Master Thesis

Imaging of the Earth’s Plasmaspherein the Extreme Ultraviolet
極端紫外光による地球プラズマ圏の撮像

Go MurakamiSpace and Planetary Science Group
Department of Earth & Planetary Science
Graduate School of Science, The University of Tokyo

February 28, 2008

Abstract
The Earth’s plasmasphere is the region where dense plasma originated from theionosphere is captured along the closed geomagnetic field lines. Intensive observations
using spacecrafts revealed the dynamical aspect of the plasmasphere changing
depending on geomagnetic conditions. The cycles of erosion and refilling of the plasma
population in the plasmasphere have been studied extensively in the past. However,conventional in-situ measurements are local observations with which it is difficult toseparate the temporal development and the spatial structure of the plasmasphere. To
overcome this difficulty, the visualization of the plasmasphere has recently beenemphasized. Helium ions are the second major component in the plasmasphere (almost
10% of the total amount), and resonantly scatter the solar EUV emission at 30.4 nm. For
this reason, detecting this emission leads us to the global imaging of the plasmasphere.
Recent advances in satellite-based imaging techniques have made it possible toroutinely obtain full global images of the plasmasphere. The Extreme Ultraviolet
Imager (EUV) on the IMAGE satellite gave us complete sequential pictures. I
statistically examined the plasmapause response to the southward turning of theinterplanetary magnetic field (IMF) using sequential EUV images of the plasmasphere.
By tracking the location of the plasmapause in sequential EUV images, I calculated theplasmapause radial velocity and then estimated the time development of the convection
electric field at the plasmapause. Consequently, my research indicates that theplasmapause response to the southward turning of the IMF takes 10-30 min, with anaverage of 18 min. This timescale suggests that the convection electric field propagatesfrom the magnetopause to the inner magnetosphere through the ionosphere.
In the near future, the Telescope of Extreme Ultraviolet (TEX) mounted in theUpper atmosphere and Plasma Imager (UPI) on the SELENE satellite will provide
global images of near-Earth plasmas from the lunar orbit. The SELENE satellite waslaunched and sent to the Moon in 2007. The UPI-TEX imager detects the resonancescattering emissions of helium ions (He II: 30.4 nm) and oxygen ions (O II: 83.4 nm) totake images of the plasmasphere and the polar wind, respectively. I calibrated theperformance of the UPI-TEX imager before the launch. As a result, the UPI-TEX has atotal sensitivity of 0.014 cps/Rayleigh/bin for the He II (30.4 nm) and 0.0036
cps/Rayleigh/bin for the O II (83.4 nm). Imaging of the plasmasphere with the 2-minexposure and the 0.09-RE resolution improves our understanding of the distribution ofplasma in the inner magnetosphere and the forces that control it. On the other hand, theremote sensing at O II with the 30-min exposure and the 0.5-RE resolution gives the firstglobal image of the polar wind and make significant contributions to our understandingof the transport routes and mechanisms of the cold O+ ions.
Furthermore, I have developed a high-resolution imaging detector with five
microchannel plates (MCPs) in a set of V- and Z-stacks and a resistive anode encoder(RAE) for the future missions. In a position-sensitive system with an RAE, the spatial
resolution depends on the signal-to-noise ratios at the anode terminals. Therefore, a highand stable electron gain of MCPs allows the position determination of eachphotoelectron event with a high spatial resolution. I investigated the effect of thenegative inter-stack potential on the PHD and the spatial resolution by means ofcalculation and experiments. As a result, applying the negative inter-stack potential
reduced the width of the PHD and improved the spatial resolution to 400×400 pixels.
The results enable us to optimize and apply the technique to the future missions.

要旨
地球近傍にはプラズマ圏と呼ばれる高密度の冷たいプラズマで満たされた領域が存在する. 地球近傍におけるプラズマの運動を支配するのは主に地球の自転が引き起こす共回転電場と太陽風が磁気圏内に引き起こす大規模な対流電場である. これらを重ね合わせた電場におけるプラズマの E B ドリフトの運動からプラズマ圏の形成を説明できる. 太陽風の変動に応じてプラズマ圏の構造が活発に変化することが過去の人工衛星による直接観測から明らかにされた. しかし直接観測では衛星のその場の情報しか得られないため, プラズマ圏の時間変化と空間変化を分離することは困難であった. こうした問題を解決するため,プラズマの可視化により地球近傍でのプラズマの振る舞いを大局的に把握しようとする試みがなされてきた. プラズマ圏を構成する主なイオンのうちヘリウ
ムイオンと酸素イオンは波長 30.4 nm と 83.4 nm においてそれぞれ共鳴散乱を起こすため, これらの光を捉えることでプラズマ圏全体の時間変動を明らかにすることができる.
2000 年に打ち上げられた IMAGE 衛星搭載の極端紫外光撮像器（EUV）はヘ
リウムイオンからの共鳴散乱光を連続的に撮像し, プラズマ圏の構造や運動を明らかにした. 私は 2000－2001 年に EUV 撮像器が捉えた画像を用いて太陽風の変動に対するプラズマ圏の応答を統計的に調べた. その結果, 太陽風の変動に応じてプラズマ圏が運動するまでに 10－30 分, 平均 18 分の遅れがあることを明らかにした. この結果は太陽風電場が磁気圏前面から電離圏を経由して夜側内部磁気圏へと伝播することを示唆している.
また 2007 年に打ち上げられた月周回衛星 SELENE（かぐや）では極端紫外光望遠鏡（UPI-TEX）によるヘリウムイオン（30.4 nm）および酸素イオン（83.4 nm）の撮像が計画されている. この観測では月周回軌道からの撮像であるため常にプラズマ圏全体を視野内に収めることができデータ数の向上が見込まれる. さらに人工衛星による酸素イオンの撮像は世界初の試みであり, 極域から流出す
るプラズマの経路の同定に大きな貢献が期待される. 私は UPI-TEX 撮像器の地
上較正試験を行い, 波長 30.4 nm において 0.014 cps/Rayleigh/bin, 波長 83.4 nm において 0.0036 cps/Rayleigh/bin の感度をもつことを確認した. UPI-TEX 撮像器は時間分解能 2 分および空間分解能 0.09 RE でプラズマ圏のヘリウムイオンを, また時間分解能 30 分および空間分解能 0.5 RE で極域から流出する酸素イオンをそれぞれ撮像でき, 科学目標を十分達成できる.

太陽風の早い変動とプラズマ圏の運動をより高精度で比較するには, プラズ
マポーズのわずかな変動を捉えられる高い位置分解能をもつ検出器が必要となる. そこで私は将来のプラズマ圏撮像計画に向け, 2 段と 3 段を組み合わせたマ
イクロチャンネルプレート（MCP）とレジスティブアノードを用いた検出器を開発した. 私はシミュレーションおよび実験を行い検出器の性能を評価した.
その結果, 2 段と 3 段の MCP の間に逆向きに電圧を印加することで波高分布の幅を狭め, 位置分解能を UPI-TEX 撮像器に比べ約 3 倍に向上できることを確認した.

Contents
1

2

General Introduction...........................................................................................1
1.1

The Earth’s Plasmasphere ..............................................................................1

1.2

Imaging the Plasmasphere .............................................................................6

1.3

Construction of this Thesis .......................................................................... 10

The Plasmapause Response to the Southward Turning of the IMF Derived

from Sequential EUV Images ................................................................................... 113

2.1

Introduction................................................................................................. 11

2.2

Data Selection ............................................................................................. 13

2.3

Analysis ...................................................................................................... 16

2.4

Discussion ................................................................................................... 23

2.5

Summary ..................................................................................................... 25

Calibration of Telescope of Extreme Ultraviolet (TEX) onboard SELENE.... 26
3.1

Introduction................................................................................................. 26

3.2

Science Target ............................................................................................. 28

3.2.1 Global Image of the Plasmasphere ........................................................... 28
3.2.2 Plasmasphere Response to the Interplanetary Magnetic Field (IMF) ........ 29
3.2.3 Plasmasphere Refilling ............................................................................ 29
3.2.4 Continuous Leakage of Plasmaspheric Materials...................................... 30
3.2.5 Inconsistency of Old and New Topics on the Plasmasphere...................... 31
3.2.6 Net Escape Rate of Earth’s Oxygen to Space ........................................... 31
3.2.7 Short- and Long-terms Variation of O+ Outflow ....................................... 32
3.3

Instruments.................................................................................................. 32

3.3.1 Mo/Si Multilayer Coated Mirror .............................................................. 35
3.3.2 Metallic Thin Filter.................................................................................. 39
3.3.3 Microchannel Plates (MCPs) and Resistive Anode Encoder (RAE).......... 41
3.3.4 Electronics and High-Voltage Power Supply ............................................ 44
3.3.5 Total Sensitivity....................................................................................... 45
3.4
4

Summary ..................................................................................................... 49

Development of High-resolution Imaging Detector using Five Microchannel

Plates and a Resistive Anode Encoder...................................................................... 50
4.1

Introduction................................................................................................. 50

4.2

Calculation .................................................................................................. 52

5

4.3

Measurement ............................................................................................... 55

4.4

Results and Discussion ................................................................................ 59

4.5

Summary ..................................................................................................... 62

Concluding Remarks......................................................................................... 64

Acknowledgements ................................................................................................... 67
References ................................................................................................................. 68

1 General Introduction
1.1 The Earth’s Plasmasphere
The magnetosphere is a region around the Earth in which the Earth’s magnetic
field is the dominant influence and controls the motions of charged particles. Figure 1.1
shows the results of a synthesis of in-situ and ground-based observations over 40 years
and represents today’s generally accepted global view of the Earth’s magnetosphere. In
the inner magnetosphere, the magnetic field configuration is dipolar like with
perturbation at its outer edge caused by the external pressure change of the solar wind.
At ionospheric altitude below 1000km in the mid and high latitude, the dipole field isnearly vertical. The charged particles in the ionosphere are free to move upward along
magnetic field lines and populate the higher-altitude regions of the Earth’s
magnetosphere. The charged particles, which are predominantly electrons and protons
with energies not exceeding several electron volts [eV], form a region of cold dense
plasma (shown in Figure 1.1) and is called the plasmasphere.
The existence of a sharp outer boundary of the plasmasphere, corresponding to
the plasmapause, was first discovered in the early 1960s [Carpenter, 1963]. Analysis ofwhistler waves recorded at ground stations established the plasmapause as a permanentfeature within the magnetosphere. After the establishment of the definition of theplasmasphere, satellite in-situ measurements were used to investigate the innermagnetosphere, particularly the plasmasphere in the 1970s. Figure 1.2 shows thedominance of H+, He+, and O+ measured by a mass spectrometer onboard the OGO 5
satellite, showing density profiles plotted versus the L-shell parameter (the distance ofthe equatorial crossing point of the field line from the Earth’s center, measured in
Earth’s radii). The density profile of He+ is very similar in shape to that of H+. The
density distribution of those ions sharply decreases by a factor of 100 or more. In Figure
1.2, the position of the outer boundary, the plasmapause, at L=4.9 is evident in both H+
and He+. The outer plasmasphere consists of about 90% H+ and about 10% He+ with
only a trace of O+.

1

Figure 1.1: A schematic diagram of a generally accepted global representation of
the Earth’s magnetosphere. This picture has been synthesized from a vast numberof localized observations made over the past four decades throughout many of theregions [Williams et al., 1992].

2Figure 1.2: A typical inbound plasmapause crossing observed by OGO 5 showing
the H+, He+, and O+ ion concentrations as a function of L and local time [Chappell
et al., 1970].
On the other hand, there have been theoretical approaches for thecomprehension of the plasmasphere [Nishida, 1966]. Nishida [1966] represented a clear
explanation of the formation of the plasmapause. There is a large-scale electric field
directed from dawn to dusk. This electric field originated from the solar wind magnetic
field and penetrates into the magnetosphere. The plasma in the magnetosphere drifts toE B direction due to this electric field (E is the convection electric field, and B is the
Earth’s magnetic field). When the corotation with the Earth is superposed on alarge-scale convection, a convection profile shown in Figure 1.3 is obtained. The
motions of cold plasmas, such as those in the plasmasphere, are explained by the E B
drift convection, since their gradient and curvature drifts are very small. In Figure 1.3, it
can be seen that the streamlines of the plasma convections can be categorized into two
distinctly different flow regimes by a separatrix between open and closed drift paths:
those which pass through the magnetosphere, and those which circulate around theEarth. The latter closed region inside the separatrix continues to be filled up to highdensities with ionospheric plasmas. Outside this separatrix, cold plasmas on an open
drift trajectory generally drift sunward from the tail, and flux tubes are filled up to3

relatively lower densities of cold plasmas due to their shorter lifetimes in themagnetosphere. As a result, the separatrix can be recognized as the outer boundary ofdense plasma, i.e., the plasmapause.

Figure 1.3: Streamlines of the magnetospheric plasma motion obtained bysuperposing the motion due to the Earth’s rotation on the flow of thesolar-wind-induced convection. The boundary between two groups of streamlines,
i.e., those that pass the tail and those that do not, is suggested to be theplasmapause (dotted curve). The region of turbulence is also indicated [Nishida,1966].
Intensive observation using spacecraft followed and revealed the dynamical
aspect of the plasmasphere changing depending on geomagnetic conditions. It became
clear that the plasmapause position is closely correlated with average magnetic activity
over a 2 to 6-hour period, e.g., the Kp index [Chappell, 1970]. The plasmasphere shrinks
with increasing the magnetic activity in a very well-behaved, predictable manner. Thisvariation in the radius of the plasmasphere depending on the magnetic activity is shownin Figure 1.4. This figure represents the superposed results over more than 40 passes ofthe satellite. The decrease in the radius of the plasmapause is accompanied by an4increased steepness in the density gradient at the plasmapause, while the total
concentration levels inside and outside the plasmapause remain approximately ~103
ions/cm3 and ~1 ions/cm3, respectively.

Figure 1.4: A composite graph of H+ concentration measured by OGO 5 showing
the reaction of the plasmasphere in the nightside region to changes in the level ofmagnetic activity. The plasmapause is found to be steep and move to lower L
values with increasing activity [Chappell et al., 1970].
When the southward IMF encounters the Earth’s northward magnetic field at
the dayside magnetopause, it leads to the process of magnetic field line merging or
reconnection, driving sunward convection in the inner magnetosphere [Dungey, 1961].
Figure 1.5 shows a schematic diagram of the plasma motion in the enhanced convection
electric field. The enhanced convection electric field triggered by a southward turningof the IMF forms new E B drift trajectories, and particles that were previously on theoutermost closed trajectories suddenly find themselves on open trajectories and drift
toward the dayside magnetopause. Thus, the outer layers of the plasmasphere are
stripped away and the plasmasphere shrinks.
5

Figure 1.5: A schematic diagram of the plasma motion in the convection electricfield enhanced by the southward turning of the IMF. The dashed line indicates thepreviously formed plasmapause. After the enhancement, the nightside plasmapause
on the previous plasmapause move inward as black arrows along the newly formed
E B drift trajectories. The dayside plasma moves sunward and the structure called
as the dayside bulge or plume forms.

1.2 Imaging the Plasmasphere
In-situ measurements along the spacecraft’s path were intensively performed in
the early 1970s, and revealed dynamic aspects of the plasmasphere. However, thein-situ measurements are local observations with which it is difficult to separate thetemporal development and the spatial structure of the plasmasphere. To overcome thisdifficulty, various efforts have sought to understand the global distribution of theplasmasphere by visualization.
Helium and oxygen ions are the second and third major components in theplasmasphere, as shown in Figure 1.2, and resonantly scatter the solar extreme
ultraviolet (EUV) emission at 30.4 nm (He II) and 83.4 nm (O II), respectively. For thisreason, detecting these emissions leads us to the global imaging of the plasmasphere.
The fundamental technology to detect the EUV emission began with the He II emission6

(30.4 nm), which was first detected on a sounding rocket experiment [Johnson et al.,1971]. After this discovery, proof-of-concept observations to give inside-out views of
the plasmasphere have been performed by both sounding rockets (e.g. Ogawa andTohmatsu, 1971; Paresce et al., 1974; Yoshikawa et al., 1997) and the spacecraft (e.g.
Weller and Meier, 1974: Chakrabarti et al., 1982; Meier et al., 1998). However, it isdifficult to extract the geophysical information on the outer plasmasphere from its
inside.
The two-dimensional He II imaging of the Earth’s plasmasphere from its
outside was first done by the Planet-B (Nozomi) spacecraft [Nakamura et al., 2000].
Yoshikawa et al. [2000a, 200b, 2001] first presented static EUV images of theplasmasphere and the inner magnetosphere. Figure 1.6 represents the 2-D image
observed by the XUV scanner onboard the Planet-B spacecraft. The image shows theplasmasphere shape averaged over the noon-duskside sector. The dipole magnetic fieldlines of L = 4 and 5 are also shown. While the inner magnetic field line corresponds tothe plasmapause, the outer magnetic field line corresponds to the outer boundary of the
dusk bulge [Yoshikawa et al., 2000a].

Figure 1.6: Two-dimensional imagery assuming the symmetrical shape of theplasmasphere observed by NOZOMI. The dipole magnetic field lines of L = 4 and5 are also shown. While the inner magnetic field line corresponds to theplasmapause, the outer magnetic field line corresponds to the outer boundary of the
dusk bulge [Yoshikawa et al., 2000a].

7

Recent advances in satellite-based imaging techniques have made it possible toroutinely obtain full global images of the plasmasphere. The Extreme Ultraviolet
Imager (EUV) on the Imager for Magnetopause-to-Aurora Global Exploration(IMAGE) satellite gave us complete sequential pictures [Burch et al., 2001a, 2001b;
Sandel et al., 2000, 2001]. The EUV instrument detects the distribution of He+ in theplasmasphere with a time resolution of 10 min and a spatial resolution of 0.1 RE at theapogee (~8 RE). The EUV instrument consists of three sensor heads and their fields ofview are joined to create a single image.
The EUV images of the plasmasphere reveal a host of distinct features andbehaviors [Sandel et al., 2001, 2003]. Figure 1.7 shows an example of the EUV images.
The sun is toward the lower right, and the Earth’s shadow extends through theplasmasphere in the opposite direction, toward the upper left. The sharp drop of
intensity corresponds to the plasmapause. This image also illustrates some aspects of theplasmasphere that have proven especially accessible by EUV imaging techniques, ashoulder (lower center), a drainage plume (right), and a pre-midnight channel (top).

Figure 1.7: Two-dimensional imagery observed by Extreme Ultraviolet Imager
onboard the IMAGE satellite on 24 May 2000, during a magnetic storm. A
drainage plume extends sunward past the right edge of the image. The Earth’s
shadow is visible in the direction away from the Sun, because only He+ illuminatedby the Sun scatters the target wavelength. The white circle shows the approximate
size and position of the Earth [Sandel et al., 2003].

8

The recent imaging instruments in the EUV region mentioned above have
similar types of optics. Figure 1.8 shows schematic cross-sectional drawings of (a) the
XUV scanner onboard the Planet-B spacecraft and (b) one of the EUV sensor heads
onboard the IMAGE satellite. Each instrument is both designed as a type of
normal-incidence telescope. Their optics consists of a multilayer coated mirror to focus
EUV photons, a metal thin filter to eliminate unwanted longer wavelength emissions,and microchannel plates to pick up photon events.

Figure 1.8: (a) A schematic cross-sectional drawing of the XUV scanner onboard
the Planet-B (Nozomi) spacecraft [Nakamura et al., 2000]. The optics consists of amultilayer-coated mirror, a thin metal filter, and microchannel plates. (b) A
cutaway drawing of the EUV sensor onboard the IMAGE satellite [Sandel et al.,2000].
In the near future, an instrument which has a similar type of optics to thoseonboard Planet-B and IMAGE will take global images of the plasmasphere from thelunar orbit. The Selenological and Engineering Explorer (SELENE) was launched by
the H-IIA rocket in September of 2007 and then put into the orbit around the Moon. The
Telescope of Extreme Ultraviolet (TEX) mounted in the Upper atmosphere and Plasma
Imager (UPI) on SELENE detects the resonance scattering emissions of oxygen ions (O9II: 83.4 nm) and helium ions (He II: 30.4 nm) to take images of near-Earth plasmas. The
UPI-TEX imager has a time resolution of 1 minutes and a spatial resolution of 0.1 RE.
The observation of the plasmasphere by UPI-TEX from the lunar orbiter will give us
not only much more sequential images with higher time resolution but also a deeper
understanding of the plasmasphere dynamics from a different perspective.

1.3 Construction of this Thesis
The main purpose of this thesis is to examine the plasmasphere dynamics bymeans of a global imaging technique in the extreme ultraviolet (EUV) spectral range.
This thesis consists of 5 chapters. Chapter 2 presents the statistical research on theplasmapause response to the southward turning of the IMF observed by the Extreme
Ultraviolet Imager (EUV) onboard the IMAGE satellite. In Chapter 3, the detailed
description of the Telescope of Extreme Ultraviolet in the Upper atmosphere andPlasma Imager (UPI-TEX) onboard the SELENE satellite is presented. Chapter 4
includes the characteristics of the newly developed high-resolution imaging detector
using five microchannel plates (MCPs) and a resistive anode encoder (RAE) for the
future missions. Concluding remarks are given in Chapter 5.

10

2 The Plasmapause Response to the Southward Turning of
the IMF Derived from Sequential EUV Images
2.1 Introduction
The plasmapause is the outer boundary of the Earth’s plasmasphere which is a
torus-shaped region filled with cold plasmas in the inner magnetosphere. Decades ofobservation show that the radial location of the plasmapause generally moves inward
during geomagnetic disturbance periods. The inward motion is explained by alongstanding hypothesis that the dynamics of the plasmasphere is controlled by thesuperposition of the corotation and convection electric field [Nishida, 1966]. The
enhanced convection electric field triggered by the southward turning of the IMF forms
new E B drift trajectories, and plasmas on the plasmapause move inward along them
[Chappell et al., 1970].
Some other mechanisms explaining the inward motion of the plasmapause
during geomagnetic disturbance periods have been proposed. Lemaire [1974, 1985]
proposed that the centrifugal force drives the plasma upwards and then produces a sharp
density gradient along the magnetic field lines tangent to the zero parallel force (ZPF)
surface. According to this mechanism, after an increase in the level of magnetic activity
the ZPF surface shifts inward and the plasmasphere is peeled off in the post midnight
sector where the convection velocity is maximum [Lemaire and Gringauz, 1998].
Carpenter and Lemaire [1997] discussed the loss of plasmas through the flow from theplasmasphere into the underlying ionosphere during periods of the enhanced convection.
They indicated the possibility of the mechanism involving dumping of plasmas into theionosphere based on the whistler data.
Although it is not completely understood how the convection electric field
governing the dynamics mentioned above propagates to the inner magnetosphere, anumber of observational studies have examined the ionospheric response to the IMF.
For example, the ionospheric convection responds to the IMF quickly (~10 min) [Ridley
et al., 1998] and almost simultaneously at the whole ionosphere [Kikuchi et al., 1996;
Ruohoniemi and Greenwald, 1998; Ridley et al., 1998; Lu et al., 2002]. Furthermore,using the ground magnetometers, Hashimoto et al. [2002] examined the response timeof the convection in the inner magnetosphere to the enhancement of the ionospheric
convection. Their study suggests that the convection electric field propagates from the11

dayside magnetopause to the nightside inner magnetosphere through the ionosphere.
With regard to experimental work on the plasmasphere, the conventional
extreme ultraviolet (EUV) photometric experiments from an inside-out view, which had
been done intensively in the 1970s, suggested that the Earth’s plasmasphere would be
imaged by He II (30.4 nm) emission [e.g., Johnson et al., 1971]. In the 1990s, theremote sensing method using the EUV emission became a powerful tool to provide
global perspectives on the plasmasphere dynamics [Williams et al., 1992]. The
fundamental technology to detect the EUV emission began with the He II emission
through a rocket experiment [Yoshikawa et al., 1997]. The two-dimensional He II
imaging of the terrestrial plasmasphere from its outside was first done by the Planet-B(Nozomi) spacecraft [Nakamura et al., 2000]. Yoshikawa et al. [2000a, 2000b, 2001]
first presented static EUV images of the plasmasphere and the inner magnetosphere.
Recent advances in satellite-based imaging techniques have made it possible toroutinely obtain full global images of the plasmasphere. The EUV instrument on the
Imager for Magnetopause-to-Aurora Global Exploration (IMAGE) satellite gave us
complete sequential pictures (one image/10 min) [Burch et al., 2001a, 2001b; Sandel etal., 2000, 2001]. Using these sequential images, Goldstein et al. [2003b, 2003c, 2004a,2005] and Spasojević et al. [2003] identified the correlation between the inward motionof the plasmapause and the southward turning of the IMF for specific events. They
found that the timescale of the plasmapause response to the IMF is about 30 min in theircase studies. Furthermore, Pierrard and Cabrera [2005, 2006] presented thesimulations of the plasmapause formation and compared the predicted plasmapause
positions with the EUV observations. They remarked the correlation between the
inward motion of the plasmapause and the southward turning of the IMF in numericalsimulations. However, they did not discuss the timescale of the plasmapause response to
the IMF in their studies.
In this study, I statistically examine the plasmapause response to the southwardturning of the IMF in order to understand the propagation mechanism of the convection
electric field in the inner magnetosphere. I survey the EUV data set in the period of May
2000–December 2001 obtained by the IMAGE satellite (http://euv.lpl.arizona.edu/euv/),
and find the events showing the clear inward motion of the nightside plasmapause.
Finally, the timescale of the plasmapause response to the southward turning of the IMF
is deduced.

12

2.2 Data Selection
The EUV instrument on the IMAGE satellite imaged the He+ distribution in theplasmasphere by detecting resonantly scattered solar 30.4-nm radiation. The measured
brightness is directly proportional to the He+ column density along the line of sightbecause the 30.4-nm emission from the plasmasphere is optically thin [Sandel et al.,2001]. The EUV imager consisted of three cameras, and their fields of view wereoverlapped to create a single image. It produced sequential images of the plasmasphere
every 10 min for generally 7–9 hours out of each 14-hour orbit, with spatial resolutionof ~0.1 RE as seen from the apogee (~8 RE).
The EUV instrument sometimes observed the inward motion of the nightsideplasmapause driven by the southward IMF, called as erosion, in a time series of global
snapshots [Goldstein et al., 2003b, 2003c, 2004a, 2005; Spasojević et al., 2003]. Figure
2.1 shows sequential images of the plasmasphere (one image/10 min) taken by the EUV
instrument on 22 April 2001. In each image the direction of the Sun is to the upper left
and the dotted circle represents the Earth. The left side data in each image are notavailable because throughout the event the camera whose field of view covered theduskside was usually turned off to prevent sunlight from entering the camera [Goldstein
et al., 2005]. In this period it is clear that the nightside plasmapause, especially at post
midnight (the lower right), become sharper and move inward by ~1.5 RE.
To investigate the inward motion of the plasmapause driven by the southwardturning of the IMF, I searched the EUV data set in the period of May 2000-December
2001, and identified 16 events according to the following criteria:–The northward IMF continued more than 10 min before the southward turning.
–The IMF turned southward more rapidly than 0.25 nT/min.
–The southward IMF continued more than 10 min after the turning.
–The EUV instrument clearly observed the sharp plasmapause on the nightside andthe whole inward motion of the nightside plasmapause in the field of view.
Figure 2.2 shows an example of the variation of the IMF BZ component(BIMF_Z) satisfying the above criteria, measured by the MAG instrument onboard the
ACE satellite [Smith et al., 1998] on 22 April 2001. The IMF turned southward at 01:27
and 02:15 UT distinctly. I used 1-min interval data set of the IMF. The arrival time ofevery IMF at the average magnetopause (10 RE) from the ACE position was estimated13

as follows.
(1) It was assumed that the every IMF measured by the MAG instrument continued topropagate and arrived at the magnetopause and that the magnitude and polarity of the
IMF did not change.
(2) I used the solar wind velocity data simultaneously measured by the SWEPAM
instrument on the ACE satellite [McComas et al., 1998] and assumed the velocity did
not change. The arrival time of every IMF at the magnetopause was calculated based on
the position of the ACE satellite.
(3) Every IMF was labeled the arrival time at the magnetopause. This method
provides more correct information than by simply assuming a constant velocity of thesolar wind.

14

Figure 2.1: Sequential images of the plasmasphere taken by the EUV instrument
onboard the IMAGE satellite on 22 April 2001. In each image the direction of the
Sun is to the upper left and the dotted circle represents the Earth. Each image is 10
min apart. The left side data in each image are not available throughout the event.
In this period the nightside plasmapause is found to be sharp and move inward by
about 1.5 RE.

15

Figure 2.2: The 1-min interval data set of the IMF BZ component (in GSM
coordinates) on 22 April 2001 measured by the MAG instrument onboard the ACE
satellite. The arrival time of every IMF has been compensated at the averagemagnetopause (10 RE) from the ACE position (~213 RE), taking into account thevariation of the solar wind velocity.

2.3 Analysis
Figure 2.3a shows two snapshots of the He+ plasmasphere taken at 02:29 and
04:32 UT on 22 April 2001 by the EUV imager. They are also included in Figure 2.1.
The black circle indicates the Earth and the white arrow points sunward in each image.
The shadow of the Earth is seen in the antisunward direction, and the bright arc at the
Earth’s dayside limb is airglow from He and O+ [Burch et al., 2001a, 2001b]. The left
side data are not available as explained in the previous section. The white squares
indicate the manually identified plasmapause along the sharp brightness gradient. In
Figure 2.3a, the sharp He+ edge is identified in the EUV image, and it corresponds tothe plasmapause [Roelof and Skinner, 2000; Goldstein et al., 2003a].

16

Figure 2.3: (a) Images of the He+ plasmasphere taken by the EUV instrument on
22 April 2001. The black circle indicates the Earth and white squares indicate themanually identified plasmapause in the EUV images. The left side data are notavailable. (b) The plasmapause of each image mapped to the magnetic equatorial
plane. The black squares indicate the manually identified plasmapause in themagnetic equatorial plane, and the dashed circles are drawn at L = 2, 4, 6. The
nightside plasmapause at 04:32 UT is about 1.5RE closer to the Earth than that at
02:29 UT.
On the assumption that the plasmapause is field aligned in the dipole magnetic
field, the field line with the minimum L (radial distance in the equatorial plane) along
the line of sight to each identified point is found, as shown in Figure 2.4. Then theplasmapause can be mapped to the magnetic equatorial plane along the field line, like17Goldstein et al. [2003a]. There is some subjectivity involved in manually extracting theplasmapause. This manual extraction of the plasmapause was repeated and thedispersion of the identified plasmapause locations was calculated. As a result, theuncertainty due to the subjectivity was estimated to be about 0.2 RE, which is consistent
with that reported in Goldstein et al. [2003a, 2005]. Figure 2.3b shows the plasmapause
locations at 02:29 and 04:32 UT on 22 April 2001 mapped to the magnetic equatorial
plane using this method. It is clear that the nightside plasmapause at 04:32 UT is about
1.5 RE closer to the Earth than that at 02:29 UT in Figure 2.3b.
The plasmapause radial velocity (Vpp) was calculated after mapping theplasmapause in the EUV images to the equatorial plane. Using centered timedifferencing of the sequential plasmapause positions, Vpp was calculated at the particular
MLT where the plasmapause was clearly identified and the inward motion was mostpronounced. For example, Figure 2.5 shows the variation of Vpp at 2.1 MLT on 22 April
2001. Vpp has the subjective error due to the subjectivity in manually extraction of theplasmapause (0.2 RE) as mentioned in the previous section. Using this value, theaverage uncertainty in Vpp was estimated to be about ±0.4 RE/hour, which is consistent
with that discussed by Goldstein et al. [2003a, 2003b]. As seen in Figure 2.2 and Figure
2.5, BIMF_Z and Vpp have very similar variations with a time lag.

Figure 2.4: A schematic diagram of the plasmasphere and the line of sight of EUV
in the meridional plane. The plasmapause in the EUV images was mapped to thegeomagnetic equatorial plane as follows. The points corresponding to theplasmapause (black dot) are manually selected along the sharp brightness gradienton each image. Next, the minimum L along the line of sight is found for eachselected point, and the field line is traced to the magnetic equator (red dot).

18

Figure 2.5: Radial plasmapause velocity (Vpp) on 22 April 2001 at 2.1 MLT where
the inward motion of the nightside plasmapause was most pronounced. Vpp < 0
indicates that the plasmapause moves inward. The prominent inward motion of theplasmapause occurs in the period between 02:29 and 04:32 UT, and the locationsof the plasmapause at both times are presented in Figure 2.3. The error of Vpp dueto the manually extraction of the plasmapause is about ±0.4 RE/hour.
The solar wind and IMF impose an electric field potential across themagnetosphere, driving the convection. Therefore to examine the correlation between
the plasmapause motion and the IMF turning, it is meaningful to express both in electricfield parameters. On the assumption that the plasmas on the plasmapause move by E B
drift (see Figure 1.5), the plasmapause motion is represent as the electric field parameter(Epp) defined as
E pp

Vpp Bdipole ,(2.1)

where Bdipole is the equatorial dipole geomagnetic field. I calculate Bdipole using a simple
dipole formula, Bdipole = BE/L3 (BE = 3.11 105 T). Epp approximately corresponds to theconvection electric field in the midnight region. The variation of Epp is dominated bythat of Vpp because Bdipole is smooth. Similarly, the variation of the IMF is represented as
the electric field of the solar wind (Esw) defined as
E sw

Vsw

BIMF_Z ,

19

(2.2)

where Vsw is the solar wind velocity. I used Vsw data measured by the SWEPAM
instrument onboard the ACE satellite.
Then I calculated the linear correlation coefficient between Epp and Esw withvarious time delays, to investigate the response time of the plasmapause to the IMF.
Figure 2.6 shows the correlation coefficient on 22 April 2001. The maximum value ofthe correlation is 0.73 when a 16-min delay is assumed. Because of the interval of the
EUV data set (10 min) and the subjective error in Vpp, the response time calculated in
my method has some uncertainty. For reference, therefore, the range of 90% of the
maximum correlation value is displayed in Figure 2.6. Both Epp and 16-min delayed Esw
are shown in Figure 2.7, where Epp multiplied by 5 is displayed to show both quantities
in the same magnitude. The number of 5 suggests that the solar wind electric field
penetrates in the inner magnetosphere by 20%. This is consistent with the result of
10–25% reported in previous studies [Goldstein et al., 2003b, 2003c, 2004b, 2005]. As
seen in Figure 2.7, Epp and Esw are clearly correlated.
For the other events, the correlation coefficients were calculated in the sameway. Table 2.1 and Figure 2.8 show the results. The best-correlated time delay between
Epp and Esw is regarded as the timescale of the plasmapause response to the IMF
variation. There are four events in which the maximum values of the correlation
coefficient are below 0.5 in Table 2.1. Among four events, there is common aspect that
the IMF has a shortly periodical structure and therefore there are several local maximum
values in the correlation coefficient. That decreases the overall maximum value in thecorrelation coefficient. Nevertheless, these events also imply that the plasmapause
responded to the southward turning of the IMF. These four events are included in myresults. Finally, my results indicate that the plasmapause response to the southwardturning of the IMF takes 10–30 min, and average 18 min.

20

Figure 2.6: The correlation coefficient between the electric field parameters at theplasmapause (Epp) and in the solar wind (Esw) on 22 April 2001. The maximum
value of the correlation is 0.73 at a 16-minute delay. The range of 90% of the
maximum correlation value is also displayed.

Figure 2.7: Electric field parameters Epp (plasmapause) and Esw (solar wind). The
variations of Epp and Esw are qualitatively the same as Vpp and BIMF_Z, respectively.
Esw is displayed with 16-minute delay, which is the best correlated time delay
between Epp and Esw. Epp has been multiplied by 5 in order to show both quantities
in the same magnitude.

21

Table 2.1: The event list showing the date, MLT where the radial plasmapause
velocity was calculated, the maximum value of the correlation coefficient between
Epp and Esw, the best correlated time delay, time delay variation among 90% of
the maximum correlation value, the solar wind (SW) propagation time from the
ACE satellite to the magnetopause in this study, and previously published studies
This study
DateaPrevious case study

MLT Correlation Delay 90% range
(min)
(min)

SW propagation(min)

29 May 2000
2 Jun 2000

4.0
0.0

0.79
0.67

21
23

16 - 27
18 - 2861
54

3 Jun 2000
8 Jun 2000

0.2
3.0

0.46
0.47

0
22

0 - 2017 - 2757
44

10 Jul 2000
28 Jul 2000

2.0
0.4

0.68
0.51

28
23

13 - 41
19 -2860
56

5 Aug 2000
17 Aug 2000

0.0
2.0

0.49
0.54

19
19

13 - 26
7 - 3351
62

28 Aug 2000
6 Feb 2001

1.5
1.6

0.58
0.66

13
10

7 - 213 - 344370

22 Apr 2001
28 May 2001

2.1
2.8

0.73
0.51

16
18

9 - 280 - 3860
46

1 Jun 2001
2 Jun 2001

3.0
1.0

0.59
0.59

0
24

0 - 2618 - 32

21 Jun 2001
26 Jun 2001

3.0
3.0

0.58
0.45

15
29

6 - 2724 - 50

Delay SW propagation
(min)
(min)

30a

63b

27c

63

71
58

32d

55

43
52

32e

54.8

: [Goldstein et al., 2003b]b: Goldestein et al. [2003b] used the IMF data set measured by the Geotail satellite. They assumed theconstant solar wind velocity of ~400 km/s and calculated the solar wind propagation time by 3.7 minbased on the Geotail position (~24 RE). I extrapolated the propagation time from the ACE position (~238RE).
c: [Goldstein et al., 2005]de: [Goldstein et al., 2003c, 2004a]: [Spasojević et al., 2003]

22

Figure 2.8: The histogram of the time delay between Epp and Esw listed in Table
2.1. It corresponds to the response time of the plasmapause to the IMF variation.

2.4 Discussion
The timescale of the plasmapause response to the southward turning of the IMF
derived from the sequential EUV images was studied for specific events. Table 2.1
shows the results of the previously published case studies. For example, Goldstein et al.
[2003b] found that the response time on 10 July 2000 is 30 min at 2.4 MLT, while it is
28 min at 2.1 MLT in this study. Also, Spasojević et al. [2003] reported that on 26 June
2001, the response time is 32 min, and by my method it is 29 min. The results of myresearch for these two events are consistent with those of their case studies. However, as
shown in Table 2.1, the response times on 2 June 2001 and 22 April 2001 in my results
are different from those reported by Goldstein et al. [2003c, 2004a] and Goldstein et al.
[2005] by ~10 min. In spite of using almost the same techniques to determine theresponse time of the plasmasphere to the IMF as theirs, my estimates show shorter
response times than those in their studies. These discrepancies would mainly result fromthe difference in the method to compensate the propagation time of the IMF from thesatellite to the magnetopause. Goldstein et al. [2003c, 2004a, 2005] assumed theconstant solar wind velocity. I also calculated the response time using their method andfound the same result. But in order to calculate the correlation more precisely, I took
into account the variation of the velocity.
The plasmapause response to the southward turning of the IMF provides

23

important information about the complex coupling of the inner/outer magnetosphere,
ionosphere, and solar wind. In this chapter, I have shown that the convection electricfield is rapidly distributed in the inner magnetosphere within 10-30 minutes, based on
the sequential EUV images of the plasmasphere. The rapid response of the convection
electric field to a sudden change in the IMF is discussed.
In recent years, many authors have argued the quick response of the plasma
convection in the ionosphere [e.g., Hairston and Heelis, 1995; Ruohoniemi andGreenwald, 1998; Ridley et al., 1998; Lu et al., 2002]. Ridley et al. [1998] investigatedthe electric potential in the ionosphere using the ground magnetometers, and found thatthe ionospheric convection starts to change with a time lag of 8.4–13.6 min (Δt1) fromthe arrival of the IMF at the magnetopause. Furthermore, the progress of global
observations by magnetometers and radars with high temporal resolution enable us toinvestigate the response of the ionospheric convection over a wide range of local times
and latitudes simultaneously. In the recent studies of the ionospheric convection, it has
been shown that the convection starts almost simultaneously (Δt2 < 2 min) at all
magnetic local time [e.g., Ruohoniemi and Greenwald, 1998; Lu et al., 2002] and atwhole latitudes [Kikuchi et al., 1996].
Following to these results, Hashimoto et al. [2002] studied a quick response of
the magnetospheric convection to sudden IMF changes and the development of partial
ring current by using the ground magnetometer data obtained from the polar cap to
low-middle latitude stations. They demonstrated that the response time of theconvection in the inner magnetosphere to the enhancement of the ionospheric
convection takes 5–11 min (Δt3). Based on this result, they suggested that theconvection electric field propagates from the dayside magnetopause to the nightside
inner magnetosphere through the ionosphere.
The timescale (Δttotal) that the convection electric field takes to propagate to theinner magnetosphere through the mechanism mentioned above can be estimated to be
13-27 min, based on the past observational studies as follows: Δttotal = Δt1 +Δt2 +Δt3.
This timescale is consistent with the time lag of 10-30 min between the southwardturning of the IMF and the plasmapause motion. Therefore, this study suggests that theconvection electric field penetrates from the magnetopause to the inner magnetospherethrough the ionosphere. But this is not an exclusive mechanism. Further study isnecessary.

24

2.5 Summary
Using the sequential images of the plasmasphere observed by the EUV instrumenton the IMAGE satellite, I examined the plasmapause dynamics. I have shown that theplasmapause response to the southward turning of the IMF takes 10–30 min, andaverage 18 min. It is consistent with the timescale derived from the ionosphericobservations. Therefore it indicates that the electric field penetrates from themagnetopause to the inner magnetosphere through the ionosphere. However, due to thelimited EUV data, only 16 events were found in my search. Future observation of theplasmasphere by the SELENE satellite from the lunar orbit will give us not only much
more sequential images with higher time resolution but also a deeper understanding of
the plasmasphere dynamics from a different perspective.

25

3 Calibration of Telescope of Extreme Ultraviolet (TEX)
onboard SELENE
3.1 Introduction
Our knowledge of the Earth’s plasmasphere has increased steadily owing to anumber of in-situ plasma measurements during the past several decades. The
plasmasphere is filled with cold ions such as H+, He+, and a small amount of O+ and isfilled up to high density (about 100 ions/cm3). The outer boundary, the plasmapause, isinterpreted as a separatrix between closed and open convection trajectories [Nishida,1966]. Under a steady convection electric field condition, the plasmapause is considered
to form a “teardrop” shape on the equatorial plane. The traditional models indicate thatthe plasmapause moves closer to the Earth when the large-scale convection electric field
is enhanced. The plasma which was previously on a closed drift path then changes to an
open drift path. As a result, the open drift path carries the plasma sunward, creating aplasmaspheric tail [Chen and Wolf, 1972] or a detached plasma region [Chappell, 1974].
One question concerning the existence or nonexistence of such a structure has remainedunresolved because the conventional in-situ observations for cold ions have an inherent
shortcoming. The in-situ observations provide only local information, which does notinclude the necessary global aspect of the plasmasphere. It is therefore difficult to draw
a definite conclusion on the above question. Global-scale observation is, therefore, anatural extension to many years of single-point in-situ observations. The primary
motivation of this field of science is to visualize the global shape of the plasmasphere.
Standard theories on the polar wind indicate that only light ions such as H+ andHe+ can overcome the terrestrial gravitational potential to escape from the polar
ionosphere due to a pressure gradient along open field line and that the O+ outflow islimited due to both its large mass and its loss by charge exchange with neutral H in theatmosphere [e.g., Axford, 1968; Banks and Holzer, 1968]. Observations by polar
orbiting satellites such as Dynamic Explorer 1 (DE-1) and AKEBONO (EXOS-D),
however, found that the upward flux of O+ is comparable to that of H+, especially underan active geomagnetic condition and during a period of a high solar activity [e.g.,
Chandler et al., 1991; Abe et al., 1993, 1996]. Furthermore, GEOTAIL observed thecold dense O+ flow of the ionospheric origin in the distant tail lobe [Seki et al., 1996,
1999, and references therein]. The difference between the rate of O+ outflow from the26polar ionosphere and that of O+ escape from the magnetosphere to interplanetary spaceindicates that there should be as-yet unknown transport processes for O+ in themagnetosphere [Seki et al., 2001]. The two-dimensional (2-D) O II imagery, which
provides a global O+ distribution, is expected to identify the transport routes andmechanisms of the cold O+ ions and to reveal a quantitative balance between supply andloss, including thermal plasma.
The remote-sensing methods using the extreme ultraviolet (EUV) emissions of
He+ (He II 30.4 nm) and O+ (O II 83.4 nm) have been expected to be a powerful tool toprovide global perspectives on the escaping processes [Chiu et al., 1986; Meier, 1991;
Williams et al., 1992] because conventional in-situ plasma particle measurement
methods cannot detect plasma particles far from the satellite or thermal plasmas. The
fundamental technology used to detect He II emission began with the rocket
experiments in the 1970s and 1980s [Ogawa and Tohmatsu, 1971; Meier and Weller,
1972; Paresce et al., 1974; Chakrabarti et al., 1982]. The measurements obtained using
this technology clarified that the optical thickness of He+ ions is very thin and that theintensity is very low (10 Rayleigh at most). They also clearly showed the feasibility ofimaging the distribution of plasma around the Earth. In the 1990s, technologicaldevelopments in EUV optics were reported in several articles [Cotton et al., 1993;
Chakrabarti, 1993], and a breakthrough in the multilayer coating technique was first
demonstrated by the sounding rocket experiment [Yoshikawa et al., 1997]. The 2-D He
II imaging of the terrestrial plasmasphere from the outside was subsequently
accomplished by both the Planet- B (Nozomi) spacecraft [Nakamura et al., 2000;
Yoshikawa et al., 2000a, 2001] and the Imager for Magnetopause-to-Aurora Global
Exploration (IMAGE) mission [Burch et al., 2001a, b; Sandel et al., 2001]. Cold He+
ions in the near-Earth plasma sheet were also optically observed by Planet-B[Yoshikawa et al., 2000b].
Optical observations on oxygen ions around the Earth have not yet been
performed because of the difficulty in rejecting the bright H Ly-α line from thegeocorona. The primitive equipment, denoted the eXtreme UltraViolet (XUV) sensor,
was developed, that has thick indium film for the band-pass filter. The efficiency ratioof the O II emission to the H Ly-α line of this equipment is high enough to detect O II
emission [Yamazaki et al., 2002]. On December 4, 2000, this equipment flew onboard
the sounding rocket SS-520-2 to investigate the mechanism of ion acceleration and/or27

heating at the cusp/cleft region. The XUV sensor successfully detected O II emission
from the uppermost part of the polar ionosphere. The altitudinal variation from 150
through 1100 km indicated the existence of O+ ions beyond the polar ionosphere andsuggested that the O+ ions energized in the cusp/cleft region may drift to the uppermost
part of the polar ionosphere.
An upgrade version of this instrument onboard the SELenological and
ENgineering Explorer (SELENE) will obtain O II images of the polar wind and He II
images of the plasmasphere from the lunar orbit. The SELENE was launched by the
H-IIA rocket in 2007 to be put into the orbit around the Moon. The aims of the
SELENE project are to carry out scientific observations of the Moon, at the Moon, andfrom the Moon. The Upper atmosphere and Plasma Imager (UPI) on SELENE takes
2-D visible and extreme ultraviolet images of atomic and plasma distribution in theupper atmosphere, ionosphere and around the Earth. The component has two telescopes;
one is Telescope for VISible light (UPI-TVIS), and the other is Telescope for EXtreme
ultraviolet light (UPI-TEX). The UPI-TVIS imager detects four emission lines (427.8,
557.7, 589.3, 630.0 nm) to simultaneously take auroral images around both of the
Earth’s polar regions. The UPI-TEX imager detects the resonance scattering emissionsof oxygen ions (O II: 83.4 nm) and helium ions (He II: 30.4 nm) to take images ofnear-Earth plasmas. In this chapter the design of the UPI-TEX imager and its calibrated
performance are presented.

3.2 Science Target
The UPI-TEX instrument provides global images of plasmas in the vicinity of
the Earth by detecting He II (30.4 nm) emission scattered by He+ ions and O II (83.4nm) emission by O+ ions. The observational targets are the plasmasphere and the polarwind. In the following subsections, the scientific motivations and open questions whichhave not been resolved by past satellite observations are presented.
3.2.1 Global Image of the Plasmasphere
The intensity of He II (30.4 nm) emission from the main body of theplasmasphere is reported as about 5 Rayleigh by Nakamura et al. [2000]. This result is
consistent with early measurements of 1-10 Rayleigh done by sounding rockets fromthe inside-out view [e.g., Ogawa and Tohmatsu, 1971] and with simulated EUV images28

of the plasmasphere [e.g., Roelof et al., 1992]. The intensity from the outer boundary of
the plasmasphere has not been clarified yet by the IMAGE mission, but reported to bein the range of 0.1-5 Rayleigh by the Planet-B satellite [Yoshikawa et al., 2000a]. The
intensity of 0.1-0.5 Rayleigh is considered to be a lower threshold level of the detection
to grasp the global shape of the plasmasphere.
3.2.2 Plasmasphere Response to the Interplanetary Magnetic Field (IMF)
I used the sequential EUV images of the plasmasphere obtained by NASA’s
IMAGE satellite and investigated the plasmaspheric response time to the solar windelectric field that was measured by the solar monitor (the ACE satellite). It has beenshown in Chapter 2 that the plasmapause response to the southward turning of the IMF
takes 10–30 min, with an average of 18 min [also reported in Murakami et al., 2007].
The result is consistent with the time scale derived from the ionospheric observations on
ground stations [Ridley et al., 1998; Ruohoniemi and Greenwald, 1998; Lu et al., 2002;
Kikuchi et al., 1996; Hashimoto et al., 2002]. Based on this accordance of data, I
concluded that the electric field penetrates from the magnetopause to the innermagnetosphere through the ionosphere. However, the result should be re-investigatedusing a more completed data set of EUV images because only 16 events were found due
to limited EUV data. The IMAGE satellite orbits near the Earth. As such, it cannot
always monitor the whole plasmasphere, and there is always an interval of at least
several hours between images, which makes it difficult to estimate the time lag between
the IMF arrival at the magnetosphere and plasmasphere evolution. TEX observation
improves this difficulty because SELENE provides the best platform to command apanoramic view of the whole plasmasphere. As cited above, the intensity of 0.1-0.5
Rayleigh is a lower threshold level to grasp the plasmaspheric response.
3.2.3 Plasmasphere Refilling
The refilling of the plasmasphere has attracted considerably theoretical interestfor several decades [see “Plasmasphere Refilling” special section in Journal ofGeophysical Research, 97(A2), 1047–1179, 1992]. Scientifically important studies onplasmasphere refilling were carried out in the 1960s and early 1970s. These wereobservational [Park, 1970; Chappell et al., 1971] and theoretical and/or modelingstudies [Banks et al., 1971; Schulz and Koons, 1972; Chen and Wolf, 1972; Grebowsky,

29

1971] that addressed various refilling problems. One of the unsolved and highlighted
issues is how plasma distribution approaches a diffusive equilibrium state under thecondition of refilling, where plasma is supplied from the conjugate ionospheres throughfield-aligned supersonic flows. Most models dealing with refilling of the plasmasphererely on initial conditions, such as plasma density, flow velocity, and temperature, in an
empty flux tube. However, there is no theoretical basis to decide on the initial
distributions, nor is there firm guidance from the observations. Because of the scarcity
of information, the initial conditions are arbitrarily chosen in the models. The choice isgenerally only based on the fact that the equatorial plasma density in a depleted flux
tube is as low as 1 cm−3.
We do not know how this low equatorial density connects with the highionospheric density. In this regard, systematic observation along a depleted flux tube
should reveal the subsequent refilling process. It is impossible to carry out suchobservations, however, using a single satellite, because it is impossible to know
latitudinal and longitudinal distributions of plasma simultaneously.
The EUV imaging is a promising methodology to overcome this difficulty.
Based on the theoretical considerations, a time resolution of 30 min is necessary to
examine the refilling issue [Wilson et al., 1992]. The measured He II (30.4 nm) intensityof the refilling trough region is 0.02-0.1 Rayleigh as cited above. TEX should achieve
this capability of time and spatial resolutions.
3.2.4 Continuous Leakage of Plasmaspheric Materials
Recent studies have revealed another feature of the plasmasphere. According toin-situ plasma measurements carried out on multiple geosynchronous satellites, thestandard teardrop model of the plasmasphere, where the bulge size and location aredetermined by the convection electric field strength, is not a completely comprehensive
description even under low geomagnetic activity [Moldwin et al., 1994]. Furthermore,
Matsui et al. [1999] have reported that GEOTAIL on a higher L value path (L~10) often
found plasmaspheric materials independent of Kp. These researchers also concluded
that cold plasma continuously moves toward the magnetopause. These results force us
to verify the classical understanding and morphology of the plasmasphere/plasmapausethat were constructed in the 1960s.
According to Matsui’s statistical analysis, the structure of cold ions along the30

GEOTAIL path (L>10) is quasi-stable, with a density of 2–10 cm−3. If it is assumed that
10% of such a population consists of He+ ions, that they spread along the dusk-to-dawn
direction of 10 RE, and that the g-factor of He II (30.4 nm) emission is 2×10−5 (ion−1 s−1),then the intensity of He II (30.4 nm) from this source is estimated at 0.026-0.13
Rayleigh. Thus, it is necessary to construct optics with a lower threshold level of about
0.02 Rayleigh to identify this unique phenomenon.
3.2.5 Inconsistency of Old and New Topics on the Plasmasphere
The subject of the continuous leakage of cold plasma introduced in Section
3.2.4 may be related to the issue of plasmasphere refilling. On one hand, in recent
studies, such as that of Matsui, it has been possible to determine the rate of ion leakage
from the plasmasphere provided that the amount of plasma supplied directly from theionosphere is not dominant. On the other hand, such as in studies of the plasmasphererefilling, it is difficult to determine the evolution of the content in the same flux tube
from the in-situ observation, because a flux tube drifts regardless of a satellite orbital
motion. The early measurements on the refilling rate of the ionospheric plasma may
include the contribution of plasma leakage from the inner plasmasphere. Consequently,
the refilling rate may be overestimated. The EUV observation provides promising clues
for the verification of the above two processes—i.e., the plasmasphere refilling and acontinuous leakage of plasmaspheric ions.
3.2.6 Net Escape Rate of Earth’s Oxygen to SpaceThe observed polar outflow rate for O+ ions corresponds to the loss of about
18% of the present-day atmospheric oxygen over 3 billion years. However, part of thisapparent loss can actually be returned to the atmosphere. Examining loss rates of four
already identified escape routes with high-altitude spacecraft observations, Seki et al.
[2002] showed that the total oxygen loss rate inferred from current knowledge is about
one order of magnitude smaller than the polar O+ outflow rate. They then suggested that
there might be a substantial return flux from the magnetosphere to the low-latitude
ionosphere and also implied that this may indicate the existence of an as-yet unknown
loss process for terrestrial oxygen. Otherwise, the net escape loss over 3 billion years
may drop to about 2% of the current atmospheric oxygen content. This issue is very
important not only for magnetosphere physics but also for the field of planetology.

31

Photometric observation is a very useful approach to resolve this issue as it isdifficult to observe the escape of cold O+ ions by in-situ measurements at energies
below 50 eV without spacecraft potential control [Seki et al., 2002].
3.2.7 Short- and Long-terms Variation of O+ Outflow
Previous studies reported that the occurrence of the O+ outflow event depends
on the direction of the IMF, the geomagnetic activity, the seasonal condition, and thesolar activity [e.g., Chandler et al., 1991; Abe et al., 1993, 1996]. Chiu et al. [1990]
predicted the intensity of O+ outflow to range from 0.3 to 1 Rayleigh. Recently the
Low-Energy Neutral Atom imager (LENA) on IMAGE observed the oxygen ion burst
outflow in response to an enhancement of the dynamic pressure of the solar wind at the
moment of the passage of a shock produced by a coronal mass ejection [Moore et al.,
2001; Fuselier et al., 2001]. These observations indicated that the O+ outflow flux has
both a long- and short-term variation; therefore, a constant monitoring and simultaneous
observation of solar wind plasma are necessary to identify O+ outflow mechanisms. The
LENA imager can detect the energetic neutral atoms within the energy range of 10-750
eV [Moore et al., 2000]. The TEX instrument, however, can detect the O II emission
scattering by the thermal O+ as well as the low energy ions, i.e., it can measure O+ witha thermal energy range at which ions will return to the atmosphere due to thegravitational force once they flow upward. The TEX instrument observes during an
isometric orbit from the Earth and is able to detect the short-term variation more easily
than the IMAGE mission.
The key technology to clarify the above issues is to build bright optics in the
spectral range of EUV. A multilayer coated mirror that has a high reflectivity in the
EUV range is of particular importance [Yoshikawa et al., 1997]. The capabilities of
instruments with these technologies have been already demonstrated by the Planet-B
mission [Yoshikawa et al., 2001]. The success of the Planet-B experiment has motivatedus to attempt continuous global EUV imaging.

3.3 Instruments
Effective imaging of the near-Earth plasma distribution requires global‘snapshots’ with wide field-of-view (FOV). Lunar orbiter (SELENE) provides us with
an ideal platform that commands a panoramic view of the near-Earth situation. The

32

telescope of EUV is mounted on a two-axis gimbal system together with Telescope ofVisible light (TVIS), as shown in Figure 3.1. The pointing accuracies of the equatorial
mounting are 3.24" around the Azimuthal and 2.68" around the Elevation axes. These
correspond to the spatial resolution of 5 km on the Earth’s surface.
The UPI-TEX imager is a type of normal-incidence telescope with a split thinmetal filter, that is composed of aluminum/carbon (Al/C) and indium (In) in order todetect the resonance scattering emissions of helium ions (He II: 30.4 nm) and oxygenions (O II: 83.4 nm). The primary scientific goals are: (1) to understand a global plasma
distribution in the inner magnetosphere, and (2) to study the outflow mechanisms ofoxygen ions from the polar ionosphere. The cold plasma in the plasmasphere also showsan active response at the high magnetic activity, which is contrary to a traditional
understanding of plasmaspheric formation. Recent observations suggest that the heavyions, such as oxygen ions, could also flow out. The remote sensing at O II by TEX gives
us the first overall image of the world and will investigate the nature of the formation.
The optics in TEX consists of a multilayer-coated mirror to increase thereflectivity at 30.4 nm, a split metallic thin filter, and microchannel plates to pick up
photon events (Figure 3.2). A high-voltage power supply (HVPS) is installed at one side,
as shown in Figure 3.3, and a preamplifier unit is contained along the other side.
Survival heaters for the main body of the instrument, the HVPS, and the preamplifier(not shown in Figure 3.3) are also installed. Table 3.1 summarizes the main features of
TEX. The key points for the O II (83.4 nm) and He II (30.4 nm) imagings are (1) toeliminate the contamination from H I (Ly-α), which is emitted from the geocorona at a
maximum intensity of 10 kR [Rairden et al., 1986] and (2) to enhance the detection
efficiency at 30.4 nm.

33

Figure 3.1: Upper atmosphere and Plasma Imager (UPI) onboard the SELENE
satellite. Two telescopes are installed on the equatorial mounting (motor-controlled
azimuthal and elevation axes) gimbal. After the insertion to the lunar orbit, thesystem is open (left panels) in order to make the two telescopes ready to point to
the Earth. Pointing accuracy of the equatorial mounting is 3.24” around Azimuthaland 2.68” around Elevation axis.

Figure 3.2: Configuration of Telescope of Extreme ultraviolet (TEX) onboard the
SELENE satellite. Main optical components are a multilayer coated mirror, a split
metal thin filter, and microchannel plates. Field-of-view is separate, O II (83.4nm,
red line in the figure) and He II (30.4nm, blue line) areas.

34

Figure 3.3: A photograph of TEX without a multilayer thermal insulator. A
high-voltage power supply is installed at this side. The preamplifiers are put at the
opposite side.

Table 3.1: Principal TEX design parameters
Prime focus
Angular resolution
Spatial resolution
Maximum time resoution
Mass
Size
Open Aperture
Format
FOV
Reflectivity of mirror
Transmittance of filter
Quantum efficiency of MCP
Sensitivity

F/1.4
0.085 deg
0.089 R E
1 min
2 kg
235 mm×φ140 mm
48.6 cm2
128×128
10.88×10.88 deg2
Typ. (at 30.4 nm)
Typ. (at 83.4 nm)
18.8%
20.2%
11.2%
2.1%
7.6%
9.8%
0.014 cps/Rayleigh/bin
0.0036 cps/Rayleigh/bin

3.3.1 Mo/Si Multilayer Coated Mirror
In this section, I discuss the design and testing of a mirror coating. To realize acompact optics, a multilayer-coated mirror is designed for normal incidence reflection
with a peak reflectivity at 30.4 nm, because most materials have much higher

35

single-layer reflectivities at 83.4 nm. For example, the normal incidence reflectance ofsilicon at 83.4 nm is about 37%. The mirror substrate is made of Zerodur, and the mirror
has a focal length of 168 mm, a diameter of 120 mm, and a spherical shape to reduceaberration.
The recent study found that an Mg/SiC multilayer is a leading candidate for ahighly reflective coating at 30.4 nm [Yoshikawa et al., 2005]. Many samples were
fabricated and tested to determine optical performance and degradation underconditions of vacuum and air. The Mg/SiC coating has a high reflectivity (>40%) at
30.4 nm, but the Mg/SiC stack is degraded even with only a short time exposure to theatmosphere [Murakami et al., 2006]. Therefore, a conventional Mo/Si multilayer
coating is used in this mission. This decision was supported by the results of somesuccessful space-borne experiments in which the Mo/Si multilayers had been used
[Yoshikawa et al., 1997, 2001]. For this space mission, 20 alternating Mo-Si layers(period of 16.49 nm, Mo thickness of 4.1 nm, Si thickness of 12.39 nm) on the Zerodur
substrate of the mirror are chosen. All layers were deposited by ion-beam sputtering.
The Nikon Company manufactured the substrate and coating of the mirror.
The calibration system consists of a windowless RF excited flow lamp filledwith source gases and connected to a JOVIN YVON LHT30 30-cm scanningmonochromator. Using an atomic source with the monochromator provides radiation
limited only by the line widths of spectral lines from the light source. In the vacuum
chamber, the mirror was set on a manipulator with three orthogonal axes of translation
and a single axis of rotation aligned perpendicularly to the incident beam. The incident
and reflected beams were detected by a photon detector. The reflectivity was determinedby measuring the ratio of reflected and incident beams. This facility has a capability ofmeasuring the reflectance (and transmittance) at any position as a function of
wavelength and incident angle. This process was repeated several times for eachwavelength to enable us to compensate for long-term drifts in the source intensity.
Figure 3.4 shows the calibration system in the EUV facility.

36

Figure 3.4: Calibration setup in the EUV facility. Flight mirror was set up as
so-called “θ-2θ” configuration in order to measure reflectivity against incident
angles, points, and wavelength.
Figure 3.5 gives the average reflectivity of the flight mirror as a function of
wavelength. The reflectivities were measured at two points (30 and 50 mm from thecenter) toward every 90-degree radial direction using a 1-mm diameter beam. The linein the figure represents a theoretical curve of Mo/Si layers. At longer wavelengths, there
is difference between the calculation and measurement, but this is due to the ambiguityof the optical parameter of Si in the EUV.
Figure 3.6 shows the dependence of the reflectivity on the incident angle to themirror surface. In this case, the reflectivities were measured at three points (30, 40, and
50 mm from the center). The reflectivities at 30.4 nm have the maximum values(17.5-19.1%) at incident angles of 12-15 deg. This result indicates that the incident
photons of 30.4-nm wavelength are pass through the surface and are reflected on the
Mo/Si periodic multilayer structure. The difference of the angles exhibiting the
maximum reflectivities in three points is mainly caused by the discrepancy of the
thickness of the periodic multilayer coating. On the other hand, the reflectivities at 83.4

37

nm are slightly higher with larger incident angles and almost uniform at three points.
This result indicates that 83.4-nm photons are reflected mostly on the mirror surface(i.e., oblique photons are reflected more efficiently) and the reflectivity does not dependon the thickness of the multilayer coating. The dependence on the incident angle and thereflectivities at different points on the mirror are taken into account when calculating
the total sensitivity.
Although my testing of the mirrors shows that the Mo/Si multilayer coating
was not particularly subject to degradation by oxidation, the coating was protected using
normal procedures for handling. The mirror was stored in boxes purged by dry N2, andwhen the mirror had been installed in the telescope, the overall instrument was kept
under as dry a N2 purge as practically possible.

Figure 3.5: Reflectivity of a Mo/Si multilayer coated mirror used in TEX. The
reflectivity achieved up to 18.8%. Unwanted longer wavelength emissions can beeliminated by a metallic thin filter.

38

Figure 3.6: Dependency of the mirror reflectivity on incident angle to the mirror
surface at wavelengths of (a) 30.4 nm and (b) 83.4 nm. The reflectivites were
measured at positions of 30 mm, 40 mm, and 50 mm from the center.

3.3.2 Metallic Thin Filter
The filter material was chosen to provide high transmittance at the target
wavelengths of 30.4 and 83.4 nm, while attenuating the unwanted Ly-α (121.6 nm)
emission from the Earth’s geocorona and the interplanetary medium. The filter is set at
12.3 mm off from the focal plane and is divided into two parts (Figure 3.7). One parthas a carbon coating of 31.2 nm on an Al film (162.8 nm), which attenuates emissions
whose wavelengths are longer than 40 nm; it allows the passage of He II (30.4 nm) witha small transmittance of He I (58.4 nm). The other part transmits O II (83.4 nm) andconsists of a 292.9-nm thick layer of In. The filter is supported by a stainless steel meshhaving a wire spacing of 0.36 mm and an open area of 75%. The Al material was
permitted to oxidize in air on one side, and In was oxidized on both sides. The oxidized
coatings stabilize the material and provide the additional benefit of reducing thetransmission of Ly-α (121.6 nm).
Figure 3.7 (lower panels) compares the calibrated transmittances withtheoretical ones. In the area of Al/C, carbon reduces transmittance at longer
wavelengths, since Al itself has a fairly flat transmittance between 17 and 80 nm.
Oxidizing, as noted above, reduces the transmittance over all wavelengths, especially at
longer wavelengths. The difference between theoretical and calibrated curves indicatesthe oxidization of the Al face. I conclude that transmittance above 120 nm is below 10−9,39although no absolute calibrated value was available. The photon rejection at longer
wavelengths is good enough to eliminate strong contamination from Earth’s Ly-α
geocorona. The In filter is used as band-pass filter to prevent He I, He II, and Ly-α
emissions from contaminating O II emission. In particular, the Ly-α is so strong that theIn filter is designed to be as thick as 292.9 nm despite low transmittance at 83.4 nm.
The right panel of Figure 3.7 shows that the transmittance at Ly-α is 106-fold lower thanthat at 83.4 nm, but it is still sufficient to reduce the contamination of Ly-α.
The dependence on the incident angle of the filter is shown in Figure 3.8,
together with the theoretical curves, showing good agreement with the theory.

Figure 3.7: Transmittances of thin metallic filters. (Left) Aluminum (162.8nm)/Carbon (31.2 nm) part. This part is used for He II (30.4 nm) transmittance.
(Right) Indium (292.9 nm) used for O II (83.4 nm) transmittance.

40Figure 3.8: Dependency of filter transmittance on incident angle.

3.3.3 Microchannel Plates (MCPs) and Resistive Anode Encoder (RAE)
A three-stage MCP stack is used for photon detection. Each MCP has a circularshape with an active area of 13.9 cm2, a bias angle of 8 deg, a 12-μm pore size, an L/D(a ratio of length to pore diameter) of 40:1 and a total resistance of 125 MOhm. Three
MCPs were placed in direct contact as a back-to-back stack. An important indicator of
the MCP stack performance is the pulse height distribution (PHD) of the single photon
events. The gain and the PHD measurements were accomplished with a 30.4-nm
illumination from the EUV lamp using the monochromator. Figure 3.9 shows the PHD
with −2.3 kV and −2.5 kV applied at the MCP face. The pulse height ratio (PHR: widthof the pulse height distribution) was somewhat broader—96% at -2.3 kV and 106% at-2.5 kV—than I would expect from a five-stage MCP stack. The dark count is 0.28 cps
cm−2 at -2.3 kV and 0.37 cps cm−2 at -2.5 kV.
Figure 3.10 shows the dependence of the quantum efficiency (QE) on theincident angle to the channel axis, measured at wavelengths of 30.4 nm and 83.4 nm.
The QEs at both wavelengths are lower with small incident angles. This corresponds toa situation where the photons strike the channel wall far from the MCP input, resulting
in a number of events with low gains below the threshold of the low level discriminator.
Behind the MCPs, the resistive anode encoder (RAE) is placed to analyze theposition of photon injection (Figure 3.11). The electron cloud emitted from thethree-stage MCP consists of an order of 106 electrons and impinges on RAE. The
electron cloud is divided over four electrodes at the corners in reverse proportional tothe distance to them. Image distortion in the RAE arises mainly because the MCP

41electron clouds at the edge run off the field. Therefore, an active area of 32×32 mm2 on
the RAE is used in order to avoid the distortion.
Figure 3.12 shows the block diagram of a position analysis in TEX electronics.
The signal at each electrode is amplified by A-225 (AMPTEK). The pulse-shaper(A-206) further amplifies the signal and recognizes the event with a threshold level.
When A-206 identifies a significant signal level, it sends a “ready” signal to the mainelectronics. PH300 holds the signal peak and ADC digitalizes the height of signal. The
main electronics computes the position.
The image linearity and resolution were measured by imaging an array of
pinholes in a metal mask illuminated by 30.4 nm beam, as shown in Figure 3.13. The
ion-etching mask was placed at 8.3 mm off from the input face. The array consists of
0.5 mm or 1.0 mm holes with 5-mm distance. The pinhole mask image (left panel)clarifies none of distortion of the image within active area of 32-mm × 32-mm. Rightpanel is a cross-section of the left panel in the red line, indicates 250-μm resolution isachieved.

Figure 3.9: Pulse-height distribution of MCPs with -2.3kV and -2.5kV applied.
Pulse-height-ratio (PHR) of 96-106% is achieved.

42

Figure 3.10: Dependency of the MCP quantum efficiency (QE) on incident angle
to the channel axis. The QEs of both wavelengths are lower with small incident
angles (<~10 deg). This corresponds to a situation where the photons strike thechannel wall far from the MCP input, resulting in a number of events with low
gains below the threshold of the low level discriminator.

Figure 3.11: Photograph of Resistive Anode Encoder. Active area is 32 mm
32mm. Resistive face is made of RuO2. An active area of 32mm
center is used in order to avoid aberration near the edges.

43

32mm at theFigure 3.12: Block diagram of position analysis electronics. Preamplifiers,
pulse-shapers with discriminators, and peak-hold ICs are all AMPTEK Co.
products. This system realizes 128-bin (resolution) for each axis.

Figure 3.13: (Left) Pinhole mask image obtained by the TEX detector. I
illuminated 30.4-nm beam onto the mask. Distance between holes is 5 mm andhole’s diameters are 0.5 mm or 1.0 mm. (Right) Cross-section of red line in the left
panel. Resolution of 32 mm/128 (250 μm) is achieved.

3.3.4 Electronics and High-Voltage Power Supply
SH-3 (Hitachi) controls the equatorial mounting gimbal system and the datainterface to two telescopes as well as recognizes commands from command receiver.
The high voltage power supply installed at the side of TEX is turned on by commands

44

from SH-3.
3.3.5 Total Sensitivity
The TEX imager has a full circular field-of-view (FOV) of 10.88 deg, and the
128×128 bins corresponding to the spatial resolution of 0.09 RE on the Earth’s surface
from the lunar orbit. One-half of the FOV is used for O II (83.4 nm) detection; the otheris for He II (30.4 nm).
The signal count rate N (cps) of the emission is represented by the following
equations,

N

10 6
4

I S

Tfilter Rmirror

E MCP (cps),(3.1)

where I is the brightness of emission in Rayleigh, S is an entrance aperture (cm2), Δ is
the sold angle of identical FOV, Tfilter is the transmittance, Rmirror is the mirror’sreflectivity, and EMCP is a quantum efficiency of the MCP. The typical values of those
instrumental parameters are summarized in Table 3.1. The telescope has a total
efficiency of 0.014 cps/Rayleigh/bin for the He II (30.4 nm) and 0.0036
cps/Rayleigh/bin for the O II (83.4 nm).
Figure 3.14 shows the map of the total efficiency for all 128×128 bins on thedetector, based on the laboratory photometric calibration. I constructed the map taking
account of following points: (1) the flat field for a uniform light source calculated using
the geometry of the telescope, and the dependences of (2) the mirror’s reflectivity, (3)the filter’s transmittance, and (4) the MCPs’ QE on the incident angle. The map will be
used to convert the signal count rate (cps) to intensity units (Rayleigh).
In the following, I estimate the signal-to-noise ratio (SNR) during theobservation. The detector behaves as an ideal photon counter so that signal corresponds
to the number of photon events collected during an exposure period (texp), and noise is
the square root of signal.
Instrumental background (noise count from MCPs in the detector unit) is a
primary noise source. The MCPs manufactured by Hamamatsu Photonics Co. had anoise count by 10.4 cps/cm2 in the previous planetary mission, although the noise countin the laboratory was below 1 cps/cm2 [Yoshikawa et al., 2000a; Yamazaki et al., 2006].
This increase was due to cosmic ray and/or high-energy particle bombardments in themagnetosphere. For the lunar environment, I assume the latter component is negligible

45

but as the worst case analysis, I assume 10.4 cps/cm2 (0.0065 cps/bin) as dark count inlunar orbiter.
The second noise source is a contamination from He I (58.4 nm) and Ly-α
emissions. As stated above, Ly-α is sufficiently reduced by the metallic filter on bothsides. The intensity of Ly-α is insensitive to my estimation of SNR. In the following
discussion, I assumed and fixed the Ly-α intensity at 5 kR. Another possible
contamination is He I (58.4 nm) emission. This is fairly well reduced by the filter, but it
has the potential to contaminate the observation. The maximum geocoronal He I (58.4nm) intensity and the interstellar He I (58.4 nm) intensity are assumed to be 500
Rayleigh [Meier, 1991] and 10 Rayleigh [Yamazaki et al., 2002], respectively.
During the measurement, a detector pixel sees the EUV signal (Nsignal),
measured in cps/bin, and the background (Nbackground: instrumental background plus
contaminations from the other EUV emissions), measured in the same units. The
number of signal counts is the total observed counts minus the background as follows:
N signal
signal

N total2
total

N signal
SNR

N background2
background2

2N backgroundN signal
signal2
signal(3.2)

N signal
N signal

2
background2N background.

Figure 3.15 shows SNRs based on the above assumptions. Expected intensities
discussed in Section 3.2 are also indicated. Here “SNR=2 or 3” is set as a norm. For He
II (30.4 nm) observation, the observation mode (2-min exposure and 0.09 RE resolution)
is able to identify the main body of the plasmasphere with a high reliability (SNR=3).
The identifications of plasmapause and refilling region in the trough discussed in
Section 3.2 are feasible with a fairly good SNR. Imaging polar wind by O II (83.4 nm)
is a challenging target, but the 30-min exposure with 0.5 RE resolution allows the polarwind to be visible with SNR=2.

46

Figure 3.14: The map of the total efficiency of TEX for all 128×128 bins, basedon the laboratory photometric calibration. The map is divided into two areas: thesensitivity for He II (30.4 nm) is represented in the upper left and that for O II
(83.4 nm) is represented in the lower right, respectively. The bins where thephotons of both wavelengths arrive are shown as a black region.

47

Figure 3.15: (Upper panel) Expected Signal-to-Noise Ratios for He II (30.4 nm)
observation. The SNR is insensitive to ionospheric contamination at 121.6 nm.
Here Ly-alpha intensity is assumed to be 5kR. The intensity of He I (58.4 nm) is
assumed to be 500R, 100R (ionosphere) and 10R(interstellar) as indicated. (Lower
panel) Expected SNR for O II (83.4 nm) observation. The observation is free of
ionospheric He I and Ly-alpha contaminations in the range of 10-500R for He I
(58.4 nm) and 1k-30kR for Ly-alpha. In this plot, the intensities of He I and
Ly-alpha are 100R and 5kR, respectively.

48

3.4 Summary
The UPI-TEX instrument onboard the SELENE satellite will image near-Earthcold plasmas at O II (83.4 nm) and He II (30.4 nm) emissions. The UPI-TEX instrument
has enough performance to detect the oxygen ion outflow, the transport route from thepolar ionosphere into the magnetosphere, and the plasmasphere. The 2-D O II imagery
is expected to identify the transport routes and mechanisms of the cold O+ ions and toreveal a quantitative balance between the supply and loss, including thermal plasma. On
the other hand, the detection of He II (30.4 nm) by the TEX instrument is devoted tocontributing to conventional plasmasphere physics.

49

4 Development of High-resolution Imaging Detector using
Five Microchannel Plates and a Resistive Anode Encoder
4.1 Introduction
Microchannel plates (MCPs) are widely used to detect ions, electrons, andphotons in laboratory and space applications. One of the position sensitive anodes for
MCPs is a resistive anode encoder (RAE) [Lampton and Carlson, 1979]. In the position
sensitive detector with an RAE, the ratio of the charge signals delivered at its four
corners or edges depends only on the X and Y centroid location of the event, and so the
X-Y coordinates of each event can be determined in essentially real time with a pair ofratio circuits fed by these signals. The simplicity both in the anode itself and in theassociated electronics has led to the resistive anode becoming one of the most popular
readout methods for photon counting systems.
In a resistive anode position-sensitive system, the spatial resolution depends on
the signal-to-noise ratio (SNR) at the anode terminals and therefore depends on the
MCP gain. Here, “signal” refers to the charge produced by a single photoelectron, and issynonymous with the gain of the electron multiplier structure. The “noise” is mainly
caused by the extraneous electrical noise introduced into the position analyzing systemand the thermally generated noise. The double (so called “Chevron” or “V-stack”) or
triple (“Z-stack”) MCPs are commonly used as a method of obtaining a high gain (>106).
Ideally the output pulses should be all the same amplitude for good pulse counting. An
MCP stack, however, does not give a certain value of the gain for every photoelectron
event. The statistical variation of the gain of a V-stack or a Z-stack MCP is generally
characterized by the pulse height distribution (PHD) with typical values of FWHM of
about 80-120%. Therefore, it is desirable that the MCP stack exhibits a high and stable
gain with a narrow PHD.
The position-sensitive detectors using an RAE and a usual V-stack or Z-stack
MCP typically give spatial resolutions of the order of 100×100 to 250×250 pixels.
Lampton and Carlson [1979] reported that their sample detector consisting of two
MCPs in a V-stack configuration and an RAE achieved the resolution of 100×100 pixels.
Recently, the detector using a Z-stack MCP and an RAE was mounted in the Telescopeof Extreme Ultraviolet (TEX) onboard the SELENE satellite. The detector had an active
area of 32×32 mm and it achieved the 250-μm resolution, corresponding to the format50

of 128×128 pixels [Yoshikawa et al., 2008] (see Chapter 3).
This degree of the spatial resolution has proven satisfactory for a number ofapplications in laboratory and space flight photon counting, but the space science
community needs a higher-resolution imaging detector to be applied on the optics forthe future missions. For example, a UV spectrometer in the Mercury mission needs an
MCP detector with a spatial resolution of 512×512 pixels in order to pick up small
atmospheric signals around Mercury with a high wavelength resolution.
One of the techniques to improve the spatial resolution of the detector with an
RAE is using five MCPs in a set of V-stack and Z-stack [Firmani et al., 1982; Floryan
and Johnson, 1989]. Firmani et al. [1982] and Floryan and Johnson [1989] investigatedthe performance of the similar type of detector incorporating five MCPs arranged in aset of V- and Z-stacks with an inter-stack gap (a V-Z gap). They each measured thespatial resolution of their own detector under a wide variety of applied potentials, andfound that a negative potential applied across the V-Z gap enabled the detector toachieve the spatial resolution of the order of 500×500 pixels. It was not completely
understood, however, why applying the negative potential across the V-Z gap improved
the spatial resolution. Firmani et al. [1982] suggested that applying the negative
potential expanded the electron cloud arriving at the Z-stack and a high stable gain wasachieved due to a large number of excited channels. On the contrary, Floryan andJohnson [1989] suggested as follows: (1) electrons with sufficiently high energies toreach the Z-stack were emitted far enough up the channels in the V-stack and highly
collimated by the channels, (2) the number of excited channels at the Z-stack was
reduced, and (3) the Z-stack MCP was completely operated in the saturated mode to
exhibit a narrow PHD.
It is necessary to understand the effect of the negative potential applied acrossthe V-Z gap on the electron distribution at the Z-stack input, in order to use thistechnique for other applications and to reduce the effort at experimental measurementsfor optimization purposes. Computer simulation gives us useful information on the
distribution of the electron cloud from the MCP output [Timothy and Bybee, 1975;
Tremsin and Siegmund, 1999; Saito et al., 2007]. Zanodvorov et al. [1980] presented a“ballistic” model of the charge cloud propagation based on the differential electron
distributions at the MCP output, which had been measured in the experiments ofBronshteyn et al. [1979]. In this approach, no interactions between the electrons in the51

cloud are taken into account, i.e., the electric field is assumed to be constant rather thanself-consistent.
In this study, I use a similar ballistic approach for calculating the distribution of
the electron cloud across the V-Z gap. The result gives us important information on theeffect of the negative potential across the V-Z gap. Then, to confirm the result of thecalculation, experiments are described to investigate the dependency of the spatial
resolutions and the PHDs on the applied negative potential. This work has a direct
impact on the way in which MCPs can be optimized to satisfy future mission
requirements on the spatial resolution, and enhances the capabilities for the mission
performance.

4.2 Calculation
The statistical nature of the electron avalanche, originating at the output of the
MCP channels, inhibits a precise calculation of the charge cloud evolution. The angles
and velocities of the output electrons exhibit large variations depending not only on the
MCP intrinsic parameters (such as the length-to-diameter ratio L/D, the secondary
electron emission coefficient of the semi-conductive layer, the end spoiling, etc.), butalso on the mode and sometimes even the history of operation. However, below I
suggest some approximation which enables me to simulate the electron clouddevelopment in the gap between the V- and Z-stacks (V-Z gap) and obtain some
information on the size of the electron cloud at the input surface of the Z-stack MCP.
Following the same approach as the past studies [Zanodvorov et al., 1980;
Tremsin and Siegmund, 1999], the trajectories of all electrons in the cloud are assumed
to be purely ballistic, and the differential distribution function at the output of the MCP,
measured by Bronshteyn et al. [1979], is used. For simplicity, it is assumed further that(1) the original size of the charge cloud at the V-stack output, determined by the numberof excited channels (~6.5 channels [Wiza et al., 1977; Tournianski et al., 2004]), is
negligible, that (2) no interactions occur between the electrons in the cloud, and that (3)the bias angle of the output MCP is 0 deg and then the electron distribution is symmetricto the channel axis. Figure 4.1 shows the coordinate system for my model. Using theelectron distribution function, it is possible to calculate the ballistic trajectories of allelectrons and to estimate the spatial distribution of the cloud at any distance from the
MCP output.

52

Figure 4.1: Schematic diagram of the coordinate system used for the calculation of
the charge cloud spreading. Potentials accelerating electrons from the V-stacktoward the Z-stack are considered positive.
I use one of the experimental distributions of Bronshteyn et al. [1979], which
was measured using a single MCP with an L/D of 60:1, 20-μm channels, and an applied
voltage of 1.4 kV. The electron distribution function at the MCP output can beapproximated by the following function [Zanodvorov et al., 1980]:
f (E, ) e

2

/2 2

2

A( )e ( E E0 ) / 2 0

2

2

2

B ( )e ( E E1 ( )) / 2 1 ( ) ,(4.1)

where E and θ are the electron output energy and angle, respectively, and τ = tanθ. The
unknown parameters in this distribution were chosen to fit the experimental data
[Bronshteyn et al., 1979] as follows:
E0 = 3 eV, σ0 = 12 eV, Δ = 0.2,
A(τ) = 50.398e7.3822τ, B(τ) = 25.992e9.8642τ,-6.9515τ

E1(τ) = 314.45e(4.2)-9.7077τ, σ1(τ) = 147.62e.

Figure 4.2 shows the distribution function f(E, θ) at various constant values of θ. The
equation of motion of an electron in the cloud at the moment t can be written as,r

2E
t sin ,
me

z

2E
t cos
me

(4.3)
qeV g t 2
2me D

53

,

(4.4)

where r is the radial distance from the channel axis, z is the distance from the V-stack
output, D is the distance between the V- and Z-stacks, Vg is the accelerating bias in the
V-Z gap, and qe and me are the electron charge and mass, respectively. In this work,
potentials accelerating electrons from the V-stack toward the Z-stack are considered
positive. Using Eq. (4.3) and (4.4), I obtain the radial distance r0 at the Z-stack input (z= D).
Using the electron distribution function (4.1) (summarized in Figure 4.2), I
calculated r0 of all electrons and estimated the cloud spatial distribution at the input
surface of the Z-stack with various values of Vg. The radial distributions of electrons,arriving at the Z-stack input with sufficient energy to excite secondary electrons (> 10eV [Martin and Bowyer, 1982; Siegmund et al., 1988]), are calculated with theparameters of D = 2.5 mm and Vg = 0, -100, and -200 V. Figure 4.3 shows the result. In
Figure 4.3, the total number of electrons arriving at the Z-stack for each Vg isnormalized. Applying the negative inter-stack potential clearly reduced the size of theelectron cloud. This means that the outer component of the electron cloud consisting of
the low energy electrons emitted at large angles (as shown in Figure 4.2) is mainly
eliminated. Therefore, the result of the calculation indicates that the negative inter-stackpotential allows the Z-stack MCP to be operated in the completely saturated mode and anarrow PHD will be achieved, as suggested by Floryan and Johnson [1989].

Figure 4.2: Energy distribution curves of electrons at MCP output at fixed output
angles, measured by Bronshteyn et al. [1979]. The electrons with high energy areconcentrated in small output angles.

54

Figure 4.3: The distribution of the electron cloud at the Z-stack input calculatedwith the ballistic model. The total number of electrons arriving at the Z-stack foreach Vg is normalized. Applying the negative inter-stack potential clearly reduced
the size of the electron cloud, from the FWHM of 0.66 mm to 0.21 mm.

4.3 Measurement
The experimental approach was also performed to check whether the result
from the calculation is realistic. The prototype detector used in the measurements
consists of an input window made of synthetic silica, a Cs2Te photocathode, five MCPs
in a set of V- and Z-stacks, and an RAE. The photograph and the schematic geometry ofthe detector are shown in Figure 4.4 and Figure 4.5, respectively. The photocathodeconverts the incoming photons into photoelectrons. The photoelectron is accelerated
towards the MCP in the cathode gap. The photoelectrons impact the walls of the MCP
channels causing a charge avalanche, giving an overall charge multiplication of ~2×107.
This charge cloud is drifted from the MCP output to the RAE, and its centroid position
can be determined. The four output charge signals from the RAE are amplified andshaped by four preamplifiers. The discriminator discards signals which are too small tobe properly processed and signals resulting from simultaneous or near-simultaneous
events on the RAE. For valid signals, the centroid position of the electron cloud is
calculated by measuring the peak voltage of the pre-amplifiers for each corner A, B, C,

55

and D of the RAE. The following algorithm is then used to obtain its position in the
X-Y coordinates:XA B,
A B C D

(4.5)

Y

B C
.
A B C D

(4.6)

Figure 4.4: A photograph of the prototype detector. The detector head consists ofan input window, a photocathode, five MCPs in a set of V- and Z-stacks, and a
resistive anode. Inside the input window, the components are sealed in vacuum.
The diameter of the active area is 25 mm.

56

Figure 4.5: A schematic diagram of the prototype detector. The detector requires
five separate potentials to operate: the potential between the photocathode and the
input surface of the V-stack (Vp), that across the V-stack (Vv), that between the
output surface of the V-stack and the input surface of the Z-stack (Vg), that acrossthe Z-stack (Vz), and that between the output surface of the Z-stack and theresistive anode (Va).
The MCPs incorporated in the detector is arranged in a set of V-stack (twoplates) and Z-stack (three plates) with a gap of 2.5mm, as shown in Figure 4.5. EachMCP has a circular shape with an active diameter of 25 mm, a bias angle of 12 deg, a
10-μm channel diameter, and an L/D (a ratio of length to channel diameter) of 40:1. The
resistances of the V-stack and the Z-stack are ~90 MΩ and ~120 MΩ, respectively. The
RAE has an active area of 23×23 mm2. In this measurement, however, I use a circular
active area of 16-mm diameter defined by an aperture mask placed in contact with theoutput surface of the Z-stack to avoid the distortion. The detector requires five separate
potentials applied as shown in Figure 4.5: the potential between the photocathode andthe input surface of the V-stack (Vp), that across the V-stack (Vv), that between the
output surface of the V-stack and the input surface of the Z-stack (Vg), that across theZ-stack (Vz), and that between the output surface of the Z-stack and the resistive anode(Va). It is possible to operate the detector with a wide variety of combinations of these
five potentials. In this work, potentials which accelerate electrons from the57

photocathode toward the anode are considered positive, and those which decelerate
electrons in this direction are considered negative. The applied voltages of Vp and Va
were constantly ~200V and ~500V, respectively, for all the measurements.
I measured the spatial resolution of the detector using the USAF 1951
resolution test target. My experimental setup is shown schematically in Figure 4.6. The
light for all tests was 633 nm light produced by an intensity-stabilized laser. The test
target was illuminated uniformly through diffusers, and the image of the test pattern was
projected onto the photocathode by a relay lens. Optical focus of the test-pattern image
on the photocathode was obtained by taking a series of spatial resolution measurements(described below) at various distances between the relay lens and the detector, and bydetermining the position where the maximum resolution was achieved. Figure 4.7
shows a typical image of the test target and a cross-sectional histogram at 6.3 lp/mm
across it, obtained with Vv of 1420V, Vg of 0V, and Vz of 1710V applied. The contrast
was calculated from the images of the test target with a simple formula, (Max – Min) /(Max + Min). Analyzing the profiles of the test-pattern images allows the resolution of
the detector to be determined.

Figure 4.6: A schematic diagram (upper) and a photograph (lower) of the setup of
the measurement. The test target is illuminated uniformly. The relay lens transmits
the image of the test target to the photocathode of the detector.

58

Figure 4.7: (a) A test target image obtained by the experimental detector. The
series of patterns encircled by dotted line are well focused. (b) A cross section
profile of (a) taken along the white line at 6.3 lines/mm. The contrast at 6.3
lines/mm is 0.53.
In this work, I recorded a number of test target images and pulse height
distributions obtained with different combinations of applied potentials. I investigated
the relationship between the negative inter-stack potential, the spatial resolution, and thepulse height distribution with the experimental approach.

4.4 Results and Discussion
In the first measurement, the dependence of the spatial resolution on the gain
and the PHD was examined. I measured the PHDs and spatial resolutions with various
potentials applied, and the results are shown in Figure 4.8. Each applied potential is also
displayed. As seen in Figure 4.8a, the peak gain increases and the distribution becomeswider as the applied potential increases. In Figure 4.8b, while the contrast linearly
increases with the gain below 1.6×107, there is no correlation between the contrast andthe gain above 1.6×107. This result means that the spatial resolution mainly depends on
the number of events with low gains (< ~0.4×107 in this measurement) which have thelow SNR. In this measurement, therefore, it was confirmed that the spatial resolution of
the detector surely depends on the gain and the PHD of the MCPs.

59

Figure 4.8: (a) Pulse height distributions (PHDs) with various high-voltage
applied. The gain increases and the distribution becomes wider as applied voltage
increases. (b) Dependence of the contrast at 6.3 lines/mm on the gain. The PHD
corresponding to each point is shown in (a) in the same color. While the contrast
linearly increases with the gain below 1.6 107, there seems to be no correlation
between the contrast and the gain over 1.6 107.
In the second measurement, the effect of the negative inter-stack potential (Vg)on the PHD was examined. Figure 4.9 shows the PHDs with Vg of 0V, -76V, and –205V
applied. The applied potentials across the V-stack (Vv) and across the Z-stack (Vz) arecontrolled appropriately so that each PHD has approximately the same gain at the peak.
The PHD clearly got narrower as Vg increased. The full width at half maximum(FWHM) of the PHD (PHR: pulse height ratio) was excellently improved from 165%
(Vg = 0V) to 38% (Vg = –205V). This narrowing of the PHD introduced by applying thenegative inter-stack potential is consistent with the prediction from my calculation(described in Section 4.2).
Finally, the dependence of the spatial resolution on Vg was examined. Figure
4.10 shows modulation transfer functions with Vg of 0V, -76V, and –205V applied. The
spatial frequencies where the contrasts equal to 0.33 are also displayed. Applying thenegative inter-stack potential improved the spatial resolution by ~14%. This result is
consistent with that obtained in the second measurement. In Figure 4.10, furthermore,
there seems to be no difference between the spatial resolutions with Vg of –76V and–205V. This characteristic can be explained by the result of the first measurement, i.e.
nearly the same numbers of events with low gains below ~0.4×107 were achieved. As a60

result, the spatial resolution of 54 µm, corresponding to 400×400 pixels in my detector,
was achieved by using V- and Z-stacks and applying the negative inter-stack potential.
The achieved resolution is lower than that of the order of 500×500 pixels
reported by Firmani et al. [1982] and Floryan and Johnson [1989]. This difference
would mainly result from the deficiency of optimizing the potentials applied between
the photocathode and the V-stack input (Vp) and between the Z-stack output and theRAE (Va). Floryan and Johnson [1989] reported that the resolution depends on Va andespecially on Vp. They achieved the highest resolution of 500×500 pixels with Vp of
700V, while it was 200V in my measurement. It is necessary to measure the spatial
resolution of the detector with higher Vp for the future missions. Nevertheless, this study
revealed the effect of the negative potential applied across the inter-stack gap on the
distribution of the electron cloud, the PHD, and the spatial resolution by the calculationand the experiments. Using the same technique I can optimize the characteristics of theposition sensitive detectors incorporating an RAE to the future missions.

Figure 4.9: Dependency of the pulse height distribution on the potential between
the output surface of the V-stack and the input surface of the Z-stack (Vg). The
distribution becomes sharper as Vg increases. The applied potentials across theV-stack (Vv) and across the Z-stack (Vz) are controlled appropriately so that each
PHD has approximately the same gain at the peak.

61

Figure 4.10: Dependency of modulation transfer function (MTF) on Vg. The PHD
with each Vg applied is shown in Figure 4.9 in the same color. Applying negative
Vg increases the spatial resolution by ~14%. There seems to be no difference
between the spatial resolutions with Vg of –76V and –205V.

4.5 Summary
I have investigated the effect of the negative potential applied across theinter-stack (V-Z) gap on the PHD and the spatial resolution by means of calculation andexperiments. The calculation with a simple ballistic model showed that the negativeinter-stack potential (Vg) reduced the size of the electron cloud by ~70%. The result
suggests that under such a condition the Z-stack MCP is operated in the completely
saturated mode and exhibits a narrow PHD. On the other hand, I measured the PHDs
and the spatial resolutions of the prototype detector with various potentials applied, andrevealed the following points.
(1) The spatial resolution depends on the PHD, especially on the number of theevents with low gains. This result is explained by the dependency of the resolution on
the SNR of the output pulses.
(2) Applying the negative Vg reduces the width of the PHD from PHR=165%(Vg=0V) to 38% (Vg=-205V). This result is consistent with that of the calculation.
(3) Applying the negative Vg improves the spatial resolution by ~14%. This result is
explained by the results of (1) and (2).

62

Finally, the spatial resolution of 54 µm, corresponding to 400×400 pixels in my detector,
was achieved. I can optimize the technique using five MCPs and a RAE and apply it to
the future missions. Further measurements, however, with a higher potential applied
between the Z-stack output and the RAE are necessary to achieve a higher spatial
resolution.

63

5 Concluding Remarks
The present study has confirmed that a remote sensing with EUV techniques is
a promising means to explore the Earth’s plasmasphere. The plasmasphere and its outerregion are full of unsolved problems which are difficult to clarify by in-situ
observations alone. Global EUV imaging enables us to shed new light to the physics ofthese regions. This thesis details the analysis on the plasmasphere images obtained by
the EUV imager on the IMAGE satellite, the instrumentation and calibration of the
UPI-TEX imager on the SELENE satellite, and the development of a high-resolution
detector for future missions.
I statistically examined the plasmapause response to the southward turning of
the IMF using sequential global images of the plasmasphere. The extreme ultraviolet(EUV) imager on the IMAGE satellite clearly observed inward motion of theplasmapause driven by the southward turning of the IMF. I surveyed the EUV data inthe 2000–2001 period and found 16 events. Using the sequential EUV images, I haveshown that the plasmapause response to the southward turning of the IMF takes 10–30
min, and average 18 min. It is consistent with the timescale derived from theionospheric observations. Therefore the result indicates that the electric field penetrates
from the magnetopause to the inner magnetosphere through the ionosphere. Chapter 2
deals with this topic.
The SELENE satellite was launched in 2007 and sent to the Moon. From thelunar orbit, the UPI-TEX will image the Earth’s plasmasphere and the polar wind at O II
(83.4 nm) and He II (30.4 nm) emissions. The instrumentation and calibration of the
UPI-TEX imager has been presented in Chapter 3. The UPI-TEX imager has a total
efficiency of 0.014 cps/Rayleigh/bin for the He II (30.4 nm) and 0.0036
cps/Rayleigh/bin for the O II (83.4 nm). For imaging the main body of the plasmasphereat He II (30.4 nm) emission, the UPI-TEX imager achieves the temporal resolution of 2
min and the spatial resolution of 0.09 RE with a high reliability (SNR=3). Imaging of theplasmasphere will improve our understanding of the distribution of plasma in the innermagnetosphere and the forces that control it. In addition, the images will be useful forestablishing the global context that can enhance the interpretation of in-situ
measurements. On the other hand, the 30-min exposure with 0.5 RE resolution allows us
to image the polar wind by O II (83.4 nm) emission with SNR=2. The remote sensing at

64

O II by UPI-TEX will give the first global image of the polar wind and make significant
contributions to our understanding of the transport routes and mechanisms of the cold
O+ ions.
In Chapter 4, I have presented the development of a high-resolution detector
using five MCPs in a set of V- and Z-stacks and an RAE for the future missions. The
effect of the negative potential applied across the inter-stack gap on the PHD and thespatial resolution has been investigated. As a result, the calculation with a simple
ballistic model showed that the negative inter-stack potential reduced the size of theelectron cloud by ~70%. The experiments furthermore revealed that applying thenegative inter-stack potential reduced the width of the PHD by ~80% and improved thespatial resolution by ~14%. Finally, the spatial resolution of 54 µm, corresponding to
400×400 pixels in my detector, was achieved. Based on these results, I can optimize andapply this technique to the future missions.

Remaining Problems
In Chapter 2, using sequential EUV images of the plasmasphere I have shownthat the plasmapause response to the southward turning of the IMF takes 10-30 min,with an average of 18 min. However, the result should be re-investigated using a more
completed data set of EUV images because only 16 events were found in my search dueto the limited EUV data. The IMAGE satellite orbits near the Earth, with the apogee of~8 RE. As such, it cannot always monitor the whole plasmasphere, and there is always
an interval of at least several hours between images, which makes it difficult to estimate
the time lag between the IMF arrival at the magnetosphere and plasmasphere evolution.
In addition, a higher temporal resolution than that of the EUV instrument (10 min) is
preferable to compare the plasmapause motion with the changes of the solar windobserved at the resolution of ~1 min.
The observation by the UPI-TEX imager improves the number of events
because SELENE provides the best platform to command a panoramic view of thewhole plasmasphere. I can re-investigate the plasmapause response to the solar windusing much more sequential images obtained by UPI-TEX. On the other hand, thetemporal resolution is not effectively improved with respect to observing theplasmapause motion. The typical velocity of the plasmapause caused by the enhanced65

convection is ~0.02 RE/min, while the UPI-TEX imager has only a spatial resolution of
0.1 RE. Therefore, the spatial resolution of 0.02 RE, e.g. corresponding to 500×500
pixels in the field of view of 10 RE, is required to distinguish the plasmapause motion at
a temporal resolution of the order of ~1 min. I have developed a high-resolution
detector using V- and Z-stack MCPs and an RAE and achieved a spatial resolution of
400×400 pixels, as described in Chapter 4. Further examination of the dependency ofthe spatial resolution on the potential applied between the Z-stack output and the RAE
is necessary to achieve a higher spatial resolution.

66

Acknowledgements
The author wishes to express his sincere appreciation to Dr. I. Yoshikawa of the
University of Tokyo for his guidance during the master’s course and for stimulating
discussions. The author would like to thank Dr. A. Yamazaki of Institute of Space andAstronautical Science for offering valuable suggestions and guidance throughout thepresent work as well as supporting the calibration of UPI-TEX. The author is grateful to
Mr. K. Yoshioka for his continuing interest and encouragement. The author also thanks
Dr. N. Iwagami, Professor M. Hoshino, Professor M. Hirahara, Dr. T. Yokoyama, andProfessor M. Nakamura for providing important comments.
The author appreciates the help received from Ms. M. Hirai with analyzing the
IMAGE EUV data. The authors would like to thank Dr. B. R Sandel for his assistance inthe processing of EUV data available at http://euv.lpl.arizona.edu/euv/. The author also
thanks Dr. N. Ness, Dr. C. Smith and the ACE science center for ACE MAG data and Dr.
D. McComas for ACE SWEPAM data. The author is grateful to all of the members of
the UPI team and the SELENE project team for their effort in achieving a successful
mission.
The author also wishes to thank Dr. S. Kameda, Mr. J. Ono, Mr. F. Ezawa, Mr.
T. Toyota, Mr. G. Ogawa, Dr. M. N. Nishino, Dr. S. Ohtsuki, and other members of
Yoshikawa Lab, Iwagami Lab, and Nakamura Lab for their emotional assistance andfriendship.

67

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