes passing through the solar core
Would require Earth (SDO/HMI) observations in combination with PHI observations at far side:
angle 150 to 210 degrees.
Run R_FULL_LRES_HCAD_GlobalHelioseismology as long as possible, 60 days would be ideal
but may not be feasible.
Compression of the images can be quite high (TBC how high).
SOOPs that should be running (quasi-)continuously
• I_DEFAULT -> currently scheduled to run throughout the whole mission. The in-situ
instruments will always contribute to this SOOP + may contribute to other SOOPs at the
same time (all SOOPs starting with 'L_').
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• L_IS_STIX -> in SAP v0 we scheduled STIX to run throughout all RS windows in its
default mode. The STIX data volume downloaded can be steered depending on the available
downlink at each time.
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7 OCTOBER 2018 OPTION E TRAJECTORY AND MEDIUM TERM PLANNING
7.1 Option E trajectory
Description of the trajectory from CREMA will be written here after the final launch date will be
known. It will subsequently be revised after launch for taking into account the actual orbital
characteristics.
7.2 Planning periods for Option E (MTPs)
The different planning periods of 6 months (called MTPs) for the nominal and extended mission
phases are the following:
NMP planning periods:
• MTP05 - 2021/01/01 - 2021/07/01
• MTP06 - 2021/07/01 - 2022/01/01
• MTP07 - 2022/01/01 - 2022/07/01
• MTP08 - 2022/07/01 - 2023/01/01
• MTP09 - 2023/01/01 - 2023/07/01
• MTP10 - 2023/07/01 - 2024/01/01
• MTP11 - 2024/01/01 - 2024/07/01
• MTP12 - 2024/07/01 - 2025/01/01
EMP planning periods:
• MTP13 - 2025/01/01 - 2025/07/01
• MTP14 - 2025/07/01 - 2026/01/01
• MTP15 - 2026/01/01 - 2026/07/01
• MTP16 - 2026/07/01 - 2027/01/01
• MTP17 - 2027/01/01 - 2027/07/01
• MTP18 - 2027/07/01 - 2028/01/01
• MTP19 - 2028/01/01 - 2028/07/01
• MTP20 - 2028/07/01 - 2029/01/01
• MTP21 - 2029/01/01 - 2029/07/01
In the next sections, we are detailing the characteristics as well as the preliminary science planning
for each one. The MTPs as shown here are based on a pure six-month division (although they may
be adapted where necessary to avoid that RSWs are split across MTPs). However actual MTPs will
follow the ESOC station-scheduling boundaries, which will only approximate this pure six-monthly
division, so some later adaptation of these boundaries is likely.
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7.2.1 MTP05 - 2021/01/01 - 2021/07/01
Plots are in GSE (geocentric solar ecliptic) coordinates, so Earth is at [0,0], the Sun is at [1,0]. The plot is the
projection of the orbit on the ecliptic plane.
First orbit in NMP.
7.2.1.1 RS window (default) placement
Period Window/GAM EXT
Start Start End Hcentric Distance Range Hgraphic Latitude Range
SC-
Sun- Earth Angle Conjunction
Safe Mode
Comms
Blackout
Largest
Comms
Roll
GAM
Restrictions
MTP
05
EGAM 1 N/A 2021-
01-01
MTP
05
South 2021-
02-13
2021-
02-17
2021-
02-27 0.71 0.64 0.59 -05.05 -05.12 -05.05 002 001 006 NO NO 00 NO
MTP
05
Perihelion 2021-
03-16
2021-
03-20
2021-
03-30 0.37 0.35 0.36 -00.85 02.12 04.10 059 083 106 NO NO 00 NO
MTP
05
North 2021-
03-30
2021-
04-02
2021-
04-12 0.37 0.43 0.49 04.59 05.12 04.89 119 136 148 NO NO 00 NO
7.2.1.2 Science planning
For the very first NMP orbit, we avoid triggered observations and run more synoptic-like programs.
• I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
• MTP5 - perihelion:
Can address objectives requiring Metis & SoloHI to observe Earth-directed Transients, as
Solar Orbiter is in quadrature with SoloHI looking towards Earth:
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o L_FULL_HRES_HCAD_Coronal_Dynamics running for 10 days
(L_FULL_HRES_HCAD_Eruption_Watch would be possible as well, but as it
involves triggers it does not seem ideal for first orbit)
o EUI/HRI and SPICE are not involved in the SOOP above. They can run individual
programs that fit in the TM corridors, and are valuable at disk center. Individual
SOOPs will be assigned later.
Current model: EUI_HRI_QS (EUV_CAD=600 LYA_CAD=600), SPICE_WAVES
Current model (SAP v0):
modelled within the baseline, currently only the perihelion window is filled, as defined above.
+ I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
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7.2.2 MTP06 - 2021/07/01 - 2022/01/01
Second NMP orbit, including 3 RS windows.
7.2.2.1 RS windows (default placement):
Period Window/GAM EXT
Start Start End Hcentric Distance Range Hgraphic Latitude Range
SC-
Sun- Earth Angle Conjunction
Safe Mode
Comms
Blackout
Largest
Comms
Roll
GAM
Restrictions
MTP
06
South 2021-
09-04
2021-
09-08
2021-
09-18 0.70 0.64 0.58 -05.07 -05.12 -05.03 161 164 169 NO
YES (RSW
Days 8 to 10) 00 NO
MTP
06
Perihelion 2021-
10-06
2021-
10-10
2021-
10-20 0.36 0.34 0.37 -00.50 02.47 04.30 134 110 087 NO NO 00 NO
MTP
06
North 2021-
10-20
2021-
10-23
2021-
11-02 0.38 0.44 0.49 04.73 05.11 04.82 075 059 047 NO NO 00 YES
MTP
06
VGAM 4 N/A 2021-
11-23
VGAM4 at 2021-11-23, so 3rd RSW need to moved out of 2021-10-23 to 2021-12-01 period, i.e.
shift 10 days earlier or move to completely different period.
7.2.2.2 New placement for the RS window:
• Keep MTP06 South and Perihelion
• MTP06 North can be moved without losing science as the max inclination of this orbit (<5º)
is negligible anyway. It can be moved either to:
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1. After the VGAM, to end 2021 or start 2022, i.e. the very end of MTP06 or start
MTP07. This option has extremely good TM but will be far from the Sun, e.g.
ending around 0.8AU.
2. Move between MTP06 south and perihelion windows. Depending on the science goal
you may want to concatenate it to any of those two, or move all 3 windows together
to cover 30 consecutive days of RS operations (may have drawback for calibration
opportunities though).
• Option 1 has the advantage that it can partially solve the problem of the huge underrun
expected for end 2021-start 2022. Also, the MTP06 perihelion window can be heavily
loaded with TM, stored on-board, and dumped once we have moved closer to Earth.
7.2.2.3 SOOP planning
I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
• MTP06-South: 10 days of Magnetic field synoptics to construct Carrington maps including
Earth data and Solar Orbiter data (still low latitude)
: L_FULL_HRES_LCAD_MagnFieldConfig (addressing objective 1.1.3.2.1 How does the
Sun's magnetic field change over time? & 1.1.3 Source regions of the heliospheric magnetic
field)
• MTP06-Perihelion: Short period (~30mins) of R_SMALL_HRES_HCAD_RSburst
• Alternatively, or combined with previous objective, the South and Perihelion window can
also be merged together and serve to address slow SW connection science goals, in
particular SOOP L_BOTH_MRES_MCAD_Farside_Connection, possibly combined
with L_SMALL_MRES_MCAD_Connection_Mosaic
Note that if you shift the South window later to improve TM and concatenate it with the
Perihelion window, one needs to take into account the conjunction period in between
(concatenated windows may end up being only just 20days long without precursor)
• MTP06-North: to be started asap after GAM, i.e. 30 Nov 2021 running until 9 Dec 2021
(implies NO precursors). Solar distance range of 0.77AU to 0.82AU. To be decided by
SWT whether this is worth pursuing (SoloHI and Metis restrictions).
Current model (SAP v0): RS windows at default locations still
• MTP06_South: L_FULL_HRES_LCAD_MagnFieldConfig throughout RSW
• MTP06_Peri: L_BOTH_MRES_MCAD_Farside_Connection for 9 days & 30
mins R_SMALL_HRES_HCAD_RSburst at end of window
• MTP06-North: L_SMALL_MRES_MCAD_Connection_Mosaic (3hrs/day) combined
with L_SMALL_HRES_HCAD_Fast_Wind rest of the time throughout RSW
(may need to be shortened due to TM restrictions)
• I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
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7.2.3 MTP07 - 2022/01/01 - 2022/07/01
3rd orbit in NMP.
7.2.3.1 RS window (default) placement
Period Window/GAM EXT
Start Start End Hcentric Distance Range Hgraphic Latitude Range
SC-
Sun- Earth Angle Conjunction
Safe Mode
Comms
Blackout
Largest
Comms
Roll
GAM
Restrictions
MTP
07
South 2022-
02-22
2022-
02-26
2022-
03-08 0.65 0.58 0.52 -12.9 -13.06 -12.77 030 026 020 NO NO 00 NO
MTP
07
Perihelion 2022-
03-22
2022-
03-26
2022-
04-05 0.32 0.30 0.32 -01.81 07.54 12.32 038 069 099 NO NO 00 NO
MTP
07
North N/A 2022-
04-05
2022-
04-15 0.31 0.38 0.44 12.32 12.94 11.53 099 122 138 NO NO 00 NO
7.2.3.2 Current scenario
This orbit has been exercised during the SOWG8 PlanningExercise2016_withResults.pptx.
I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
MTP7 - South:
Start of the RSW is beyond 0.55AU so we can offpoint with Metis still observing. This fits the
requirements of:
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• R_SMALL_HRES_LCAD_Composition_vs_Height (streamer case for which we need
Metis)
• L_SMALL_HRES_MCAD_Suprathermal_Popul for objective 3.3.1.3 Role of shocks in
generating SEPs.
We could choose several limb pointings and switch between them during few days.
MTP7 - perihelion:
• R_SMALL_HRES_HCAD_AR_Dynamics would fit at start of perihelion (good TM)
• Objectives requiring Metis & SoloHI to observe Earth-directed Transients, as Solar Orbiter
is in quadrature with SoloHI looking towards Earth:
o L_FULL_HRES_HCAD_Eruption_Watch, or alternatively:
o L_FULL_HRES_HCAD_Coronal_Dynamics (similar as above but less on-disk
activities)
• Combine with connection science goals on slow solar wind sources. These require modelling
to find the most likely connection point.
SOOPs L_SMALL_MRES_MCAD_Connection_Mosaic combined
with L_SMALL_HRES_HCAD_SlowWindConnection (higher res/cadence RS to explore
source region in more detail) or L_SMALL_MRES_MCAD_Ballistic-connection
If we happen to point at a coronal hole boundary, also the
SOOP L_SMALL_HRES_LCAD_CH_Boundary_Expansion fits.
MTP7 - Perihelion+ North: these 2 concatenated windows could also be used
for L_FULL_LRES_MCAD_Coronal_Synoptic SOOP
Modelled in SAP v0:
I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
MTP7 - South: first 2 days of limb pointing:
• 1 day R_SMALL_HRES_LCAD_Composition_vs_Height
• 1 day L_SMALL_HRES_MCAD_Suprathermal_Popul
MTP7 - perihelion: we prefer to model slow wind connection science over eruption watch because of the time in the
solar cycle (rising phase) which simplifies slow wind structure and lower chances for eruptions:
• L_SMALL_MRES_MCAD_Connection_Mosaic (~3hr/day) interleaved
with L_SMALL_HRES_HCAD_SlowWindConnection (~21hrs/day)
(can be modelled in 2 big blocks, i.e. ~1 day of first SOOP and 9 days of 2nd SOOP)
• if it turns out there is extra TM available, we
model R_SMALL_HRES_HCAD_AR_Dynamics on the 1st day of perihelion RSW
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MTP7 - north: L_FULL_LRES_MCAD_Coronal_Synoptic throughout RSW
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7.2.4 MTP08 - 2022/07/01 - 2023/01/01
4rd orbit in NMP.
7.2.4.1 RS window (default) placement
Period Window/GAM EXT
Start Start End Hcentric Distance Range Hgraphic Latitude Range
SC-
Sun- Earth Angle Conjunction
Safe Mode
Comms
Blackout
Largest
Comms
Roll
GAM
Restrictions
MTP
08
South 2022-
08-10
2022-
08-14
2022-
08-24 0.65 0.59 0.53 -12.86 -13.07 -12.83 167 171 178
YES (RSW
Days 9 & 10)
YES (RSW
Days 3 to
10)
00 NO
MTP 08
Perihelion 2022-09-06
2022-09-10
2022-09-20
0.32 0.30 0.32 -02.57 06.79 12.03 127 097 066 NO NO 00 NO
MTP
08
North N/A 2022-
09-20
2022-
09-30 0.32 0.37 0.43 12.03 13.00 11.72 066 042 026 NO NO 00 NO
7.2.4.2 Current scenario
I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
Connection science during solar minimum or rising phase, using all 3 RS windows + 1 instance of
the RS burst observations
MTP8 - South window
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• PHI magnetogram data at far side using L_FULL_HRES_LCAD_MagnFieldConfig for the
magnetic field modelling
Take some time to combine PHI data with Earth magnetogram data and construct the full solar
magnetic model.
MTP8 - Perihelion
Use perihelion extension window to update the model and choose the RS target:
• keep synoptic program during 10-20 days chasing the connection
point: L_BOTH_MRES_MCAD_Farside_Connection, possibly combined
with L_SMALL_MRES_MCAD_Connection_Mosaic (~3hrs per day)
• At the end of this perihelion window , we also plan a short period (~30mins)
of R_SMALL_HRES_HCAD_RSburst at the chosen target as this perihelion RSW has good
TM downlink which even improves when moving further towards Earth
MTP8 - North
• keep synoptic program chasing the connection
point: L_BOTH_MRES_MCAD_Farside_Connection, possibly combined
with L_SMALL_MRES_MCAD_Connection_Mosaic (~3hrs per day)
Alternative for MTP8 - Perihelion+ North:
• These windows have high TM downlink, and could be suitable to schedule a few instances
of R_SMALL_HRES_HCAD_RSburst to enhance chances of catching energetic particle
events (see 3.2.6 Effects of energetic particles propagating downward in the
chromosphere). TM return can even be improved by moving the 2 windows later,
starting at perihelion instead of before. In the case of starting at Perihelion itself, the
distance range covered during the two concatenated windows would be 0.3-0.5 AU.
• these 2 concatenated windows could also be used
for L_FULL_LRES_MCAD_Coronal_Synoptic SOOP
Modelled in SAP v0:
modelled exactly following the above assumptions in first 3 sections (south/peri/north) + I_DEFAULT outside of RS
windows, L_IS_STIX on inside all RS windows
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7.2.5 MTP09 - 2023/01/01 - 2023/07/01
5th orbit in NMP, 3 RS windows, reaching max latitude 13º.
RS window (default) placement
Period Window/GAM EXT
Start Start End Hcentric Distance Range Hgraphic Latitude Range
SC-
Sun- Earth Angle Conjunction
Safe Mode
Comms
Blackout
Largest
Comms
Roll
GAM
Restrictions
MTP 09
South 2023-01-25
2023-01-29
2023-02-08
0.65 0.58 0.52 -12.89 -13.06 -12.77 002 002 008 NO NO 00 NO
MTP
09
Perihelion 2023-
02-22
2023-
02-26
2023-
03-08 0.32 0.30 0.32 -01.88 07.47 12.29 066 097 127 NO NO 00 NO
MTP
09
North N/A 2023-
03-08
2023-
03-18 0.32 0.37 0.44 12.29 12.94 11.55 127 150 166 NO NO 00 NO
Current scenario
I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
• MTP 9 - South: Start of the RSW is beyond 0.55AU so we can offpoint with Metis still
observing. This fits the requirements of:
o R_SMALL_HRES_LCAD_Composition_vs_Height (streamer case for which we
need Metis)
o L_SMALL_HRES_MCAD_Suprathermal_Popul for objective 3.3.1.3 Role of
shocks in generating SEPs.
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o L_SMALL_MRES_MCAD_Connection_Mosaic interleaved
with L_SMALL_MRES_MCAD_Ballistic-connection
We could choose several limb pointings and switch between them during few days.
• MTP9 - perihelion: Objectives requiring Metis & SoloHI to observe Earth-directed
Transients, as Solar Orbiter is in quadrature with SoloHI looking towards Earth:
o L_FULL_HRES_HCAD_Eruption_Watch, or alternatively:
o L_FULL_HRES_HCAD_Coronal_Dynamics (similar as above but less on-disk
activities)
• MTP9 - Perihelion+ North: these 2 concatenated windows could also be used
for L_FULL_LRES_MCAD_Coronal_Synoptic SOOP (with current definition, this SOOP
includes L_FULL_HRES_MCAD_CME_SEPs, TBC on SoloHI's contribution)
Modelled in SAP v0:
I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
We are close to solar max so higher chance for eruptions:
MTP9 - South: 2 days of limb pointing (start RSW):
• 1 day R_SMALL_HRES_LCAD_Composition_vs_Height
• 1 day L_SMALL_HRES_MCAD_Suprathermal_Popul
• 1 day L_SMALL_MRES_MCAD_Connection_Mosaic
• 7 days L_SMALL_MRES_MCAD_Ballistic_Connection (in reality the mosaics and ballistic
connection are interleaved, however this is too much detail for modelling at this stage.)
MTP9 - perihelion:
• 1 day of L_FULL_HRES_HCAD_Eruption_Watch at start, and 1 repetition on day 5 -> prioritize 1or2 events in EUI and PHI buffers, to be flushed after 2nd SOOP instance. Metis keeps max 2 CMEs per day.
• Rest of the window (8 days) we schedule L_FULL_LRES_MCAD_Coronal_Synoptic
MTP9 - north: keep L_FULL_LRES_MCAD_Coronal_Synoptic for whole window
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7.2.6 MTP10 - 2023/07/01 - 2024/01/01
6th orbit & part of 7th orbit in NMP: 4 RS windows, first S and N windows reach max latitude 13º,
last S window reaches 21º already.
RS window (default) placement
Period Window/GAM EXT
Start Start End Hcentric Distance Range Hgraphic Latitude Range
SC-
Sun- Earth Angle Conjunction
Safe Mode
Comms
Blackout
Largest
Comms
Roll
GAM
Restrictions
MTP
10
South 2023-
07-13
2023-
07-17
2023
07-27 0.65 0.59 0.53 -12.86 -13.07 -12.83 166 162 155 NO NO 00 NO
MTP 10
Perihelion 2023-08-09
2023-08-13
2023-08-23
0.32 0.30 0.32 -02.64 06.71 12.00 100 070 038 NO NO 00 NO
MTP
10
North N/A 2023-
08-23
2023-
09-02 0.32 0.37 0.43 12.00 13.01 11.74 038 014 002 NO NO 00 YES
MTP
10
VGAM 5 N/A 2023-
09-28
MTP
10
South 2023-
12-11
2023-
12-15
2023-
12-25 0.58 0.51 0.45 -20.95 -21.32 -20.47 038 045 055 NO NO 00 NO
MTP10-North placement conflicts with VGAM5
RS window placement
• Move concatenated perihelion and North window earlier by ~5 days to avoid GAM. This
will slightly decrease the downlink rates for these windows, but as we are moving towards
Earth, SSMM can be filled up and all TM can be dumped right after the first 3 RS windows.
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• First South window can potentially be moved for better TM if we want to sacrifice the
moderate latitude (13º).
Potentially we could move the 3 RSWs together to make 1 concatenated window : this
could be helpful for L_FULL_LRES_MCAD_Coronal_Synoptic SOOP
and R_FULL_HRES_HCAD_Density_Fluctuations (which has a short duration ~8hrs)
• To check feasibility of this option, we should simulate what happens if
o we give more TM downlink share to IS payload during early MTP09
o let RS produce a lot of TM during MTP09 RS windows (and store on-board)
o lower the TM downlink share to IS payload during the bad comms period
overlapping MTP09 and MTP10 = April-June 2023 (EPD may be in far mode
anyway for last part of that period)
o let all payload produce at regular or better rates during the first 3 MTP10 RS
windows
• If we care more about latitude than solar distance during the 2nd south window, it could
potentially be moved earlier to gain TM. May help moving into MTP11 which has very bad
TM constraints.
Planning
I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
• MTP10-Full First Orbit (S + P + N): move windows together to make 1 concatenated
window and run L_FULL_LRES_MCAD_Coronal_Synoptic SOOP (see above)
o Extra idea for MTP10-Perihelion: Short period (~30mins)
of R_SMALL_HRES_HCAD_RSburst as this perihelion RSW has good TM
downlink which even improves when moving further towards Earth
• MTP10-2nd South: connection science goals on slow solar wind sources. These require
modelling to find the most likely connection point.
SOOPs L_SMALL_MRES_MCAD_Connection_Mosaic combined
with L_SMALL_HRES_HCAD_SlowWindConnection (higher res/cadence RS to explore
source region in more detail) or L_SMALL_MRES_MCAD_Ballistic-connection
If we happen to point at a coronal hole boundary, also the
SOOP L_SMALL_HRES_LCAD_CH_Boundary_Expansion fits: first time at 20º ->
preferred option for now (first opportunity at 20º
Modelled in SAP v0:
I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
• MTP10-Full First Orbit (S + P +
N): L_FULL_LRES_MCAD_Coronal_Synoptic throughout the 3 RS windows
+ during best TM of perihelion window 10mins of R_SMALL_HRES_HCAD_RSburst
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+ on 1st and last day of North window 8 hrs
of R_FULL_HRES_HCAD_Density_Fluctuations
• MTP10-2nd South: L_SMALL_HRES_LCAD_CH_Boundary_Expansion for 10 days
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7.2.7 MTP11 - 2024/01/01 - 2024/07/01
7th orbit & most of 8th orbit in NMP: 5 RS windows, S and N windows reach max latitude 21º.
RS window (default) placement
Period Window/GAM EXT
Start Start End Hcentric Distance Range Hgraphic Latitude Range
SC-
Sun- Earth Angle Conjunction
Safe Mode
Comms
Blackout
Largest
Comms
Roll
GAM
Restrictions
MTP 11
Perihelion 2024-01-03
2024-01-07
2024-01-17
0.30 0.28 0.31 -05.99 10.39 19.63 105 137 170 NO NO -01 NO
MTP
11
North N/A 2024-
01-17
2024-
01-27 0.31 0.36 0.43 19.63 21.20 18.96 170 164 148 NO NO 85 NO
MTP
11
South 2024-
05-09
2024-
05-13
2024-
05-23 0.59 0.52 0.46 -20.83 -21.33 -20.71 112 105 094 NO NO 00 NO
MTP
11
Perihelion 2024-
06-01
2024-
06-05
2024-
06-15 0.30 0.28 0.31 -05.50 10.91 19.82 042 010 023 NO NO 00 NO
MTP 11
North N/A 2024-06-15
2024-06-25
0.31 0.37 0.43 19.82 21.16 18.84 023 049 066 NO NO 00 NO
I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
MTP11 - First Perihelion window:
• one of the 3 closest perihelion windows, interesting for high-resolution science
like R_SMALL_HRES_LCAD_FineScaleStructure during few days
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• Metis, SoloHI can benefit from being close in as well:
run L_FULL_LRES_MCAD_Coronal_Synoptic for rest of the time
• (If it fits TM wise, L_FULL_MRES_MCAD_Flare_SEPs could be run here as well
(depending on what has been scheduled in rest of MTP) )
MTP11 - First North window: very close to the Sun at a reasonable latitude
• suitable for L_SMALL_HRES_HCAD_Fast_Wind
MTP11 - Second Perihelion window:
• one of the 3 closest perihelion windows, interesting for high-resolution science
like R_SMALL_HRES_LCAD_FineScaleStructure (can be run on different targets for
example)
• If it fits TM wise, L_FULL_MRES_MCAD_Flare_SEPs could be run here as well
(depending on what has been scheduled in rest of MTP)
• L_BOTH_MRES_MCAD_Flare_SEPs to run through this perihelion and upcoming north
window: benefits of close distance + radial alignment with Earth. SPP would be big asset as
well.
MTP11 - Second North window:
• Objectives requiring Metis & SoloHI to observe Earth-directed Transients, as Solar Orbiter
has 45 degrees separation angle with Earth.
We will be looking down on North pole at 21º (distance range 0.31-0.34AU), around solar
maximum.
SOOP for this RSW:
o L_FULL_HRES_HCAD_Eruption_Watch
o OR L_FULL_HRES_MCAD_CME_SEPs
• Possibly combined with connection science goals on slow solar wind sources. These require
modelling to find the most likely connection point.
SOOPs L_SMALL_MRES_MCAD_Connection_Mosaic combined
with L_SMALL_HRES_HCAD_SlowWindConnection (higher res/cadence RS to explore
source region in more detail) or L_SMALL_MRES_MCAD_Ballistic-connection
If we happen to point at a coronal hole boundary, also the
SOOP L_SMALL_HRES_LCAD_CH_Boundary_Expansion fits.
• The start of the final north window (or end final perihelion window) is also very close to the
Sun at a reasonable latitude which makes it ideal
for R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure targeted at a CH. Good
enough TM.
• L_FULL_MRES_MCAD_Flare_SEPs to be continued after perihelion window
• Also, suitable for L_SMALL_HRES_HCAD_Fast_Wind
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Alternative plan for full set of second orbit - final 3 RSWs:
• During this period, from 9 May to 25 June, we can run RS observations with minimal
interruption (only 7 interwindow days between the South window and the extension start of
the concatenated period).
SC goes from -21º to +21º passing through one of the closest perihelia. This is particularly
interesting for L_BOTH_LRES_MCAD_Pole-to-Pole SOOP.
TM is low in that MTP period which is OK since this SOOP has quite low TM needs.
Modelled in SAP v0:
I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
MTP11 - First Perihelion window:
• R_SMALL_HRES_LCAD_FineScaleStructure -> duration 2 days, with offpointing
• L_FULL_LRES_MCAD_Coronal_Synoptic -> duration 8 days
MTP11 - First North window: L_SMALL_HRES_HCAD_Fast_Wind duration 3 days - EUI and PHI flush limited
volume (still to be set in SOOP)
MTP11 - 2nd orbit (3 RSWs):
• L_BOTH_LRES_MCAD_Pole-to-Pole SOOP - 30 days in total
• 1 break for R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure at last day of
perihelion window - duration 1hr
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7.2.8 MTP12 - 2024/07/01 - 2025/01/01
end of 8th and 9th orbit in NMP: 3 RS windows, S and N windows reach max latitude 21º, best
perihelia: at 0.28AU.
RS window (default) placement
Period Window/GAM EXT
Start Start End Hcentric Distance Range Hgraphic Latitude Range
SC-
Sun- Earth Angle Conjunction
Safe Mode
Comms
Blackout
Largest
Comms
Roll
GAM
Restrictions
MTP
12
South 2024-
10-09
2024-
10-09
2024-
10-19 0.59 0.52 0.46 -20.86 -21.33 -20.65 103 110 119 NO NO 00 NO
MTP 12
Perihelion 2024-10-28
2024-11-01
2024-11-11
0.31 0.28 0.30 -07.40 08.79 18.97 168 161 127 YES (RSW Days 1 to 4)
YES (RSW Days 1 to 5)
93 NO
MTP
12
North N/A 2024-
11-11
2024-
11-21 0.30 0.36 0.42 18.97 21.30 19.31 127 101 083 NO NO -10 NO
MTP
12
VGAM 6
(EMP) N/A
2024-
12-21
I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
• MTP12 - South window:
• MTP12 - Perihelion window:
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o 1 of the 3 closest perihelion windows, interesting for high-resolution science
like R_SMALL_HRES_LCAD_FineScaleStructure
o L_FULL_HRES_LCAD_MagnFieldConfig as we are on the back side: will allow us
to construct 360º magnetic field + very low TM rate, suitable for this orbit
o OR L_FULL_MRES_MCAD_Flare_SEPs (requires medium TM so feasibility to be
modelled)
• MTP12 - North
o suitable for L_SMALL_HRES_HCAD_Fast_Wind
o continue L_FULL_HRES_LCAD_MagnFieldConfig
o could be combined with L_FULL_MRES_MCAD_Flare_SEPs
• MTP12 Perihelion+North: these concatenated windows could be used
for L_FULL_LRES_MCAD_Coronal_Synoptic SOOP
Modelled in SAP v0:
I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows
south window: /
perihelion: L_FULL_HRES_LCAD_MagnFieldConfig for 10 days
north : L_FULL_HRES_LCAD_MagnFieldConfig (6 days) + L_FULL_MRES_MCAD_Flare_SEPs for
4 days
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7.2.9 MTP13 - 2025/01/01 - 2025/07/01 (EMP)
1st orbit in EMP: 3 RS windows, S and N windows reach max latitude 28º.
RS window (default) placement
Period Window/GAM EXT
Start Start End Hcentric Distance Range Hgraphic Latitude Range
SC-
Sun- Earth Angle Conjunction
Safe Mode
Comms
Blackout
Largest
Comms
Roll
GAM
Restrictions
MTP
13
South 2025-
02-22
2025-
02-26
2025-
03-08 0.59 0.53 0.48 -27.64 -28.17 -27.17 039 032 023 NO NO -08 NO
MTP 13
Perihelion 2025-03-20
2025-03-24
2025-04-03
0.34 0.33 0.34 -06.12 10.84 22.44 028 051 075 NO NO 00 NO
MTP
13
North 2025-
04-03
2025-
04-06
2025-
04-16 0.35 0.41 0.46 25.35 28.14 26.16 090 110 126 NO NO 00 NO
• MTP13 Perihelion and North window seem ideal to do CME objectives or any transients
moving towards Earth:
o Solar Orbiter and Earth are in quadrature
o 2 RSW with very good TM downlink
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o very high latitude in North window (28º) during solar declining phase: chance of
CMEs affecting polar reversal!
o SOOP: L_FULL_HRES_HCAD_Eruption_Watch
o Also suitable for CME initiation, see
SOOP R_SMALL_HRES_HCAD_AR_Dynamics (close to the Sun at Earth side)
o Alternatively, these 2 concatenated windows could be used
for L_FULL_LRES_MCAD_Coronal_Synoptic SOOP
o Also L_FULL_MRES_MCAD_Flare_SEPs would fit as one of its objectives needs
quadrature + L_FULL_MRES_MCAD_CME_SEPs
• MTP13 Perihelion can also be used for connection science goals on slow solar wind
sources. These require modelling to find the most likely connection point.
SOOPs L_SMALL_MRES_MCAD_Connection_Mosaic combined
with L_SMALL_HRES_HCAD_SlowWindConnection (higher res/cadence RS to explore
source region in more detail) or L_SMALL_MRES_MCAD_Ballistic-connection
If we happen to point at a coronal hole boundary, also the
SOOP L_SMALL_HRES_LCAD_CH_Boundary_Expansion fits.
• MTP13 North also suitable for L_SMALL_HRES_HCAD_Fast_Wind
• MTP13 - full orbit also makes a fast scan through big range of latitudes, and can
alternatively be used for L_SMALL_HRES_HCAD_Fast_Wind to address 1.1.1.1 Low FIP
fast wind origins. But TM is not very good towards the end!
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7.2.10 MTP14 - 2025/07/01 - 2026/01/01
2nd orbit in EMP: 3 RS windows, S and N windows reach max latitude 28º, farther perihelia: at
0.33AU.
RS window (default) placement
Period Window/GAM EXT
Start Start End Hcentric Distance Range Hgraphic Latitude Range
SC-
Sun- Earth Angle Conjunction
Safe Mode
Comms
Blackout
Largest
Comms
Roll
GAM
Restrictions
MTP
14
South 2025-
07-22
2025-
07-26
2025-
08-05 0.59 0.53 0.48 -27.67 -28.16 -27.09 176 177 167 NO NO -01 NO
MTP
14
Perihelion 2025-
08-17
2025-
08-21
2025-
08-31 0.34 0.33 0.34 -05.61 11.39 22.78 115 092 068 NO NO 00 NO
MTP
14
North 2025-
08-31
2025-
09-03
2025-
09-13 0.36 0.41 0.46 25.58 28.12 26.04 053 033 018 NO NO 00 NO
• MTP14 - perihelion: another good opportunity to schedule 20-30mins
of R_SMALL_HRES_HCAD_RSburst at perihelion day, at least if you can hold onto the
data until the big underrun. Together with the North window this is a time with very high
TM downlink, and could be suitable to schedule a few instances
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of R_SMALL_HRES_HCAD_RSburst to enhance chances of catching energetic particle
events (see 3.2.6 Effects of energetic particles propagating downward in the chromosphere).
• MTP14 - NorthB also suitable for L_SMALL_HRES_HCAD_Fast_Wind (high latitude
and close in)
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7.2.11 MTP15 - 2026/01/01 - 2026/07/01
3th and 4th orbit in EMP: 6 RS windows, in the first orbit S and N windows reach max latitude 28º,
in the 2nd orbit they reach 31º already; perihelia at 0.33AU and 0.3AU.
RS window (default) placement
Period Window/GAM EXT
Start Start End Hcentric Distance Range Hgraphic Latitude Range
SC-
Sun- Earth Angle Conjunction
Safe Mode
Comms
Blackout
Largest
Comms
Roll
GAM
Restrictions
MTP
15
South 2025-
12-19
2025-
12-23
2026-
01-02 0.59 0.54 0.49 -27.51 -28.19 -27.39 028 035 045 NO NO 05 NO
MTP
15
Perihelion 2026-
01-14
2026-
01-18
2026-
01-28 0.34 0.33 0.34 -05.08 11.93 23.10 096 119 143 NO NO 00 NO
MTP
15
North 2026-
01-28
2026-
01-30
2026-
02-09 0.35 0.40 0.45 24.59 28.19 26.51 153 174 170 NO NO 82 NO
MTP
15
VGAM 7 N/A 2026-
03-15
MTP
15
South 2026-
05-05
2026-
05-09
2026-
05-19 0.50 0.43 0.38 -30.01 -31.29 -28.91 113 099 080 NO NO 00 NO
MTP
15
Perihelion 2026-
05-22
2026-
05-26
2026-
06-05 0.31 0.30 0.32 -15.16 06.26 22.12 047 021 006 NO NO 00 NO
MTP
15
North 2026-
06-05
2026-
06-09
2026-
06-19 0.34 0.41 0.46 27.83 31.30 29.59 026 047 061 NO NO 00 NO
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MTP15 - 1st South:
MTP15 - 1st Perihelion:
• L_FULL_MRES_MCAD_Flare_SEPs requires perihelion and medium TM. Close to solar
max so higher chance for flares.
MTP15 - 1st North: also, suitable for L_SMALL_HRES_HCAD_Fast_Wind (high latitude and
close in)
MTP15 - 2nd orbit:
• Use the last orbit (3 windows) for L_BOTH_LRES_MCAD_Pole-to-Pole SOOP.
• Alternatively, or in combination, use those for L_SMALL_HRES_HCAD_Fast_Wind! This
MTP has good TM to support it, it falls during declining phase with higher chance on
extended CHs. Also, best quadrant for context from Earth.
• Alternatively, or in combination, these 3 RSWs can also be favorable to
run R_FULL_HRES_HCAD_Density_Fluctuations with Metis leading (thus disk center
pointing!)
• Perihelion+North is also suitable for L_FULL_MRES_MCAD_Flare_SEPs: benefits of
close distance + radial alignment with Earth. SPP would be big asset as well.
MTP15 - 2nd perihelion & 2nd North:
• close to the Sun at Earth side, so suitable for R_SMALL_HRES_HCAD_AR_Dynamics
MTP15 - 2nd North: also, suitable for L_SMALL_HRES_HCAD_Fast_Wind (high latitude and
close in)
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7.2.12 MTP16 - 2026/07/01 - 2027/01/01
5th orbit in EMP: 3 RS windows, S and N windows reach max latitude 31º; perihelia at 0.3AU.
RS window (default) placement
Period Window/GAM EXT
Start Start End Hcentric Distance Range Hgraphic Latitude Range
SC-
Sun- Earth Angle Conjunction
Safe Mode
Comms
Blackout
Largest
Comms
Roll
GAM
Restrictions
MTP
16
South 2026-
09-16
2026-
09-20
2026-
09-30 0.50 0.43 0.38 -30.08 -31.27 -28.73 116 129 147 NO NO 00 NO
MTP 16
Perihelion 2026-10-03
2026-10-07
2026-10-17
0.32 0.30 0.31 -17.54 03.18 20.02 179 156 129 YES (RSW Days 1 to 4)
YES (RSW Days 1 to 4)
-17 NO
MTP
16
North 2026-
10-18
2026-
10-22
2026-
11-01 0.35 0.41 0.47 28.07 31.29 29.49 104 084 069 NO NO 00 NO
MTP16 - full orbit:
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• During this orbit, RS observations can run from pole to pole with minimal interruption. This
is particularly interesting for L_Pole-to-Pole SOOP.
• Alternatively, if RS windows can be linked
together: L_FULL_LRES_MCAD_Coronal_Synoptic (also requires long duration RS
observations)
• Other candidate: run R_FULL_HRES_HCAD_Density_Fluctuations at each RS window
(not too far out for Metis to still see the density fluctuations). 8 hour per window should be
enough. Metis is leading, thus disk-center pointing!
MTP16 - South:
MTP16 - Perihelion:
• L_FULL_MRES_MCAD_Flare_SEPs could be run at perihelion. Requires medium TM.
MTP16 - North:
• North window is one of closest high-lat windows, and thus suitable
for L_SMALL_HRES_HCAD_Fast_Wind to address 1.1.1.2 Origin of the small-scale X-ray
and UV jets in polar coronal holes
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7.2.13 MTP17 - 2027/01/01 - 2027/07/01
6th orbit + first part of 7th orbit in EMP: 4 RS windows, N and both S windows reach max latitude
31º; perihelion at 0.3AU.
RS window (default) placement
Period Window/GAM EXT
Start Start End Hcentric Distance Range Hgraphic Latitude Range
SC-Sun-
Earth
Angle
Conjunction
Safe Mode
Comms
Blackout
Comms
Roll
GAM
Restrictions
MTP
17 South
2027-
01-29
2027-
02-02
2027-
02-12 0.50 0.43 0.37 -30.16 -31.25 -28.54 NO NO NO
MTP
17 Perihelion
2027-
02-15
2027-
02-19
2027-
03-01 0.32 0.30 0.31 -17.03 03.86 20.50 NO NO NO
MTP
17 North
2027-
03-02
2027-
03-06
2027-
03-16 0.35 0.41 0.47 28.30 31.27 29.39 NO NO NO
MTP
17 South
2027-
06-13
2027-
06-17
2027-
06-27 0.50 0.44 0.38 -29.81 -31.31 -29.33 NO NO NO
MTP17-perihelion:
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• IF PERIHELION WINDOW GETS SHIFTED FEW DAYS EARLIER (also good for
TM):
connection science goals on slow solar wind sources. These require modelling to find the
most likely connection point.
SOOPs L_SMALL_MRES_MCAD_Connection_Mosaic combined
with L_SMALL_HRES_HCAD_SlowWindConnection (higher res/cadence RS to explore
source region in more detail) or L_SMALL_MRES_MCAD_Ballistic-connection
If we happen to point at a coronal hole boundary, also the
SOOP L_SMALL_HRES_LCAD_CH_Boundary_Expansion fits.
• close perihelion also candidate for R_SMALL_HRES_LCAD_FineScaleStructure
MTP17 - full orbit (first 3 windows):
• RS observations can run from pole to pole with minimal interruption. This is particularly
interesting for L_BOTH_LRES_MCAD_Pole-to-Pole SOOP.
• Alternatively, or in combination, run L_SMALL_HRES_HCAD_Fast_Wind to address
science goal 1.1.1.1 Low FIP fast wind origins: fast scan through big range of latitudes, and
low-latitude coronal hole would be ideal
• Other candidate: run R_FULL_HRES_HCAD_Density_Fluctuations at each RS window
(not too far out for Metis to still see the density fluctuations). 8 hour per window should be
enough.
MTP17-North = close high-lat window
• same SOOP as above L_SMALL_HRES_HCAD_Fast_Wind can be used to address science
goal 1.1.1.2 Origin of the small-scale X-ray and UV jets in polar coronal holes
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7.2.14 MTP18 - 2027/07/01 - 2028/01/01
second part of 7th orbit + 8th orbit in EMP: 5 RS windows, N and S windows reach max latitude
31º; 2 perihelia at 0.3AU.
RS window (default) placement
Period Window/GAM EXT
Start Start End Hcentric Distance Range Hgraphic Latitude Range
SC-Sun-
Earth
Angle
Conjunction
Safe Mode
Comms
Blackout
Comms
Roll
GAM
Restrictions
MTP
18 Perihelion
2027-
06-30
2027-
07-04
2027-
07-14 0.32 0.30 0.32 -16.52 04.54 20.97 NO NO NO
MTP 18
North 2027-07-15
2027-07-19
2027-07-29
0.35 0.41 0.47 28.53 31.25 29.28 NO NO NO
MTP
18 South
2027-
10-26
2027-
10-30
2027-
11-09 0.50 0.44 0.38 -29.89 -31.30 -29.17 NO NO NO
MTP
18 Perihelion
2027-
11-12
2027-
11-16
2027-
11-26 0.32 0.30 0.32 -15.99 05.21 21.43
YES (RSW Days
6 and 7)
YES (RSW
Days 6 and 7) NO
MTP
18 North
2027-
11-26
2027-
11-30
2027-
12-10 0.34 0.41 0.46 27.42 31.31 29.74 NO NO NO
MTP18 - full 8th orbit (i.e. RS windows 3, 4 and 5):
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• During this orbit, RS observations can run from pole to pole with minimal interruption. This
is particularly interesting for L_Pole-to-Pole SOOP.
• Alternatively, if RS windows can be linked
together: L_FULL_LRES_MCAD_Coronal_Synoptic (also requires long duration RS
observations)
• Other candidate: run R_FULL_HRES_HCAD_Density_Fluctuations at each RS window
(not too far out for Metis to still see the density fluctuations). 8 hour per window should be
enough. Metis is leading, thus disk-center pointing!
• Yet other alternative for this orbit with a fast scan through big range of
latitudes: L_SMALL_HRES_HCAD_Fast_Wind to address 1.1.1.1 Low FIP fast wind
origins
MTP18 - First North:
• North windows are one of closest high-lat windows, and thus suitable
for L_SMALL_HRES_HCAD_Fast_Wind to address 1.1.1.2 Origin of the small-scale X-ray
and UV jets in polar coronal holes
• 1st North window also suitable study CME initiation and structure (close to the Sun), with
pointing to ARs. Earth context preferred for modelling and CME
context: R_SMALL_HRES_HCAD_AR_Dynamics
• L_FULL_MRES_MCAD_Flare_SEPs : benefits of close distance + high latitude + radial
alignment with Earth. SPP would be big asset as well.
MTP18 - Second North:
• As above, also this North window is one of closest high-lat windows, and thus suitable
for L_SMALL_HRES_HCAD_Fast_Wind to address 1.1.1.2 Origin of the small-scale X-ray
and UV jets in polar coronal holes
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7.2.15 MTP19 - 2028/01/01 - 2028/07/01
9th orbit in EMP: 3 RS windows, N and S windows reach max latitude ~33º; relatively far
perihelion at 0.38AU.
RS window (default) placement
MTP19-Perihelion:
• connection science goals on slow solar wind sources. These require modelling to find the
most likely connection point.
SOOPs L_SMALL_MRES_MCAD_Connection_Mosaic combined
with L_SMALL_HRES_HCAD_SlowWindConnection (higher res/cadence RS to explore
source region in more detail) or L_SMALL_MRES_MCAD_Ballistic-connection
If we happen to point at a coronal hole boundary, also the
SOOP L_SMALL_HRES_LCAD_CH_Boundary_Expansion fits.
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• perihelion window also suitable to study CME initiation and structure (close to the Sun),
with pointing to ARs. Earth context preferred for modelling and CME
context: R_SMALL_HRES_HCAD_AR_Dynamics
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7.2.16 MTP20 - 2028/07/01 - 2029/01/01
10th orbit in EMP: 3 RS windows, N and S windows reach max latitude ~33º; relatively far
perihelion at 0.38AU.
RS window (default) placement
MTP20-perihelion:
• RS burst SOOP (R_SMALL_HRES_HCAD_RSburst). We could aim at different types of targets or even plain disk center to
discover new physics in Solar orbiter's highest cadence data, also possibly in unexpected locations. Note that this perihelion is relatively far out: at 0.38AU.
• L_FULL_MRES_MCAD_Flare_SEPs
MTP20-North:
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• good opportunity for detailed pole analysis, TM very good for big volume of PHI polar
data. SOOP to be added
7.2.17 MTP21 - 2029/01/01 - 2029/07/01
11th and last orbit in EMP: 3 RS windows, N and S windows reach max latitude ~33º; relatively
far perihelion at 0.38AU.
RS window (default) placement
MTP21-South(+Perihelion):
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• connection science goals on slow solar wind sources. These require modelling to find the
most likely connection point.
SOOPs L_SMALL_MRES_MCAD_Connection_Mosaic combined
with L_SMALL_HRES_HCAD_SlowWindConnection (higher res/cadence RS to explore
source region in more detail) or L_SMALL_MRES_MCAD_Ballistic-connection
If we happen to point at a coronal hole boundary, also the
SOOP L_SMALL_HRES_LCAD_CH_Boundary_Expansion fits.
Could be extended to perihelion window
• South window also good opportunity for detailed pole analysis, TM very good for big
volume of PHI polar data. SOOP to be added
• Close high-lat window, interesting for L_SMALL_HRES_HCAD_Fast_Wind, to address
science goal 1.1.1.2 Origin of the small-scale X-ray and UV jets in polar coronal holes
MTP21-Perihelion:
• see above: connection science goals on slow solar wind sources. These require modelling to
find the most likely connection point.
SOOPs L_SMALL_MRES_MCAD_Connection_Mosaic combined
with L_SMALL_HRES_HCAD_SlowWindConnection (higher res/cadence RS to explore
source region in more detail) or L_SMALL_MRES_MCAD_Ballistic-connection
If we happen to point at a coronal hole boundary, also the
SOOP L_SMALL_HRES_LCAD_CH_Boundary_Expansion fits.
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8 SIMULATIONS
This section will contain all simulations performed once the full mission timeline is known. This
will be done after the knowledge of the launch date and final orbit is stabilized. The simulations will
include the SSMM fill state, the fill state per instrument stores, power consumption etc.
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9 APPENDIX
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