ISSN 0919-8296

ANNUAL REPORT [APRIL 2004 — MARCH 2005] ICRR UNIV. OF TOKYO

2004

INSTITUTEFORCOSMIC RAY RESEARCH
UNIVERSITY OF TOKYOANNUAL REPORT(APRIL 2004 – MARCH 2005)

Editorial Board
TAKEUCHI, Yasuo
MIYOKI, Shinji­c Institute for Cosmic Ray Research, University of Tokyo
5-1-5, Kashiwanoha, Kashiwa, Chiba 277-8582 Japan
Telephone: (81)4-7136-3102
Facsimile: (81)4-7136-3115
WWW URL: http://www.icrr.u-tokyo.ac.jp

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Universal Academy Press, Inc.
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E-Mail: general@uap.co.jp
WWW URL: http://www.uap.co.jp

TABLE OF CONTENTS
Preface1

Research Divisions
Neutrino and Astroparticle DivisionHigh Energy Cosmic Ray Division
Astrophysics and Gravity Division
Observatories and a Research Center

40

Norikura Observatory
Akeno Observatory
Kamioka Observatory
Research Center for Cosmic Neutrinos
Appendix A.

ICRR International Workshops

49

Appendix B.

ICRR Seminars

50

Appendix C.

List of Publications — 2004 fiscal year

51(a) Papers Published in Journals(b) Conference Papers(c) ICRR Report
Appendix D.

Doctoral Theses

55

Appendix E.

Public Relations

55(a) ICRR News(b) Public Lectures(c) Visitors
Appendix F.

Inter-University Researches

57

Appendix G.

List of Committee Members

58(a) Board of Councillors(b) Executive Committee(c) Advisory Committee
Appendix H.

List of Personnel

59

PREFACE
Starting from April 2004, Japanese national universities have become independent corporations and were set up outside of the government organizations. We therefore are supposed to have much more autonomy. However the new mechanism of the budget allocationfrom the government to the universities is inadequate to initiate big and long term scientificprojects to be conducted in the universities. It is also known to be difficult to increase thenumber of scientific staffs for the long term programs. A new contrivance for the budget
allocation from the government is urgently needed in order to keep activities in basic sciencein Japanese universities and institutions.
The Institute for Cosmic Ray Research, the inter-university institute, is closely associatedwith Japanese universities and foreign institutes. The experimental facilities located at theinstitute are jointly established and used by about 350 physicists. We hope that we willcontinue to keep good collaborations and to keep providing important results on cosmic rays
physics and related subjects although the university system has been drastically changed.
The institute has been re-organized into three research divisions; Neutrino and Astroparticle Division, High Energy Cosmic Ray Division, and Astrophysics and Gravity Division. In each division, a few different experiments are conducted.
The next big project of the institute is the gravitational wave detector (LCGT), by which
the direct detection of the gravitational wave (GW) is aimed. The detection of the GW will
prove the correctness of the Einstein‘s theory of general relativity and can provide the detailsof the dynamic feature of the universe.
The study of the cosmic ray had often played a leading role when the particle physics had
made a significant development. The pions and muons were discovered in the cosmic rays
and those discoveries had brought a fruitful study of the elementary particles using particle
accelerators. The discovery of the neutrino mass has again brought a triumph to the cosmic
ray research and it has opened up a new field.
At Kamioka underground observatory, the Super-Kamiokande experiment is continuously producing interesting results. There are some other projects in Kamioka, for example,
dark matter experiments and geo-physics experiments, which make use of advantages of theunderground environments. The construction of the prototype gravitational wave antenna isalso in progress.
In 2003, the long-waited Telescope Array project has been approved and its construction
has begun. The existence of the highest energy cosmic rays beyond the ZGK cut-off indicatedby the AGASA group is a big puzzle. If it is confirmed, it may suggest a new physics. Thereis a steady flow of data from the experiment at Yangbajing (Tibet), and from the cosmic
gamma-ray telescope (CANGAROO) deployed over the desert of Woomera (Australia).
The underground physics and the ground based cosmic ray measurements overseas are
the back-born of the institute.

Yoichiro Suzuki,
Director of ICRR

Fig. 1. The ICRR building at Kashiwa, Chiba, Japan.

Fig. 2. The inner detector of SuperKamiokande-I during the initial construction. The purified water is under filling.
One can clearly see the 20-inch diameter photomultiplier tubes and a person
working on a boat.

Fig. 3. The system of four imaging atmospheric Cherenkov telescopes of
10 m diameter of CANGAROO projectfor detection of very high energy
gamma-rays The whole system is inoperation since March 2004 in
Woomera,South Australia..

Fig. 4. Tibet-III air shower array(37000 m 2) at Yangbajing, Tibet(4300 m in altitude).

Fig. 5. Air fluorescence telescopes (left)
and a scintillator surface detector (right)of the Telescope Array experiment under construction in Utah, USA for the
study of extremely high energy cosmic
rays.

Fig. 6. Cryogenic mirror suspension system for Large
Scale Cryogenic Gravitational Wave Telescope.

Fig. 7. Wide-view telescope of 2.5 m diameter (left telescope) in Arizona, USA for the Sloan Digital Sky Survey project.

Fig. 8. A public lecture held byResearch Center for Cosmic
Neutrinos.

Organization
Board of Councillors

Faculty Council

Research Division
Neutrino and Astroparticle DivisionHigh Energy Cosmic Ray Division
Astrophysics and Gravity Division

Director
Vice-Director

Executive Commitee

Facilities
Norikura Observatory
Akeno Observatory
Kamioka Observatory
Research Center for Cosmic Neutrinos

Advisory Comittee

Administration Division

Number of Staff Members in 2004
Associate

Research

Technical

Resea rch

Professor
4
5
3
0
12

Associate
7
6
2
0
15

Staff
3
13
1
0
17

Fellows
8
17
8
0
33

Professor
Neutrino and Astroparticle Div.
High Energy Cosmic Ray Div.
Astrophysics and Gravity Div.
Administration
Total

3
2
3
0
8

Administrators
andSecretaries
8
3
0
11
22

FY 1999–2004 Budget
Personnel
expenses
Nonpersonnel
expenses
Total

1999

2000

2001

2002

2003

2004

569 050589 879

418 475

460 332

434 874444 498

1 411 063

1 423 789

1 518 584

1 518 065

1 785 449

2 054 950

1 980 113

2 013 668

1 937 059

1 978 397

2 220 323

2 499 448
(in 1 000 yen)

33
46
17
11
107

1

RESEARCH DIVISIONS

Neutrino and Astroparticle Division
Overview
Super-Kamiokande Experiment
K2K Experiment
XMASS ExperimentHigh Energy Cosmic Ray Division
Overview
CANGAROO-II & -III Project
Study of the Most Energetic Cosmic Rays — AGASA Collaboration
Seeking for an Origin of Super-GZK Events — Telescope Array Experiment
Experimental Study of High-energy Cosmic Rays by the Tibet Air Shower Arrayin the Tibetan Highlands (The Tibet ASγ Collaboration)
Developement of Wide Angle High Resolution Detector
Chacaltaya Observatory of Cosmic Physics

Astrophysics and Gravity Division
Overview
TAMA Project
LCGT Project
Construction of CLIO at Kamioka
Sloan Digital Sky Survey
Theory Group
Implications of the Curvaton on Inflationary Cosmology
Spacetime Symmetries on the Supermanifolds
EWIMP Dark Matter Detections
Significant Effects of the Second KK Particles on LKP Dark Matter Physics
Hadronic Electric Dipole Moments in Supersymmetric Grand Unified Theories
Big-Bang Nucleosynthesis and Hadronic Decay of Long-Lived Massive Particles

2

NEUTRINO AND ASTROPARTICLE DIVISION
Overview
The Neutrino and Astroparticle division consists ofKamioka Observatory and Research Center for Cosmic Neutrinos (RCCN) and performs neutrino physics and astroparticle physics. Kamioka Observatory is located at Kamioka,
Hida-shi, Gifu prefecture. The observatory has undergroundexperimental sites at 1000 m underground and research facilities on the ground.
The main experiment at Kamioka Observatory is SuperKamiokande (SK). The SK detector was built from 1991 to
1995 and data taking was started in 1996. In 1998, the SK
collaboration announced the discovery of neutrino oscillations
using atmospheric neutrinos. Another evidence for neutrino
oscillations was found in 2001 using solar neutrinos by comparing solar neutrino data from SK and SNO. The SK detector has been used as the far detector of the artificial neutrino
beam experiment (K2K) performed from 1999 and neutrino
oscillation was confirmed in 2002. After the discovery of neutrino oscillations, detailed studies of neutrino oscillations havebeen performed at SK. In the analysis of atmospheric neutrinos, oscillation parameters, oscillation mode (whether νμ oscillates to ντ or νsterile ) and LE dependence of oscillation
have being studied. In the solar neutrino analysis, oscillation
parameters has been determined in 2002 and now searchingfor energy spectrum distortion and day/night time variation
expected from the obtained solution. SK has been monitoring
neutrinos from supernova burst. If a supernova burst occurs at
the distance to the center of our galaxy, SK is able to detect
about 8,000 neutrino events. SK is also searching for nucleon
decay as the direct evidence of Grand Unified Theories. A
high intensity neutrino beam experiment using J-PARC (T2K)
is expected to start in 2009 and SK detector will be the far
detector of the experiment. High precision measurement ofoscillation parameters and the third oscillation pattern (fromthe third neutrino mass eigenstate to the first eigenstate) will
be investigated by T2K.
Another activity of the Neutrino and Astroparticle division is a multi-purpose experiment using liquid xenon aiming
at the detection of cold dark matter, neutrino absolute massusing neutrinoless double beta decay, and low energy solar
neutrinos. An R&D study for the liquid xenon detector is being performed at Kamioka observatory.
Recent progress of research activities in the Neutrino andAstroparticle division is presented.

Super-Kamiokande Experiment[Spokesperson : Yoichiro Suzuki]
Kamioka Observatory, ICRR, Univ. of Tokyo, Gifu, 5061205
In collaboration with the members of:
Kamioka Observatory, ICRR, Univ. of Tokyo, Japan; RCCN,ICRR, Univ. of Tokyo, Japan; Boston Univ., USA; BNL,

USA; Univ. of California, Irvine, USA; California State
Univ., Dominguez Hills, USA; Chonnam National Univ,, Korea; Duke Univ., USA; George Mason Univ., USA; Gifu
Univ., Japan; Univ. of Hawaii, USA; Indiana Univ., USA;
KEK, Japan; Kobe Univ., Japan; Kyoto Univ., Japan; LANL,
USA; Louisiana State Univ., USA; Univ. of Maryland, USA;
Univ. of Minnesota, USA; Miyagi Univ. of Education, Japan;
SUNY, Stony Brook, Japan; Nagoya Univ., Japan; Niigata
Univ., Japan; Okayama Univ., Japan; Osaka Univ., Japan;
Seoul National Univ., Japan; Shizuoka Seika College, Japan;
Shizuoka Univ., Japan; Sungkyunkwan Univ., Korea; Tohoku
Univ., Japan; Univ. of Tokyo, Japan; Tokai Univ., Japan;
Tokyo Inst. of Tech., Japan; Inst. of Experimental Physics,
Poland; Univ. of Washington, USA.

Introduction
Super-Kamiokande(SK) is a large water Cherenkov detector, located 1000 m underground in Kamioka mine, Japan.
50 kton of pure water is contained in a stainless steel tankof 39.3 meters in diameter and 41.4 m in height. SK took
data from April 1996 to July 2001 (SK-I phase) using 11,146
20-inch photomultipliers(PMTs) for inner detector and 1,885
8-inch PMTs for outer detector. After the accident in November 2001, the detector was reconstructed in 2002 using about
5200 20-inch PMTs. The detector has been running as thesecond phase of the experiment(SK-II) since December 2002.
In this report, those results are described.

Atmospheric neutrinos
Cosmic ray interactions in the atmosphere produce neutrinos. The prediction of the absolute flux has an uncertainty
of at least 20%. However, the flavor ratio of the atmospheric neutrino flux, νμ  ν̄μ  νe  ν̄e , has been calculated to an accuracy of better than 5%. Another importantfeature of atmospheric neutrinos is that the fluxes of upwardand downward going neutrinos are expected to be nearly equalfor Eν (a few GeV) where the geomagnetic effect on primarycosmic ray is negligible.
SK-I observed 12,180 fully-contained (FC) events and 911
partially-contained (PC) events during 1489 days of data taking. FC events deposit all of their Cherenkov light in the innerdetector, while PC events have exiting tracks which deposit
some Cherenkov light in the outer detector. The neutrino
interaction vertex was required to have been reconstructedwithin the 22.5 kiloton fiducial volume, defined to be  2 mfrom the PMT wall.
The FC events were classified into “sub-GeV” (Evis 
1330 MeV) and “multi-GeV” (Evis  1330 MeV) samples.
The numbers of observed and predicted events for sub- andmulti-GeV energy regions in SK are summarized in Table 1.
The prediction is based on the recent precise measurements of

MC

Sub-GeV
e-like3353 2879.8
μ -like3227 4212.8
R 0.658 0.016 0.035
Multi-GeV
e-like746
680.5
μ -like(FC+PC) 1562 2029.5
R 0.70200 032
030 0.101

50
45
40
35
30
25
20
15
10
5
0
-1

-0.5

0

0.5

Number of Events

100
80
60
40

Number of Events

Number of Events

-0.5

0

0.5

140
120
100
80
60
40

-0.5

0

cosΘ

0.5

1

-0.5

0

cosΘ

300
200

350

-0.5

0

0.5

1

cosΘ
Multi-GeV μ-like + PC

60
40

0.5

1

-0.8 -0.6 -0.4 -0.2 0cosΘ
Upward Through Going μ
400
350
300
250
200
150
100
50
0
-1 -0.8 -0.6 -0.4 -0.2 0cosΘ

Fig. 2. The zenith angle distributions for multi-ring sub-GeV μ -like(upper left) and multi-ring multi-GeV μ -like (lower left) samples.
The zenith angle distributions of upward stopping muons (upper
right) and upward through-going muons (lower right). The solid
histograms show expectations without neutrino oscillations. The
dashed histograms show the expected flux for the νμ  ντ oscillation with sin2 2θ  10 and Δm2  21  10 3 eV2 .

300
250
200
150
100
50

20

80

400

0
-1

1

Number of Events

Number of Events

0
-1

500

100

cosΘ
Multi-GeV e-like

0-1

Sub-GeV μ-like

Sub-GeV e-like

100

0
-1

1

cosΘ
Multi-GeV Multi-ring

20

450
400
350
300
250
200
150
100
50
0
-1

120

Upward Stopping μ20

Number of Events

DataR) ratio

Number of Events

Table 1. Summary of the atmospheric μ edata μ eMC (measurement.

Number of Events

3
Sub-GeV Multi-ring

0
-1

-0.5

0

0.5

1

cosΘ

Fig. 1. The zenith angle distributions for sub-GeV e-like, sub-GeV
μ -like, multi-GeV e-like and multi-GeV (FC+PC) μ -like events.
cosΘ 1 means down-going particles. The solid histograms show
the MC prediction for the no neutrino oscillation case. The dashedhistograms show the Monte Carlo prediction for νμ  ντ oscillations with sin2 2θ  10 and Δm2  21  10 3 eV2 .

primary cosmic rays by BESS and AMS and a three dimensional calculation of the neutrino flux by Honda et al. The
hadronic interaction model of cosmic rays is also improved in
the calculation.
Among FC events, single-ring events are identified as elike or μ -like based on a Cherenkov ring pattern. All the PC
events were assigned to be multi-GeV μ -like. Using the number of e-like and μ -like events, the ratio of μ e was obtained
and it is significantly smaller than the expectation as shown in
the table.
The zenith angle distributions for the sub- and multi-GeV
samples are shown in Fig. 1. The μ -like data from SK exhibited a strong up-down asymmetry in zenith angle (Θ) while
no significant asymmetry was observed in the e-like data. The
data were compared with the Monte Carlo expectation without
neutrino oscillations and the best-fit expectation for νμ  ντ
oscillations. The oscillated Monte Carlo reproduced well thezenith angle distributions of the data. Some fraction of themulti-ring events is also subdivided into e-like and μ -likeevents using the event pattern of the most energetic Cherenkov
ring in each event. Fig. 2 shows the zenith angle distributionof multi-ring events and they also agree well with the expectations from neutrino oscillations.
Energetic atmospheric νμ ’s passing through the Earth interact with rock surrounding the detector and produce muons
via charged current interactions. These neutrino events areobserved as upward going muons. Upward going muons
are classified into two types. One is “upward through-going
muons” which have passed through the detector, and the otheris “upward stopping muons” which come into and stop insidethe detector. The mean neutrino energies of upward throughgoing muons and upward stopping muons are 100 GeV
and 10 GeV, respectively. SK has observed 1856 upward
through-going muons and 458 upward stopping muons during1646 days’ live time. Fig. 2 shows the zenith-angle distributions of those upward muons. They agree with the expectations assuming neutrino oscillations.
We carried out a neutrino oscillation analysis using theentire SK-I atmospheric neutrino data. Fig. 3 shows the allowed neutrino oscillation parameter regions for νμ  ντ oscillations. The best fit oscillation parameters are sin2 2θ  1 0
and Δm2  2 1  10 3eV2 . The allowed oscillation parameter range is obtained to be sin2 2θ  0 92 and Δm2  1 5 
3 4  10 3eV2 at 90% C.L. [1].
The atmospheric neutrino data is well described by theneutrino oscillations shown above. In this case, the survival
probability of νμ is given by a sinusoidal function of LE,
where L is the travel distance, E is the neutrino energy. However, the sinusoidal LE dependence of the survival probability of νμ has not yet been directly observed. We used a
selected sample of these atmospheric neutrino events, thosewith good resolution in LE, to search for an oscillation max-4

Data/Prediction (null oscillation)

-2

2

2

Δm (eV )10

99% C.L.
90% C.L.
68% C.L.

1.8
1.6
1.4
1.2
1
0.8
0.6
0.4
0.2
0
1

-3

10

0.7

0.75

0.8

0.85

0.9

0.95

10

Fig. 3. Allowed region of νμ  ντ neutrino oscillation parameters obtained by SK by using contained atmospheric neutrino events andthe upward-going muon events.

2

10

3

10

4

L/E (km/GeV)

1

sin22θ

10

Fig. 4. Ratio of the data to the MC events without neutrino oscillation(points) as a function of the reconstructed L E together with thebest-fit expectation for 2-flavor νμ  ντ oscillation (solid line). The
error bars are statistical only.

-2

2
2

imum in the LE distribution.
The neutrino energy, E, was estimated from the total energy of charged particles observed in the inner detector. The
flight length of neutrinos, L, which ranges from approximately
15 km to 13,000 km depending on the zenith angle, was estimated from the reconstructed neutrino direction. The neutrino direction was taken to be along the total momentum
vector from all observed particles. Since the correlation between neutrino directions and the directions of observed particles are taken into account in Monte Carlo simulations, weapplied same analysis both for real events and Monte Carloevents. We applied a cut to reject low energy or horizontalgoing events since they have either large scattering angles orlarge dLdΘzenith .
Fig. 4 shows the observed LE distribution after taking ratio to the Monte Carlo events without neutrino oscillations. In the figure, a dip, which should corresponds tothe first oscillation maximum, is observed around LE 
500 kmGeV. The distribution was fit assuming νμ  ντ oscillations. The best-fit expectation shown in the figure corresponds to sin2 2θ  Δm2   1 00 2 4  10 3 eV2 . Fig. 5
shows the contour plot of the allowed oscillation parameterregions. The result is consistent with that of the oscillation
analysis using zenith angle distributions. The observed LE
distribution gives the first direct evidence that the neutrino survival probability obeys the sinusoidal functions as predictedby neutrino flavor oscillations [2].
Another interest is an observation of ντ in the atmospheric
neutrinos since there has not been apparent evidence for theappearance of ντ charged current interactions due to νμ  ντ
oscillations. We have performed a search for the ντ appearance by using the SK-I atmospheric neutrino data. The analysis is based on two statistical methods: a likelihood analysis
and a neural network. Since τ ’s produced in the Super-K detector would immediately decay into many hadrons, the event
pattern would be similar to that of high-energy multi-ring νe

Δm (eV )10

99% C.L.
90% C.L.
68% C.L.
-3

10
0.7

0.75

0.8

0.85

0.9

0.95

1

2

sin 2θ
Fig. 5. 68, 90 and 99% C.L. allowed oscillation parameter regions for
2-flavor νμ  ντ oscillations obtained in the L E analysis.

events. Using the statistical difference in τ -like and background events, we have derived the variables to select an enriched sample of ντ charged current events in the atmospheric
neutrino data. The differences appear in the energy spectrum,
the number of charged pions in the final state, the fraction of
lepton energy with respect to neutrino energy and so on. We
optimized a combination of the variables and defined a likelihood function so that the signal to noise ratio becomes the
maximum. We also use the neural network for the analysis.
Before these analyses, we applied a pre-selection: (1) fiducial
events, (2) multi-GeV, and (3) the most energetic ring is e-like.
After the pre-selection and cuts based on the likelihoodand neural network, we fit the zenith angle distribution of aντ enriched sample to a combination of the expected ντ andthe atmospheric neutrinos (νμ and νe ) including oscillations
with the oscillation parameter, sin2 2θ  1 0 and Δm2  2 1 Number of Events

5
-2

10

100

2

Δm (eV )

80

2

60

40
-3

10

20

0
-1

SK 90% C.L.
SK 99% C.L.
CHOOZ 90% CL exclude

0-0.8 -0.6 -0.4 -0.20

0.2 0.4 0.6 0.8

1

cosΘ

Fig. 6. The zenith angle distribution for tau candidate events in thelikelihood analysis. The dashed histogram (background neutrinos)
and the solid line (excess by ντ ) show the best fit for the data.

10 3 eV2 . Fig. 6 shows the fitted zenith angle distribution.
Using τ selection efficiencies estimated by Monte Carlo study,
we concluded the number of tau events with SK-I exposure is0
155 5012
72 for the likelihood analysis and 155 50 53 forthe neural network while 79 30 and 79 29 are expected foreach analysis. Thus the fitted results are found to be consistent
with pure νμ  ντ oscillations and ντ appearance.
The two flavor neutrino oscillations successfully described
the SK-I atmospheric neutrino data. However, any contributions by electron neutrinos have not been observed yet. We extended our neutrino oscillation analysis in order to treat threeneutrino flavors. For the analysis, Δm223  Δm213  Δm2 
Δm212 was assumed. If the parameter θ13 in the mixing matrix of lepton sector (MNS matrix) is finite, neutrino oscillations among νμ  νe may be observed. Moreover, the mixing parameter is affected by potentials caused by matter andoscillations are expected to have a resonance around 5 GeV.
Therefore, we can expect an increase at upward-going MultiGeV e-like samples. Fig. 7 shows the result of the three-flavor
neutrino oscillation analysis. Though there was no significant
excess of electrons, we set an upper limit on θ13 . More statistics is needed to have better sensitivity.
Atmospheric neutrino events are being accumulated during SK-II. A preliminary analysis has already shown the distortion of the zenith angle distribution and consistency with
SK-I results.

Solar neutrinos
The observed flux by the solar neutrino experiments(Homestake, Kamiokande, Super-Kamiokande, Gallium experiments, and SNO charged current) were lower thanthe expectations from the Standard Solar Model (SSM), which
has been known as “solar neutrino problem”. The precise
measurements of solar neutrinos by SK and SNO establishedthat the solar neutrino problem is due to neutrino oscillations.
SK detects 8 B solar neutrinos through neutrino-electron

0.1

0.2

0.3

0.4

0.5

2

sin θ13

Fig. 7. The allowed region for (Δm2 , sin2 θ13 ). We assumed positiveΔm2 . The solid and dotted contours correspond to allowed regions
obtained by this analysis. The hatched region corresponds to 90%
C.L. excluded area by CHOOZ experiment.

scattering, ν  e  ν  e. The electron energy, direction andthe time of the events are measured. The ν  e  ν  e methodof the solar neutrino detection is sensitive not only to νe but
also to νμ and ντ . The cross section of νμ  e and ντ  escatterings are about 17 of the νe  e cross section.
In order to perform precise solar neutrino measurement,
the energy scale, the energy resolution, the angular resolution
and the vertex position resolution were precisely calibratedusing LINAC system and 16 N events generated by a DT neutron generator. The absolute energy scale was calibrated withan accuracy of 0.6 %.
SK-I has observed about 22,400 solar neutrino events
above 5 MeV during 1496 days of live time from May 1996
to July 2001. This number corresponds to the 8 B solar ν fluxof 2 35 0 02 (stat.) 0 08 (sys.) 106 cm 2 s 1 , which is40.6 0.4 (stat.) 11 43 (sys.)% of the BP04 SSM prediction.
Sudbury Neutrino Observatory(SNO) observes solar neutrinos by using the charged current(CC), neutral current (NC) interactions, and electron scattering(ES). The CC flux, to which6
2 s 1 from theonly νe contributes, was 1.68 00 10
11 10 cm
391-day Salt Phase SNO data set. Figure 8 shows flux of νeand νμ  ντ obtained by SK flux, SNO CC flux and SNO NC
flux. The measured total neutrino flux was consistent with
the SSM prediction as shown in the figure. It also indicateda non-electron type neutrino flux, though all the solar neutrinos should be generated in electron type. This comparisonof the flux results between SK and SNO was initially done in2001 when the first SNO data was published, then this comparison became the first strong experimental evidence of thesolar neutrino oscillation among active neutrinos.
The measured energy spectrum and day/night differencein SK-I were used to constrain neutrino oscillation parameters. We have adopted an binning free method for day/nightvariation. Since the direction of the sun is precisely known
by the time of each event, the expected oscillation probabil-

Δχ

2

8
SNO CC

7 SK-I

Δm in eV

2

5
3σ
2σ
1σ

4

10

SSM ±1σ

1

10

0

1

2

6

10

νe flux (10 /cm /sec)

10

3

4

5

6

2

Fig. 8. νe flux and νμ  ντ flux contour obtained by SK and SNO data.

Zenith Seasonal

4

Spectrum2νe→νμ/τ (95%C.L.)

2

0

±1σ

-20

-4
-5
-6
-7
-8
-9
-10
-11

10

-4

10

-3

10

-2

10

-1

1

8 3σ
6
4 2σ
2
2 1σ

Δχ

Data/SSM

10 2
tan (Θ)

2

4

6

82
Δχ

Solar/KamLAND
Zenith Seasonal
Spectrum
ν →ν
e

μ/τ

95%C.L.

-4

x10
2

-30

8

10

12

14

16

18
20
E in MeV

Fig. 9. SK-I energy spectrum compared with the expected Spectrum from LMA(top) and D/N Asymmetry (bottom). The predictions (solid lines) are for tan2 θ  055 and Δm2  63  10 5 eV2
with φ8 B  096SSM and φhep  36SSM. The gray bands are the1σ ranges corresponding to the fitted value over the entire range
5-20 MeV: A  18  16%.

ity through the earth can be calculated on an event by event
basis. Using a maximum likelihood method to accommodate
this day/night variation and combining the energy spectrum
shape information, we have done the new oscillation analysis.
The energy range of the solar neutrino spectrum is divided into
21 energy bins and the day/night asymmetry value for eachenergy bin is shown in Fig.9. Assuming the total 8 B flux predicted by SSM, the best fit oscillation parameters in the LMA
region is obtained at tan2 θ  0 55 and Δm2  6 3  10 5eV2 .
Using the expected energy and time dependence of the event
rate, a fitted day/night asymmetry value is obtained to be1 8 1 6(stat)11 32(syst)% using the binning free maximum
likelihood method. It is statistically improved from the simple
day/night difference described above. The obtained day/night
asymmetry value is consistent with the expected asymmetryof 2.1% for the best fit oscillation parameters.
Figure 10 shows results of oscillation analysis. The dark

1

KamLAND

2

LMA best-fit D/N Asymmetry
6

Δm in eV

Asymmetry in %LMA best-fit Spectrum

-10

-50

SK 1496 Days

6

Fig. 10. Dark gray region shows the excluded region by SK-I spectrum and time variation analysis. The light gray region shows anallowed region assuming SSM flux prediction. Overlaid are the areas excluded just by the day/night and seasonal variation (hatchedregions inside thick black lines). The graphs at the top (and right)
show the χ 2 difference as a function of tan2 θ (Δm2 ) alone where
the Δm2 (tan2 θ ) is chosen to minimize χ 2 .

10

-40

10

10

2 SNO NC

0.75
0.7
0.65
0.6
0.55
0.5
0.45

10

10

3

0

8

2

6

6

2

νμ+ντ flux (10 /cm /sec)

6

solar

0

1σ 2σ

3σ

0.1 0.2 0.3 0.4 0.5 0.6 0.7 0.8 0.9
2 4 6 82
2
sin (Θ)Δχ
Fig. 11. Allowed areas at 95% C.L of all the solar data and KamLAND.
The combined area and the best-fit point are also shown. The
graphs at the top (and right) show the χ 2 difference as a functionof sin2 θ (Δm2 ) only.

gray region in the figure shows an excluded region by SK
spectrum and time variation without using absolute flux. The
light gray region shows an allowed region assuming SSM prediction. Taking into account the results of other solar neutrino
experiments, a global oscillation analysis is performed. The
combined fit to SK data and all the solar neutrino experimentsshows that only the LMA region remains as the allowed region.
Figure 11 shows the allowed regions from all the solar
neutrino experiments and KamLAND. The obtained allowed
region by the global solar neutrino analysis overlaps quite wellwith the reactor ν̄e oscillation results from KamLAND.
An extensive study to obtain the solar neutrino flux below
5.0 MeV in SK-I was done in this year. In this energy region,
the main reason why these events weren’t used so far, was aData/SSM

Event/day/kton

7
0.22

0.21

0.2

4.5−5.0MeV

1
SK-II 478day 22.5kt(Preliminary)

0.8
0.6

0.19

0.18

0.4

0.17

0.2

0.16

Line: SK-I 1496day average

0

0.15

SK−I 466days
20.5kton (Preliminary)

0.14

0.13

0.12

−1

−0.8

−0.6

−0.4

−0.2

0

0.2

0.4

0.6

0.8

1

cosθsunFig. 12. Angular distribution of the special data sample below 5.0 MeV
in SK-I. The dashed line is expected background shape in thisanalysis. The solid line is the best-fit.

high background level. Therefore, we have selected very stable and high trigger efficiency run periods, then applied new
analysis tools to reject the remaining very low-energy background events. The live time and energy threshold of thisanalysis were 466 days and 4.5 MeV, respectively. Figure 12
shows the angular distribution of this special data sample between 4.5 MeV and 5.0 MeV. We obtained a clear signal in thisenergy region. The observed number of events were 629+128126(stat.) events. This correspond to the 8 B solar neutrino
flux of 3 13 0 63(stat.).
In this year, we also obtained the preliminary resultsfrom SK-II. The run period for this analysis was from Dec.
24, 2002 up to Oct. 1, 2004. The total live time in thisperiod was 478 days. The analysis energy threshold was
8.0 MeV. The obtained solar neutrino flux was 2 45  0 08 
0 07 (stat.)106 cm 2 s 1 . This was consistent with the
SK-I result. The systematic errors are under studying. The
measured day-time and night-time fluxes are 2 44  0 11 
0 11 (stat.) and 2 35  0 11  0 10 (stat.), respectively. These
are also consistent with SK-I. The preliminary energy spectrum was also obtained. Figure 13 shows the energy spectrum
from the SK-II 478 days data. There was no significant distortion.

Search for nucleon decay
Proton decays and bound neutron decays (nucleon decays
in general) is a most dramatic prediction of Grand Unified
Theories in which three fundamental forces of elementary particles are unified into a single force. Super-Kamiokande (SK)is the world largest detector to search for nucleon decays andit has accumulated data of 91.6 ktyrs (SK-I) and 25.9 ktyrs(SK-II) resulting in 117.5 ktyrs data in total. Various nucleon
decay modes have been looked for in the SK but we found no
significant signal excess so far.
A proton decay into one positoron and one neutral pion(p  e π 0 ) is one of most popular decay mode. This decay

5

10

15
20
Energy(MeV)

Fig. 13. Energy spectrum of the solar neutrino signal from SK-II
478 days data. The horizontal axis is the total energy of the recoil
electrons. The vertical axis is the ratio of the event rates between
the observation and the BP04 SSM prediction. The horizontal line
shows the SK-I flux value.

mode is mediated by super-heavy gauge bosons and discoveryof the signal would give us the information of the mass of thegauge mesons. To discriminate the signal from atmospheric
neutrino background, we reconstruct the number of particles(Charenkov rings) and reconstruct the total visible energy corresponding to parent proton mass and total momentum corresponding to proton’s Fermi momentum. Even the photo coverage area is about half (19%) in SK-II, we achieved high detection efficiency of signals as 41% and low background levels as
0.1 events in 25.9 ktyrs of SK-II. Because there are no candidate events in SK-I + SK-II data, we obtained lower limiton the partial lifetime of the proton, τ B pe π 0  6 9  1033
years at a 90% confidence level.
Moreover, we looked for SUSY favored decay modes
which include K mesons in final state; p  ν̄ K  , n  ν̄ K 0 ,p  μ  K 0 , and p  e K 0 . In p  ν̄ K  search, we looked
for 236 MeV/c monochromatic muons from the decay of K  .
Figure 14 shows the comparison between data and fitting results of muon momentum distribution for single-ring μ -likeevents. We observed no excess of signal. In any other modes,
there are no significant signal excess. Therefore we conclude
that there is no evidence of nucleon decays and we calculatedpartial lifetime limits taking into account systematic uncertainties. Obtained limits are 2 3  1033, 1 3  1032, 1 3  1033,
1 0  1033 , years at 90% confidence level for p  ν̄ K  ,n  ν̄ K 0 , p  μ  K 0 , and p  e K 0 modes, respectively. We
have finalized this results using full SK-I data and submitted
to a journal.

Supernova neutrinos
Kamiokande and IMB observed neutrino burst from supernova 1987a. Those observations confirmed that the energy release by neutrinos is about several1053erg. SuperKamiokande is able to detect several thousand neutrino events
if it happens near the center of our galaxy.
During 1703.9 days of SK-I data taking period, there was
no evidence of such supernova explosions, and the 90% C.L.
upper limit on the rate of supernova explosions within 100 kpc

8

data
atmνMC (best fit)+p→νK MC

number of events

60

40

20

0
200

225

250

275

300

μ momentum (MeV/c)
Fig. 14. The comparison between data and fitting results of muon momentum distribution for single-ring μ -like events. The filled circles
show data with statistical errors. The solid line shows p  ν̄ K 
MC. The dashed line shows the best fitted atmospheric neutrino
MC with free normalization.

is obtained to be 0.49 explosions per year.
After the reconstruction of SK-II, the realtime supernova
monitor system was restarted again as in SK-I. If a supernova
is detected, this monitor reports it to shift persons in 5–10 minutes later. This monitor also sends a signal to SNEWS (the SuperNova Early Warning System) which takes coincidence of
signals from SK, SNO, and LVD. Until now, no coincidence
events were observed.
Due to the half photocoverage of SK-II, the energy threshold for the real time monitor was raised from 6.5 MeV to
8.5 MeV. However, since dominant signals of supernova neutrinos have around 15 MeV energy, the decrease of the detection efficiency is small, and SK-II still has 100% efficiency forgalactic supernovae.

Bibliography

In collaboration with the members of:
ICRR, University of Tokyo, Japan; Kyoto University, Japan;
KEK, High Energy Accelerator Research Organization,
Japan; Kobe University, Japan; Osaka University, Japan; Hiroshima University, Japan; Okayama University, Japan; Niigata University, Japan; Tokyo University of Science, Japan;
Tohoku University, Japan; Miyagi Edu. University, Japan;
Chonnam University, Korea; Seoul National University, Korea; Dongshin University, Korea; Korea University, Korea; University of California at Irvine, USA; State University of New York at Stony Brook, USA; Boston University,
USA; University of Washington at Seattle, USA; Louisiana
State University, USA; Massachusetts Institute of Technology,
USA; University of Hawaii at Manoa, USA; Duke University, USA; TRIUMF, Canada; University of British Colombia,
Canada; University of Barcelona, Spain; University of Valencia, Spain; INRM, Rossia; CEA Saclay, France; University ofRome, Italy; University of Geneva, Swiss; University of Warsaw, Poland; A. Soltan Institute for Nuclear Studies, Poland
The discovery of the neutrino oscillation by SuperKamiokande in 1998 has lead to the conclusion that neutrino
has a finite mass. The discovery implies the existence of new
physics beyond the standard model of the elementary particles
at a huge energy scale. Since the observed neutrino mass difference of Δm2 310 3eV2 suggests the oscillation lengthis about a few hundred km for one GeV neutrinos, this phenomena can be experimentally tested by using the artificiallycreated neutrino beam with the detector placed at a few hundred km away from the neutrino production point.
The k2k experiment was planned to detect such oscillation effect and to determine the oscillation parameters precisely. The neutrino beam created at KEK has a mean energyof 1.3 GeV and was sent to the Super-Kamiokande detector,
250 km west of KEK, every 2.2 seconds with the duration of
1.1 μ second. In the site of KEK, at the distance of 300 m fromthe target, we have prepared near detectors which mainly consist of 1 kt water Cherenkov detector, fine grained scintillation
fiber detector (FGD), a lead glass calorimeter, and the muon
ranger. The supplemental detectors are also placed. The ar-[1] Super-Kamiokande Collaboration, “A Measurement ofAtmospheric Neutrino Oscillation Parameters by SuperKamiokande I”, Submitted to Phys. Rev. D, hepex/0501064.
[2] Super-Kamiokande Collaboration, “Evidence for an oscillatory signature in atmospheric neutrino oscillation”,
Phys. Rev. Lett. 93 (2004) 101801.
[3] K. Kobayashi et al. (Super-Kamioknade Collaboration), “Search for nucleon decay via modes favored by supersymmetric grand unification models in
Super-Kamiokande-I”, Submitted to Phy. Rev. D, hepex/0502026.

K2K Experiment[Spokesperson : K. Nishikawa]
Kyoto University, Kyoto 606-8502

Fig. 1. The arrangement of the near detectors.

2000
0

0

0.5

120

1

1.5
2
pμ (GeV/c)

(c)

80
40
0

0

1

2
3
4
pμ (GeV/c)

(b)

10000
5000
00

20
60

120
180
θμ (deg.)

125nsec

15

(d)

200

10

100
0

SK event timing (1bin=125/2 (nsec))

0

30

60
90
θμ (deg.)

Fig. 2. Muon momentum (pμ ) and direction (θμ ) distributions: (a) thepμ distribution of 1KT fully contained 1-ring μ -like sample, (b) 1KT
θμ for the same sample, (c) SciFi pμ for 2-track QE sample, and (d)SciBar θμ for 2-track nonQE sample. Open circles represent data,
while histograms are MC predictions using the best fit systematic
uncertainties.

rangement of the front detectors at the beginning of the experiment is shown in Fig. 1. We have analyzed data taken from
June 1999 to February 2004, which corresponds to 8.9  1019
protons on target (POT). The near detector data from this period include 2 3  1019 POT without the lead glass (K2K-IIa),and then 1 9  1019 POT (K2K-IIb) with a fully-active scintillator detector (SciBar) in its place.
There are several beam monitors, proton profile monitors,
the muon monitor and the pion monitor, placed along the neutrino beam line. The decay length is about 200 m, therefore
the beam is not scaled exactly by 1/r2 law. The flux ratio of thenear to the far detector was obtained by the beam Monte Carlocalculation by taking into account of the spread and emittanceof the beam at the target, the production of the pions, the focusing effect by the HORN and the decay of pions to neutrinos. The beam calculation was validated by the pion monitor
which measured the produced pion directions and momentum.
The beam flux and the spectrum were measured by thefront detectors. ICRR group has a responsibility of providing
the Super-Kamiokande data and the construction of the 1 kt
front detector and the analysis of the data from those detector
components. The coverage of PMT in the 1 kt front detector is same as the Super-Kamiokande detector with the 40%of the total inner detector surface. The 1kt detector provides
information of the absolute neutrino flux and the spectrum of
the neutrino beam. Figure 2 shows (a) the muon momentum
spectrum and (b) the muon direction by using the single ringfully contained events at 1 kt detector. The observed spectrum
is well reproduced by the Monte Carlo simulation in whichbeam energy spectrum and neutrino interaction in water are
taken into account. Combining the 1 kt data and FGD data
together with constraint of the pion monitor measurement, theexpected neutrino energy spectrum was obtained for the far
detector. All the systematic errors due to detector bias or uncertainty of neutrino interactions were carefully taken into account.
The synchronization of the timing between KEK 12 GeV
Proton Synchrotron and Super-Kamiokande detector was performed using a GPS system. Figure 3 shows the timing
distribution of the fully contained events observed at Super-

5

0

0

500
1000
event timing (nsec)

Fig. 3. The timing distribution of fully contained event observed at
Super-Kamiokande with respect to the beam injection timing. Periodical structure coming from proton beam bunches is seen in the
Super-Kamiokande, 250 km away from the accelerator.

Events / 0.2 (GeV)

4000

Entries / 10 (deg.)

(a)

6000

Entries / 5 (deg.)

8000

Entries / 0.25 (GeV/c)

Entries / 0.1 (GeV/c)

9

16
12
8
4
0

0

1

2

3

4
5
Eνrec (GeV)

Fig. 4. The reconstructed Eν distribution from single ring fully contained events at SK. Points with error bars are data. The solidline is the best fit spectrum. The dashed line shows the expected
spectrum without oscillation. Both histograms are normalized by
the number of events observed(57).

Kamiokande with respect to the proton beam injection timing. The time of flight of neutrinos from KEK to SuperKamiokande (830 μ sec) is corrected in the figure.
Neutrino events were clearly observed near ΔT=0 and the
spread in the timing was consistent with the duration of beam
injection at KEK 12 GeV PS. In total, 107 events were identified as K2K beam induced events at Super-Kamiokande.
Using the large number of observed neutrino events at front
detectors, expected number of neutrino events at SuperKamiokande was calculated to be 1511210 (syst) assuming nooscillation.
57 μ -like single ring events were observed and the neutrino energy spectrum obtained from the energy and scattering angle of the observed muons is shown in Fig. 4. From thedeficit of observed number of events together with the shape
distortion of neutrino energy spectrum, we have concluded
that probability of null oscillation is 0.0050% (4.0σ ). The
allowed region of Δm2 and sin2 2θ is shown in Fig. 5. The obtained result for neutrino oscillation parameters are consistentΔm2 (eV2)

10

10-1

Overview

68%
90%
99%

10-2

10-3
0

0.2

0.4

0.6

0.8 1
sin22θ

Fig. 5. Allowed regions of νμ -ντ oscillation parameters. Dashed, solidand dot-dashed lines are 68.4%, 90% and 99% C.L. contours, respectively.

with that from atmospheric neutrino oscillations.
K2K is expected to finalize physics results using full K2K
data in near future including νμ disappearance, νe appearance,and detailed neutrino interaction studies.

Bibliography[1] K2K collaboration, Detection of accelerator produced
neutrinos at a distance of 250 km, Phys. Lett. B, 511,
178–184, 2001.
[2] K2K collaboration, Indication of neutrino oscillation in
a 250 km long baseline experiment, Phys. Rev. Lett., 90,
041801, 2003.
[3] K2K collaboration, Search for Electron Neutrino Appearance in a 250 km Long-baseline Experiment, Phys.
Rev. Lett., 93, 051801, 2004.
[4] K2K collaboration, Measurement of single pi0 production in neutral current neutrino interactions with waterby a 1.3 GeV wide band muon neutrino beam, submitted to Phys. Lett. B, hep-ex/0408134.
[5] K2K collaboration, Evidence for muon neutrino oscillation in an accelerator-based experiment submitted toPhys. Rev. Lett., hep-ex/0411038.

XMASS Experiment[Spokesperson : Yoichiro Suzuki]
Kamioka Observatory, ICRR, Univ. of Tokyo, Gifu, 5061205
In collaboration with the members of:
Kamioka Observatory, ICRR, Univ. of Tokyo, Japan; RCCN,ICRR, Univ. of Tokyo, Japan; Waseda Univ., Japan; Yokohama National Univ., Japan; Miyagi Univ. of Education,
Japan; Nagoya Univ., Japan; Saga Univ., Japan; Gifu
Univ.,Japan; Seoul National Univ., Korea; INR-Kiev, Russia;
Univ. of California, Irvine, USA; Sejon Univ., Korea.

XMASS is a multi-purpose experiment using liquid xenon
aiming at the detection of cold dark matter, neutrino absolute
mass using neutrinoless double beta decay, and low energy
solar neutrinos. An R&D study for the liquid xenon detector
is being performed at Kamioka observatory.
Astronomical observations suggest that there is “dark matter” (non-luminous particles with mass) in the universe. Oneof the most likely candidates for dark matter is a weakly interacting massive particle (WIMP) like the lightest supersymmetry particle. A recoil of xenon by dark matter produces
scintillation light in liquid xenon.
The Super-Kamiokande measurements show that neutrinos have masses. However, we do not know yet the absolute
mass of neutrinos and the whether neutrino masses are Majorana type or Dirac type. Xenon nuclei with mass number 136
is one of the double beta nuclei which is best suited for thisresearch.
The solar neutrino spectrum is measured only above
5 MeV by SK and SNO so far. The spectrum of low energy solar neutrinos (pp and 7 Be neutrinos and etc.) are not measured
yet. If a 10-ton class liquid xenon detector is constructed, it
will be able to detect those neutrinos by ν +e scattering with arate of about 10 events/day.
Liquid xenon has the following advantages:
A large light yield of 42,000 photons/MeV, which is as
good as NaI(Tl) scintillator, enables detection of smallenergy signals like dark matter recoil.
Because of the higher atomic number of xenon(Z=54)
and higher density of liquid xenon (3g/cm3), external
gamma-ray background can be reduce in a short distance from the detector wall by self-shielding.
175 nm scintillation light of liquid xenon can be readout by typical PMTs of bi-alkaline photocathode with aquartz window.
Purification is easier than other materials (e.g. distillation is possible).
Isotope separation is possible. It is possible to enrich
136
Xe for double beta decay and deplete 136 Xe for solar
neutrino measurements.
A 3 kg fiducial volume liquid xenon detector has been developed for the R&D study and test data was taken in 2003
and 2004.
We plan to make an 800 kg detector (Fig.1) aiming tosearch for dark matter down to 10 45  10 44 cm2 in crosssection which is more than two orders of magnitude better
than the current best limits in the world.

R&D study by a 3 kg fiducial volume detector
Figure 2 shows the 30 liter (3 kg fiducial volume) liquid
xenon detector developed for the R&D study. It is a 30 cm cubic detector viewed by 54 2-inch PMTs. The PMTs are Hamamatsu R8778 which were developed by collaborating with
Hamamatsu and can be used for low background purposes andalso can be used at liquid xenon temperature (170K). The
detector is placed in a low background setup which consists of
5 cm-thick pure copper (OFHC), EVOH sheets, 15 cm-thickevent rate (dru)

11

1

-1

10

-2

10

0

Fig. 1. Schematic view of the proposed 800 kg liquid xenon detector
aiming at dark matter search. The detector has a sensitivity of
10 45  10 44 cm2 in cross section which is more than two ordersof magnitude better than the current best limits in the world.

500

1000

1500

2000

2500
3000
energy (keV)

Fig. 3. Backround spectrum measured by the 3 kg fiducial volume
detector. The dashed, dotted and solid histograms show the rateof the events which were observed in all volume, 20 cm and 10 cm
cube centered at the center of the detector, respectively.

Bibliography[1] Y. Suzuki et al., “Low energy solar neutrino detection byusing liquid xenon,” Aug. 2000, hep-ph/0008296.
[2] S. Moriyama, “XMASS EXPERIMENT I,” proceedingsof the International Workshop on Technique and Application of Xenon Detectors, Kashiwa, Japan, 3–4 Dec.
2001.

Fig. 2. 3 kg fiducial volume liquid xenon detector constructed for R&D
study

lead, 10 cm-thick boric acid, and 15 cm-thick polyethylene for
reducing gamma rays, neutron and radon backgrounds.
Test data was taken in December 2003 and August 2004.
Figure 3 shows the background distributions for the latest measurement. The large reduction of backgrounds in asmaller detector volume is observed at several hundreds keV,which demonstrates self-shielding of the detector. The reduction of backgrounds is as expected by Monte Carlo simulations.
As the internal source of backgrounds, 85 Kr in xenon is
one of potential sources. In order to purify xenon, a test distillation system was constructed which is able to process 0.6 kg
Xe/hour and achieve better than 1/1000 reduction in krypton
concentration. In March 2004, we processed original xenon
which contains 3 ppb mol/mol krypton with this distillation
system and measured krypton concentration by an API-MS
detector after the process. As a result, it was found that theprocessed xenon contains only 3 3 1 1 ppt mol/mol krypton
and that the 1/1000 reduction had been achieved.

[3] Y. Takeuchi, “Recent status of the XMASS project,”presentation in the International Conference on HighEnergy Physics (ICHEP), Beijing, China, 16–22 Aug.
2004.
[4] S. Moriyama, “XMASS experiment,” proceedings of theInternational Workshop on the Identification of Dark
Matter, Edinburgh, Scotland, 6–10 Sep. 2004.

12

HIGH ENERGY COSMIC RAY DIVISION
Overview
Three major research activities of the High Energy Cosmic
Ray division are the study of very high energy gamma-raysby the CANGAROO group, extremely high energy cosmic
rays by the Telescope Array (TA) group, and very high energy
cosmic rays and gamma-rays by the Tibet ASγ Collaboration.
Other activities, such as experiments utilizing the Akeno observatory, the Norikura observatory, the Mt. Chacaltaya observatory (jointly operated with Bolivia), and the emulsionpouring facilities are closely related to inter-university joint
research programs. Also the activity for development of highresolution telescope is intensively going on.
The CANGAROO project (Collaboration of Australia andNippon for a GAmma-Ray Observatory in the Outback) isa set of large imaging Cherenkov telescopes to make a precise observation of high-energy air showers originated by TeV
gamma-rays. It started as a single telescope with a relatively
small mirror (3.8 m in diameter) in 1992. In 1999 a new telescope with a 7-m reflector has been built, and now it has a
10-m reflector with a fine pixel camera. The main purposeof this project is to explore the violent, non-thermal universe
and to reveal the origin of cosmic-rays. An array of four
10-m telescopes has been completed in March 2004 so that
more sensitive observation of gamma-rays is realized with its
stereoscopic imaging capability of Cherenkov light. We havedetected several gamma-ray sources in the southern sky anddetailed study of these sources are now ongoing.
At the Akeno observatory, a series of air shower arraysof increasing geometrical sizes have been constructed and operated to observe extremely high energy cosmic rays (EHECRs). The Akeno Giant Air Shower Array (AGASA) was
operated from 1991 to January 2004 and covered the groundarea of 100 km2 as the world largest air shower array. In
13 years of operation, AGASA observed a handful of cosmic
rays exceeding the theoretical energy end point of the extragalactic cosmic rays (GZK cutoff) at 1020 eV. The Telescope
Array (TA), a large plastic scintillator array with air fluorescence telescopes under construction in Utah, USA, will succeed AGASA and measure the EHECRs with an order of magnitude larger aperture than that of AGASA to unveil the origin
of super-GZK cosmic rays discovered by AGASA.
An air shower experiment aiming to search for celestial
gamma-ray point sources started in 1990 with Chinese physicists at Yangbajing (Tibet, 4,300 m a.s.l.) and has been successful. This international collaboration is called the Tibet
ASγ Collaboration. An extension of the air shower array wascompleted in 1995 and an emulsion chamber has been combined with this air shower array since 1996 to study the primary cosmic rays around the knee energy region. After successive extensions carried out in 1999, 2002 and 2003, thetotal area of the air shower array amounts to 37,000 m2 . The
sun’s shadow in cosmic rays affected by the solar magnetic
was observed for the first time in 1992, utilizing its good an-

gular resolution at multi-TeV energy region. From this experiment with better statistics, we expect new information to
be obtained on the large-scale structure of the solar and interplanetary magnetic field and its time variation due to the
11-year-period solar activities.

CANGAROO-II & -III Project[Spokespersons : R. W. Clay, M. Mori, and T. Tanimori]In collaboration with the members of:
Institute for Cosmic Ray Research, University of Tokyo,
Kashiwa, Chiba 277-8582, Japan: Department of Physics,
Graduate School of Science, Kyoto University, Sakyo-ku,
Kyoto 606-8502, Japan: RSAA, Australian National Univ.,
ACT 2611, Australia: Dept. of Physics, Univ. of Adelaide,
SA 5005, Australia: Department of Physics, Yamagata University, Yamagata, Yamagata 990-8560, Japan: Institute ofSpace and Astronautical Science, Japan Aerospace Exploration Agency, Sagamihara, Kanagawa 229-8510, Japan: Faculty of Management Information, Yamanashi Gakuin University, Kofu, Yamanashi 400-8575, Japan: Department of
Physics, Tokai University, Hiratsuka, Kanagawa 259-1292,
Japan: Department of Physics, Konan University, Kobe,
Hyogo 658-8501, Japan: Ibaraki Prefectural University of
Health Sciences, Ami, Ibaraki 300-0394, Japan: Faculty ofEngineering, Shinshu University, Nagano, Nagano 480-8553,
Japan: Solar-Terrestrial Environment Laboratory, Nagoya
University, Nagoya, Aichi 464-8602, Japan: National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588,
Japan: School of Allied Health Sciences, Kitasato University,
Sagamihara, Kanagawa 228-8555, Japan: Faculty of Science,
Ibaraki University, Mito, Ibaraki 310-8512, JapanTo summarize CANGAROO-II and -III of the fiscal year
2004, three important things happened;
1. H.E.S.S.1 published the confirming results to the
CANGAROO-II.
2. H.E.S.S. published the conflicting results to the past
CANGAROO-I.
3. The stereoscopic analysis was partially established.
We report them as follows.
CANGAROO-II results
The result of the supernova remnant RX J1713.7-3946 [1]was confirmed by H.E.S.S. [2]. H.E.S.S. observed shell structure of this supernova remnant and the flux is consistent ( 0.7
Crab) with our one as shown in Figure 1. The spectral index
slightly differs, however, the deviation is only 2.2 σ level. We*1 H.E.S.S. is the new array of four 12m-diameter Cherenkov telescopes
built in Namibia by the group leaded by Max Planck Institute for Nuclear
Physics (Germany).

13
CANGAROO-II

Radio emission
ASCA

Fig. 1. Differential gamma-ray fluxes obtained by CANGAROO-II and
H.E.S.S. Copied from Fig. 3 of Ref. [2].

10

10

10

-6

10

-2

-1

E dN/dE (m s TeV)

10

10

-5
-6

-8
-9

10

-7

EGRET

-7

-5

10

-4

10

-3

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-2

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-1

1

10

2

10

10

-8
HESS June/July
HESS July/August

10

CANGAROO-II-9

Whipple10-1

1Fig. 3. The “significance map” obtained by the CANGAROO-II telescope is shown by the blue contours. The telescope tracking center is at the north-west rim of RX J0852.04622. The red contours
with levels at 20%, 45%, 65%, and 80% are an ASCA GIS image.
The green contours show the 4850 MHz radio emission.

10
Energy (TeV)

Fig. 2. Energy spectrum E 2 dN dE of γ -rays from the Galactic Centre.
Full circles: H.E.S.S. ‘July/August 2003’ data set. Full triangles:
H.E.S.S. ‘June/July 2003’ data set. The line indicates a power-law
fit to the ‘July/August’ spectrum. Open squares: CANGAROO-II
spectrum from Summer 2001 and 2002 (Tsuchiya et al. [3]). Open
triangle: Whipple flux from 1995 through 2003 (Kosack et al. [5]),converted to a differential flux at the peak detection energy assuming a Crab-like spectrum. The inset shows the EGRET flux from1991 to 1996 (Mayer-Hasselwander et al. [6]) (circles) comparedto fits to the CANGAROO-II (dashed line) and H.E.S.S. (solid line)
spectra. Due to the poor angular resolution of EGRET (1Æ ) the flux
shown may include other sources. Copied from Fig. 4 of Ref. [7].

will observe this in summer of 2005 in order to investigate thespectral dependence on the shell positions using stereoscopic
observation techniques.
The Galactic Center result [3, 4] was also confirmed by
H.E.S.S. [7] and VERITAS [5]. The fluxes obtained by thesethree groups are summarized in Fig. 2. The flux level at low
energy differed by one order at the lowest energy binning.
We investigated our analysis, however, at present we have notfound any faults in our code. The coming observation in thesummer of 2005 will be awaited. Also these experimental data
gave significant upper limits in the relic density of the cold
dark matter of TeV mass.
The second supernova remnant RX J0852.0-4622 was revealed by CANGAROO-II [8]. The emission profile is shownin Fig. 3.
Recently and preliminary, it was confirmed by H.E.S.S.
[9]. CANGAROO-III (stereoscopic system) had observed thissupernova remnant in 2004 Feb. Our preliminary analysis
confirmed these results. We are preparing publication on thisresult. The flux level is consistent with H.E.S.S.
CANGAROO-I enigma
H.E.S.S. published two serious conflictions withCANGAROO-I results [11, 10]. They are concerning withthe supernova remnant SN 1006 and pulsar PSR B1706-44.
The H.E.S.S.’s upper limits corresponded to 1/10 of the
CANGAROO-I fluxes. CANGAROO-III had observed these
two objects and obtained enough statistics. We are doing ourbest in order to publicize our most recent results on these two
objects in the summer conferences of 2005.
CANGAROO-III results
The CANGAROO-III stereoscopic observations started
from 2003 Dec. Presently the work was concentrated on developing the calibration and analysis softwares. Also nicething is the improvement of the camera system. By introducing the optimized-size PMT and Winston-cone light-guide,
the light correction efficiency improved by a factor of two.
We now can stably observe μ ’s Cherenkov rings, the radius ofwhich is about 1.3 degrees as shown in Fig. 4. We have developed a high efficiency and clean selection of these events.
After applying cuts, we can select them [12] without reducing
the acceptance. These events can be used for calibration ofreflectivity and spot size of mirrors. Previously, we estimatedthe absolute energy via using cosmic-ray rate and night sky
background. With these, our measurement became more accurate and they would help in solving the differences in theabsolute fluxes in between the major Cherenkov telescopes.
Using the tuned Monte-Carlo inputs, we can see the standard candle, i.e., the Crab. The angular spectrum of the Crabis shown in Fig. 5. The angular resolution is consistent withthat obtained by the μ -ring study. Also the flux is consistent
with the world average. The Hillas parameter distributions
are consistent with the Monte-Carlo predictions. The signalto-noise ratio is worse compared to that of H.E.S.S. which ismainly due to the spot size deterioration of the mirror system

number of events

14

25000
20000
15000
10000
5000
0

0 0.5 1 1.5 2 2.5 3 3.5 4curvature(=1/r) (1/degree)

Number of events

Fig. 4. Curvature (= inverse of radius) distributions for μ -rings. The
blank histogram is for all events and the hatched for the selected
ones.

Fig. 6. Upper limit (the red arrow) for the gamma-ray emission fromthe Galactic Disk, together with the other wave-length data.

150

-1

Integral Flux (cm sec )

100
-10

10

Durhram
Narrabri

-2

50

-11

J1 (steady)

10

0
0

0.1

0.2

0.3
2

0.4
2

θ (deg )Buckland ParkCANGAROO-I
J3 (burst)-1210
J2 (steady)

-13

10
-14

Fig. 5. θ 2 distribution for the Crab. The point with error bars are madefrom the signal region and the histogram from the background region. The red line is the cut position.

[13]. The alignment work is necessary in 2005.
The Galactic disk is the strongest source of gamma-raysin the GeV energy region and could be the source of TeV
gamma-rays. Two regions of the Galactic disk, around  b 
13Æ  0Æ  and 19 5Æ 0Æ , which are not visible by the northern telescopes, were scanned for the first time at TeV energies by the CANGAROO-III stereoscopic system. We could
not find any statistically significant signals and obtained upper
limits on diffuse gamma-ray flux above 600 GeV as shown in
Fig. 6. We discussed the spectral index measured by EGRET
using this result [14].
The giant radio galaxy Cen A, which is the nearest(z  0 0018) active galactic nuclei and the first claimed
TeV gamma-ray source in 1975 [15], was observed by
CANGAROO-III. Again we could not see any statistically significant signal and obtained upper limits as shown in Fig. 7.
The limit corresponded to be 1/10 of Crab flux. These are also
lower than the previous measurements by an order. The discussion of the AGN unified picture was carried out based on

10
-15

10
10

3

10

4

10

5

Energy (GeV)

Fig. 7. Upper limits obtained for Cen-A (the red arrows), together withthe previous measurements (the black). The line is the Crab flux.

this result [16].
Prospect
We would like to publish these first results after summarizing the above-mentioned calibrations and etc. Also in thesummer conference, we might open our results on SN 1006
and PSR B1706-44 based on the results of the stereo observations.

Bibliography[1] “The acceleration of cosmic-ray protons in the supernova remnant RX J1713.7-3946,” R. Enomoto et al.,Nature, 416, 823 (2002).

15[2] “High-energy particle acceleration in the shell of a supernova remnant,” F. A. Aharonian et al., Nature, 432,
75 (2004).
[3] “Detection of sub-TeV Gamma-rays from the Galactic Center Direction by CANGAROO-II” K. Tsuchiya,
R. Enomoto, L. T. Ksenofontov, M. Mori, and T. Naito,
et al. [CANGAROO-II Collaboration], Astrophys. J.
606, L115 (2004) [arXiv:astro-ph/0403592].
[4] “Very High Energy Gamma-ray Observations of the
Galactic Center with the CANGAROO-II Telescope,”
K. Tsuchiya, Ph. D. Thesis, Univ. of Tokyo, Jan. 2005.
[5] Kosack, K., Badran, H. M., Bond, H. I., et al. 2004, ApJ,
608, L97.
[6] Mayer-Hasselwander, H. A., Bertsch, D. L., Dingus, B.
L., et al., 1998, A&A, 335, 161.
[7] “Very high energy gamma rays from the direction of
Sagittarius A ”, A. Aharonian et al., A&A, 425, L13(2004).
[8] “Detection of Gamma Rays around 1 TeV from
RX J0852.0-4622 by CANGAROO-II,” H. Katagiri,
R. Enomoto., L. T. Ksenofontov, and M. Mori et al.
[CANGAROO Collaboration], Astrophys. J. 619, L163
(2005) [arXiv:astro-ph/0412623].
[9] “Another Shell-Type Supernova Remnant - RX J0852.04622,” H.E.S.S. Collaboration, http://www.mpihd.mpg.de/hfm/HESS/public/som/current2.htm[10] “Search for TeV emission from the region around PSR
B1706-44 with the H.E.S.S. experiment,” F. Aharonian
et al., A&A, 432, L9 (2005).
[11] “Upper limits to the SN1006 multi-TeV gamma-ray flux
from H.E.S.S. observations,” F. Aharonian et al., A&A,
in press.
[12] Master thesis in Japanese, Y. Adachi, Univ. of Tokyo.
2005.
[13] Master thesis in Japanese, R. Kiuchi, Univ. of Tokyo.
2005.
[14] “Very High Energy Gamma-ray Observations of the
Galactic Plane with the CANGAROO-III Telescopes,”
M. Ohishi, Ph. D. Thesis, Univ. of Tokyo, Mar. 2005.
[15] Grindlay, J. E., Helmken, H. F., Hanbury Brown, R.,
Davis, J., Allen, L. R. 1975, ApJ 201, 82.
[16] “Observation of TeV Gamma-rays from the Active Radio Galaxy Centaurus A with CANGAROO-III Imaging
Atmospheric Cherenkov Telescope,” S. Kabuki, Ph. D.
Thesis, Univ. of Tokyo, Jun. 2005.

Study of the Most Energetic Cosmic Rays
— AGASA Collaboration[Spokesperson : M. Teshima]
Max-Plank Institute for Physics, Germany
In collaboration with the members of:
Kinki University, Osaka; Yamanashi University, Kofu;
Saitama University, Saitama; Musashi Institute of Technology, Tokyo; Tokyo Institute of Technology, Tokyo; Nishina
Memorial Foundation, Tokyo; Hirosaki University, Hirosaki;
Osaka City University, Osaka; RIKEN, Saitama; Max-Plank
Institute for Physics, Germany; Ehime University, Matsuyama; Fukui University of Technology, Fukui; Communications Research Laboratory, Tokyo; National Institute of Radiological Sciences, Chiba; University of Chicago, U.S.A.;
Chiba University, Chiba.

Introduction
The AGASA is a giant air shower array covering an areaof 100 km2 . It has been in operation in Akeno village andits vicinity, 130 km west of Tokyo, in order to study EHECRs above 1019 eV. In this energy region, distinctive featuresin the energy spectrum and arrival direction distribution are
expected.
If the cosmic rays are of extragalactic origin, photopion
production between cosmic rays and primordial microwave
background photons becomes important at energies above 6 1019 eV with a mean free path of about 6 Mpc. Therefore
a cutoff in the energy spectrum is expected at around several
times 1019 eV, and is called Greisen-Zatsepin-Kuzmin (GZK)
cutoff [3]. Cosmic rays exceeding this cutoff energy must be
originated from nearby extra-galactic sources within 50 Mpc.
For cosmic rays of this origin, expected arrival directions are
relatively isotropic and should have some correlation with thematerial distribution around our galaxy; super galactic planeor clusters of galaxies.
On the other hand, if they are galactic origin, their arrival
directions should have a strong correlation with the galactic
structure, since the gyro-radius of protons of energy above
1019 eV are comparable with the galactic scale. A studyof correlations of EHECRs with the galactic structure and/orwith the large scale structure of galaxies is important.
As a mechanism of producing these EHECRs, the diffusive shock acceleration is most widely accepted. Shocks atradio lobes of Active Galactic Nuclei, termination shocks in
the local group of galaxies and others have been discussed to
be possible sites [4]. The acceleration above 1020 eV in dissipative wind models of cosmological gamma-ray burst is also
proposed [5]. Since various theoretical difficulties exist in accelerating cosmic rays up to the highest observed energies, it
is also discussed that these cosmic rays are possibly the decay
products of some massive particles produced at the collapse
and/or annihilation of cosmic topological defects which could
have been formed in a symmetry-breaking phase transition in
the early universe [6].

16Energy Spectrum
In order to calculate the primary energy of giant air showers observed by AGASA, the particle density at a distance of
600 m from the shower axis (S(600)) is used as an energy
estimator, which is known to be a good parameter [7]. The
conversion factor from S(600) [m 2 ] to the primary energy
E0 [eV] is derived by simulation [8] as
E0  2 0  10

17



26

10

J(E) E 3 [m −2sec−1 sr −1 eV 2 ]The details of the AGASA are described in Chiba et al. [1]and Ohoka et al. [2]. The array is operational since 1991, andthe exposure of 5.81016 m2 s sr above 1019 eV has been
accumulated by the end of the AGASA operation (January2004). Extensive air showers with zenith angles smaller than45Æ and with core locations inside the array area were used forthe analysis. In the energy range above 1019 eV, we only usedsuch events that passed through an eye scan to confirm theircore positions and energies.

2

25

10

6
10

3

3
24

10

Uniform sources

23

S 6001 0

10

19

(1)

The details of determination of arrival directions andS(600) are described in Takeda et al. [10].
The differential primary energy spectrum above 1018 5 eV
is shown in Fig. 1 [11]. The bars represent statistical errors only. The resolution of energy determination is about
25% at 1020 eV, which is estimated by analyzing the simulated showers. The dominant contribution of the resolution
comes from the fluctuation of the shower development and thelimited number of sampled particles by the scintillator. The
overall systematic error in determining the energy scale of the
primary cosmic ray is also analyzed by the simulation. The
dominant source of the error comes from the unknown composition of the primary and the unknown cross section of thehadronic interaction at extremely high energy. It is estimatedto be 18% [12].
The energy spectrum extends beyond 1020 eV and we observed 11 events above the GZK cutoff energy. The expectedenergy spectrum with the GZK cutoff is shown in Fig. 1 by thedashed curve. The expected number of events above 1020 eV
is 1.8 taking the energy resolution of AGASA into account.
The energy spectrum is therefore extending beyond 1020 eV
without the GZK cutoff.

Arrival Direction
Above GZK energy
The arrival direction of the highest energy cosmic rays
above 4  1019 eV is shown in Fig. 2 [13] in the equatorial
coordinates. Small circles and squares show events above
4  1019 eV and 1020 eV, respectively. They appear to be
distributed uniformly over the observed sky. We have carried out several tests for global anisotropy but could not findany large scale structures in this distribution. However, wefound one triplet and five doublets of events clustering within
2 5Æ angular separation as indicated by shadowed circles inFig. 2. The chance probability for creating these clusteringfrom the uniformly distributed cosmic rays was evaluated as
low as Pch  10 3 . No outstanding astronomical objects havebeen found in the directions of these clusters within 50 Mpc.

20

10

10

Energy [eV]Fig. 1. The energy spectrum observed by AGASA. The vertical axis
is multiplied by E . Error bars represent the Poisson upper andlower limits at 68% and arrows are 90% C.L. upper limits. Numbers attached to points show the number of events in each energy bin. The dashed curve represents the spectrum expected forextragalactic sources distributed uniformly in the Universe taking
account of the effect of energy resolution.

AGASA + A20
+60o

Dec
o

+30

24

h

h

RA

12

0

h

Galactic Plane

-30
-60

o

o

Supergalactic Plane
Fig. 2. The arrival direction of the highest energy cosmic rays above
4  10 eV in the equatorial coordinates.

In Fig. 3 [14], we show the space angle distribution between two cosmic ray events above 4 1019 eV. We observea clear peak at small correlation angles below 3Æ . The shapeof the peak is consistent with the resolution of the angulardetermination of AGASA. This result suggests the existenceof point like compact sources. The magnetic field in the intergalactic spaces is required to be nG or less for charged
particles to form these clusters.
Anisotropy around 1018 eV
The harmonic analysis in the right ascension was carriedout using about 216,000 events of cosmic rays observed by
AGASA [15]. The value of k  14 at 1018 eV is seen in the17

Density [arbitrary unit]

The Highest Energy Gamma Rays

0
0

20

40

Separation Angle [deg]Fig. 3. The space angle distribution between two events above
4 10 eV. We observe a clear auto-correlation within three degrees.

The highest energy gamma ray primary is a good discriminator of EHECR origins. The top-down theories such as the
decay of topological defects and super-heavy particles predicta high level of gamma ray composition whereas the bottom-up
theories predict none. Since the gamma ray has less probability of generating hadronic particles, the muon content of theshower is expected to be significantly less compared to theproton or nucleus primaries. The muon density in the EHECR
air showers was measured using 27 proportional counters of
2.8-10 m2 area installed in the region of AGASA array. The
muon density 1000 m away from the shower core was determined by fitting the lateral distribution of the muon component. The result was compared with the prediction of Monte
Carlo simulations with gamma ray, proton and iron primaries.
As a result, upper limits on the gamma ray fraction with 95%
confidence level are determined to be 34%, 59% and 63% forcosmic rays with energies more than 1019 eV, 1019 25 eV and
1019 5 eV [16].

Declination[degree]

From AGASA to Telescope Array
90

4
3

30

-30

2

anti G.C.

1

AGASA had ceased taking data on January 5th, 2004 after 13 years of fruitful operation. As described in this report,
its contribution to the physics of ultra-high energy cosmic ray
was remarkable. The analysis using the whole data sample
is now proceeding. We continue to build a succeeding experiment, Telescope Array, in order to identify the origin ofcosmic rays over 1020 eV discovered by AGASA.

0
G.C.

-1

Bibliography-2 [1]
-90

-3
300

200
100
Right Ascention[degree]0

N. Chiba et al., Nucl. Instrum. Methods A 311, 388(1992).

[2] H. Ohoka et al., Nucl. Instrum. Methods A A372, 527(1997).

Fig. 4. The statistical significance of the deviations of the arrival direction distribution of the cosmic rays above 10 eV in the equatorial
coordinates.

[3] K. Greisen, Phys. Rev. Lett. 16, 748 (1966); G. T. Zatsepin and V. A. Kuzmin, Pisma Zh. Eksp. Teor. Fiz. 4,
144 (1966).

first harmonic. The phase of the maximum amplitude corresponds to 20 hr in the right ascension. The statistical significance is high, corresponding to a chance probability of 10 6.
In Fig. 4, the arrival direction distributions in the equatorial coordinates are shown. They show significant deviations
from the expected isotropic distribution. Here, the energy region of 1018 eV  1018 4 eV is selected which maximize thek-value of the harmonic analysis. The AGASA is insensitivefor events with declination less than 25Æ as long as we use
showers with zenith angles less than 60Æ . We have chosen acircle of 20Æ radius to evaluate the excess in Fig. 4. A 4.5 σ
excess (obs./exp. = 506/413.6) near the galactic center region
can be seen in the figure. We also observe a deficit corresponding to 4 0 σ in the vicinity of anti-galactic center. An
excess of 3.9 σ significance (obs./exp. = 3401/3148) is seenin the Cygnus region.

[4] Various acceleration models are discussed in Astrophysical Aspects of the Most Energetic Cosmic Rays, ed. M.
Nagano and F. Takahara (World Scientific, Singapore,1990) pp. 252–334.
[5] E. Waxman, Phys. Rev. Lett. 75, 386 (1995); M. Vietri,
Ap. J. 453, 883 (1995); M. Milgrom and V. Usov, Ap. J.
449, L37 (1995);[6] Recent summary and references are described in P. Bhattacharjee, Proc. of ICRR Symp. on Extremely High Energy Cosmic Rays : Astrophysics and Future Observatories ed. M. Nagano, (Inst. of Cosmic Ray Research,
University of Tokyo., 1997) p. 125.
[7] A. M. Hillas et al, Proc. 12th ICRC, Hobart 3, 1001
(1971).
[8] H. Y. Dai et al., J. Phys. G: Nucl. Part. Phys. 14, 793(1988).

18[9] S. Yoshida et al., J. Phys. G: Nucl. Part. Phys. 20, 651(1994).
[10] M. Takeda et al., Phys. Rev. Lett. 81, 1163 (1998).
[11] M. Takeda et al., Proc. 28th ICRC, Tsukuba 1, 381(2003).
[12] M. Takeda et al., Astroparticle Phys. 19, 447 (2003).
[13] M. Takeda et al., Ap. J. 522, 225 (1999); astroph/0008102.
[14] M. Teshima et al., Proc. 28th ICRC, Tsukuba 1, 401(2003).

Fig. 1. Detector Arrangement of Phase-1 Hybrid TA.

[15] N. Hayashida et al., Astroparticle Phys. 10, 303 (1999).
[16] K. Shinozaki et al., Ap. J. 571, L117 (2002).

Seeking for an Origin of Super-GZK Events— Telescope Array Experiment[Spokesperson: M. Fukushima]
High Energy Cosmic Ray Div. ICRR, Univ. of Tokyo,
Kashiwa, 277-8582 Chiba
In Collaboration with the members of:
Chiba University, Chiba, Japan; Communications ResearchLaboratory, Tokyo, Japan; Ehime University, Matsuyama,
Japan; Hiroshima City University, Hiroshima, Japan; ICRR,
Univ. of Tokyo, Tokyo, Japan; KEK, High Energy Accelerator Organization, Tsukuba, Japan; Kinki University, Osaka,
Japan; Kanagawa University, Yokohama, Japan; Kochi University, Kochi, Japan; Montana State University, Bozeman
MT, USA; Musashi Institute of Technology, Tokyo, Japan;
National Institute of Radiological Sciences, Chiba, Japan; Osaka City Univ., Osaka, Japan; Rutgers University, Piscataway,
New Jersey, USA; Saitama Univ., Saitama, Japan; Shibaura
Institute of Technology, Saitama, Japan; Tokyo Institute ofTechnology, Tokyo, Japan; University of New Mexico, Albuquerque, New Mexico, USA; University of Utah, Salt Lake
City, Utah, USA; Yamanashi Univ., Kofu, Japan;

Extremely High Energy Cosmic Rays
The origin of the extremely high energy (EHE) cosmic
rays is an enigma. Eleven events observed by AGASA clearly
demonstrate the existence of particles above 1020 eV in theuniverse. As predicted by Greisen Zatsepin and Kuzmin(GZK), such particles collide with the microwave backgroundand quickly loses its energy by the pion photoproduction. The
universe is not transparent for EHE particles, and we have tolook for the origin of such particles within 50 Mpc (= GZKhorizon) of our galaxy.
Powerful extragalactic objects such as the active galactic
nuclei, radio galaxies and the gamma ray bursts may be suspected as the acceleration site of EHE proton or nucleus. The
search for the astronomical origin in the arrival direction of
the AGASA events, however, failed to find such objects within
100 Mpc of our galaxy. The origins beyond this limit may betolerated, but only if we assume special mechanism to allow a

longer propagation beyond the GZK limit, such as the violation of special relativity or the EHE neutrinos as the carrier ofsuch energy. In addition, it is generally perceived that the astronomical acceleration above 1020 eV is difficult to achieve.
It is fascinating to imagine that the EHE cosmic ray is
generated by the decay of super-heavy elementary particles
in the present universe. The energies above 1020 eV are easily
obtained if the particle mass is at the Grand Unification scaleof 1015 GeV. Such particles may be either surviving as a relicparticle of the Big Bang or presently generated by the decay of
topological defects. An abundant generation of EHE gamma
rays and neutrinos, in place of protons and nuclei, is expectedin the decay of such particles.

Overview of the Experiment
The Telescope Array (TA) experiment is proposed to investigate the nature of super-GZK cosmic rays and draw a decisive conclusion on its origin.
In the first stage of construction, TA is a hybrid detector
composed of a large plastic scintillator array and 3 sets of air
fluorescence telescopes installed around the ground array andoverlooking it (phase-1 Hybrid TA, see Fig. 5)[1,2,3]. The
ground array will give an aperture approximately an order ofmagnitude larger than that of AGASA. The fluorescence telescope will cross-calibrate the shower parameters measured by
the ground array by employing a direct calorimetry in the atmosphere and by imaging the shower axis and shower longitudinal development. It will also supply information on theparticle species of the primary cosmic rays by measuring thedepth of the shower maximum (Xmax ). The whole detector
will be installed in the West Desert of Utah, USA, about 140mile south west of Salt Lake City.

Hybrid TA : Ground Array
The ground detector consists of an array of 576 plastic
scintillators deployed in a grid of 1.2 km spacing covering
the ground area of 760 km2 [4]. The detector acceptance isconstant for cosmic rays with energies more than 1019 5 eVentering the detector with zenith angles less than 45Æ .
The scintillation light is read out by a wave length shifter
fiber of 1 mm diameter installed in a groove produced on
the surface of the scintillator. One layer of the scintillation

19

Fig. 2. Scintillation detector of TA

counter consists of two sets of 4 plates of 150 cm25 cm andwith a thickness of 1.2 cm shown in Fig. 6 and is read out by
96 fibers. The final scintillation counter is composed of twolayers of 3 m2 , each of which is read out by a separate PMT.
A coincidence is taken between 2 layers for obtaining a stable
trigger condition against environmental background and PMT
noise.
The signal from the PMT will be continuously digitizedwith a 12-bit flash ADC with 50 MHz sampling. For a PMT
signal exceeding 1/3 of the muon peak, a complete wave formof 4 μ s is stored locally in a memory with a time stamp
supplied by the Global Positioning System (GPS). This rateof local buffering is expected to be less than 1 kHz for the
3 m2 counter. The relative timing between remotely separated
counters should be maintained within 20 ns by the GPS forthe good resolution of arrival direction. A list of triggered
events containing the pulse height and timing information of
the hits is sent to a central DAQ system every second. If a
trigger condition of 3 or more muons is adopted, the trigger
list would contain 100 events. An air shower event is identified by the central DAQ software by requiring clustered hitsof counters with a good timing of coincidence. The air shower
event rate will be significantly less than 1 Hz when we requireat least 3 counters are hit at the same time.
The data acquisition system will be composed of a wireless local area network (LAN) using 2.4 GHz spread spectrum
technology. Considering a limited reach of the presently available models, two stages of data concentration will be necessary to reach the central DAQ from individual counters. The
speed of less than 1 Mbps is, however, sufficient for the expected trigger rate and DAQ throughput even when 2 or morestages of data collection are necessary. Five communication
towers would be needed: three for nearby telescope stations
and two for the places where communication may be difficult
by only three towers.
The total electrical power consumed by the PMT, ADC,
GPS and LAN will be of 5 W. The scintillation detector is
locally generated by the solar panel of  120 W, installed on
top of the platform at a 60 degree angle. Behind and below thepanel will be a metal enclosure containing the storage battery
and electronics. A 10-feet antenna would be attached to eachplatform. A simple and robust counter design together withan independent GPS, wireless-LAN and local power generator will make the system easy to deploy and maintain in theFig. 3. A view of a TA fluorescence station in Utah.

desert.

Hybrid TA : Fluorescence Telescope
The fluorescence measurement is made at 3 stations surrounding the ground array[5]. The stations form a triangle
with a length of sides 3040 km. Twelve reflecting telescopes
are installed at each station and cover the sky of 3Æ  34Æ inelevation and 108Æ in azimuth looking toward the center of theground array (see Figures 5 and 7). The stereo acceptance ofthe detector is 670 km2 sr for E  1020 eV by requiring atleast one station is within 45 km from the shower center. The
acceptance is 4 times larger than that of AGASA even assuming 10% duty factor. Approximately 60% of the events for E
20 eV are observed by all 3 stations.
 10
The field of view of each telescope is 18 0Æ in azimuth
and 15 5Æ in elevation. The telescope has a spherical dish of
3.3 m diameter composed of 18 hexagonal segment mirrors(see Fig. 8). The focal length is 2960 mm and the spot size on
the focal plane is less than the specification (20 mm) in the full
width for which 90% of light is concentrated according to themeasurement of the mirrors for the first telescope station. The
air shower image is recorded by a camera composed of 16162-inch PMTs placed on the imaging plane (see Fig. 9). EachPMT pixel covers 1.1Æ 1.0Æ patch of the night sky. The shapeof the photocathode is made hexagonal for the maximum lightcollection efficiency. The PMT high voltage is supplied by ableeder with zener diodes to ensure a stable gain ( 2%) underthe variation of the night sky background, and a good linearity( 5%) up to 16 k photoelectrons in 200 ns. The 8-stage PMT(Hamamatsu R6234) is operated at 800 Volt with a gain of
105.
A signal from the PMT is amplified by a factor of 4 with a
pre-amplifier and is sent to a Strecher Amplifier by a factor of
2 and a 12-bit ADC with 40 MHz sampling in a Signal Digitizer and Finder (SDF) module. The digitized four successive
signals are summed and sent to a data buffer, which is adopted
to cover a dynamic range corresponding to 32 K photoelectrons with the aid of an analog sum of 16 channels with thegain of 1/16. The digitized signal is further fed to a Signal
Finder part to search for a fluorescence signal in 25.6 μ s time

20

Fig. 6. Test setup of readout and trigger for the fluorescence telescope.

Fig. 4. Telescope installed in the telescope building in Utah

Fig. 5. Prototype Camera of TA

window. The timing and pulse height of all PMTs are collected for one camera and a shower track is searched from hit
pattern in a Track Finder (TF) module. Test setup of the readout and trigger system is shown in Fig. 10. The information ofa trigger is broadcasted to all cameras in the station througha Central Trigger Decision (CTD) module and data are collected via Ethernet to build a complete event.

For the precise measurement of energy, the absolute calibration of the telescope sensitivity is important. We have tounderstand the efficiency of telescope optics, mirror reflectivity, filter transmission, PMT quantum efficiency and the electronics sensitivity. Each component will be calibrated pieceby piece at the production stage and will be monitored insitu by various light emitting pulser system such as a diffused
Xenon lamp, a blue LED and a tiny YAP (YAlO3 :Ce) scintillator with 241 Am source. The present uncertainty of the PMT
quantum efficiency (QE) and the dinode collection efficiency(CE) of produced photoelectrons is each 10% and the systematic uncertainty of PMT sensitivity reaches to 14% in
total according to the producer. We made a simple pulsed
UV light calibration source with 5% absolute intensity error
by using a Rayleigh scattering of nitrogen laser light by purenitrogen gas molecules. We checked the value of company
supplied PMT efficiency by using this system. The preliminary result shows that the combined error of QE+CE can bemade at 5% level and the actual error of the company supplied
QE and CE values are smaller than quoted by the company.
The ultra-violet fluorescence light generated by the air
shower in the atmosphere is scattered and lost along the pathof transmission to the telescope. The responsible processes
are Rayleigh scattering by the air molecule and Mie scattering by the aerosol. The Rayleigh scattering is elementary
and the loss can be calculated and corrected with an accuracy better than 5% knowing the density distribution of theatmosphere. The amount of Mie scattering differs from place
to place and changes with time reflecting the aerosol distribution in the air. The primary task of the atmospheric monitoring is to determine the loss by the Mie scattering with anaccuracy better than 10%. It is particularly important for cosmic rays dropping in the region of the ground array. These
events will be measured by the telescope and the ground array independently, and are important means to cross-calibrate
energies measured by both systems. The atmospheric monitoring will be made by the lidar system located near each

21

Fig. 7. Lidar system located near the telescope station in Utah.

station. It consists of a pulsed Nd:YAG laser and a light receiver attached on an alt-azimuth mount (See Fig. 11). Using
this system, we can shoot laser at any direction and measure
the 3-dimensional distribution of atmospheric extinction coefficients within 10 km of each station[6]. An infra-red CCD
camera is mounted on the same alt-azimuth and monitors theexistence of low clouds in the night sky. A few more additional laser guns will be located in the center of the groundarray area. The analysis of the side scattered light from theseguns will be made by using the absolute calibration of the laser
power and the light receiver. Field tests of prototype monitoring system have been made since 1997 in Utah and Akeno.
The most critical work of the air shower experiment as acalorimeter is to calibrate absolute energy scale of the detector
at the energies of our concern, namely at  1020 eV. One suchpossibility of end-to-end absolute energy calibration is to inject an electron beam of energy 10100 MeV from the linac
vertically up in the atmosphere. If we take 40 MeV electron
beam with 109 particles per bunch, the total energy is equivalent to 1017 eV. Such amount of fluorescence lights generated 100 m away is clearly detectable by the fluorescencetelescope. The energy of the linac beam is deposited in the airby the low energy electrons and gammas. Note that the situation is similar to the actual air shower; a large part of the air
shower energy is deposited in the air by the low energy electrons and positrons below the critical energy. We are in theprocess of designing such calibration device for the hybrid TA
experiment.
Power for these fluorescence detectors will be supplied bydiesel generators. Fuel for the genrators will be stored on-sitein large tanks.

Construction
The phase-1 Hybrid Telescope Array will be built by thecollaboration of Japanese and American physicists. The groupconsists of cosmic ray physicists who have been working in
AGASA and HiRes experiments, and high energy physicists
who worked in the accelerator experiments in US and Japan.
The construction of TA started in the summer of 2003 by the
Grants-in-Aid for Scientific Research (Kakenhi).
On December in 2004, 20 scintillation ground detectors were deployed by helicopter for the engineering test.
Figure 12 shows one of the detectors deployed in the field.
Data transfer was successful between the detector and the temporary communication tower. The timing distributions of sig-

Fig. 8. A scintillator ground detector deployed in the field.

Fig. 9. The timing distributions of signals from the scintillator detector
by a cosmic ray muon.

nals from the scintillator detector by a cosmic ray muon are
shown in Fig. 13.
On October 2004, the construction of one of the three telescope buildings was completed. The structure of twelve telescopes were installed early in 2005 and attached two composite mirrors as shown in Fig. 7. The test observation of air
shower using this ground array and the telescopes will be performed in 2005. The full deployment of scintillator grounddetectors will begin in mid-to late 2005. The construction of
the 2nd telescope station will also start in 2005. The full detector will be completed and operational in 2007.

Bibliography[1] The Telescope Array Project: Design Report. July,
2000, available by the web http://www-ta.icrr.utokyo.ac.jp/index.html.
[2] M. Fukushima, Proc. of APPI2003 workshop, KEK
Proceedings 2003-6.
[3] M. Fukushima et al., Proc. of 28th International Cosmic

22

Experimental Study of High-energy Cosmic
Rays by the Tibet Air Shower Array in the
Tibetan Highlands (The Tibet ASγ Collaboration)[Spokesperson : Masato Takita]
ICRR, Univ. of Tokyo, Kashiwa, Chiba 277-8582
In collaboration with the members of:Hirosaki Univ., Hirosaki; Utsunomiya Univ., Utsunomiya;
Saitama Univ., Urawa; Kanagawa Univ., Yokohama; Yokohama National Univ., Yokohama; Shonan Inst. of Technology,
Fujisawa ; Tokyo Metropolitan Coll. of Aeronautical Eng.,Tokyo ; NACSIS, Monbukagakusho, Tokyo ; Konan Univ.,
Kobe, Shibaura Inst. of Technology, Tokyo; China
Since 1990, the Tibet air shower array has been successfully operated at Yangbajing (4 300 m above sea level) in Tibet. Some results obtained from this experiment have beenreported elsewhere. The first air shower array (Tibet-I) was
updated in late 1994 by increasing the number of scintillation
detectors from 65 to 221, and then this array (called Tibet-II)
has been fully operating. The Tibet-II array consists of 185
first-timing (FT) plastic scintillation detectors of 0.5 m2 each.
With a FT detector, equipped is a plastic scintillator plate in
3 cm thickness each viewed by a 2 φ PMT (HPK H1161).
They are placed at a grid point with 15 m spacing, as in TibetI. Among these, each of the 52 FT detectors also contains a
1 5 φ PMT (HPK H3178) for a wide dynamic-range measurement of the particle density. The FT array is also surroundedby 32 density detectors of 0.5 m2 each to obtain a good core
location for an individual air shower event. The effective areafor detecting an contained air shower event is about 8 times
as large as Tibet-I. The air shower events have been accumulated at a rate of about 150 Hz to efficiently detect air showers
with energies around 10 TeV. In 1996, a high-density array
was constructed near the center of the Tibet-II array. This array (HD-array) consists of 110 plastic scintillation detectors,
each of which was deployed at a grid point with 7.5 m spacing. The HD array is sensitive to air showers with energies
around a few TeV. The trigger rate of the HD array has been
set to about 150 Hz. Using this HD array, in 1999, we succeeded in observing multi-TeV γ ray signals from the Crab
Nebula at 5.5 σ confidence level. This was the first detectionof multi-TeV γ -ray signals by a conventional air shower array.
Subsequently, we also detected multi-TeV γ rays successfully
at 3.7 σ level from Mrk501 which was in a highly flaring state
between March 1997 and August 1997.
In the late fall of 1999, the area of the HD array was enlarged up to 22000 m2 (Tibet-III). It is equipped with 533 plastic scintillation counters (497 FT + 36 density PMTs) in total.

Displacemet (N−S)
of the Moon (degree)

2000
Feb.26

0.4
0.2
0.0
−0.2
−0.4

Deficit rate by the Moon
( Observed / Expected )Ray Conference, Vol. 2,1025(2003).
[4] S. Kawakami et al., Proc. of 28th International Cosmic
Ray Conference, Vol. 2,1033(2003).
[5] F. Kakimoto et al., Proc. of 28th International Cosmic
Ray Conference, Vol. 2, 1029(2003).
[6] T. Yamamoto et al., Nucl. Instr. and Methods A488,
191(2002).

Jun.5

Sep.13 Dec.22

2001
Apr.1

Jul.10

Oct.18North

χ2 / dof = 7.8 / 18

South

1.4
1.2
1.0
0.8
0.6

χ2 / dof = 16.8 / 18

51500

51600

51700

51800 51900
MJD

52000

52100

52200

Fig. 1. Tibet-III long-term stability in pointing accuracy and angularresolution.

The trigger rate has been set approximately to 680 Hz.
The performance of the Tibet air shower array has been
well examined by observing the moon’s shadow in cosmic
rays, and the deficit map of cosmic rays around the moon
demonstrates the angular resolution to be around 0.9Æ at a few
TeV for the Tibet-III array. The long-term stability in pointing
accuracy (less than 0.02 o ) and angular resolution is demonstrated in Fig. 1.
Multi-TeV γ -ray signal was successfully detected at 5σlevel from the Crab (the standard candle in γ -ray astronomy) by the Tibet-III array. We also succeeded in observing multi-TeV gamma-ray flares at 5.1σ level from Markarian
421 which was in a very active phase during the year 2000
and 2001. We observed a clear long-term positive correlation
between X-ray data (RXTE satellite) and our TeV γ -ray datafrom Mrk421 in the active period.
We also searched for multi-TeV diffuse γ rays from thegalactic plane. As there was no significant signal, flux upper
limits were obtained from the inner galaxy (20o  l  55o,
2o  b  2o) and the outer galaxy (140o  l  225o, 2o b  2o ).
Then, autumn 2002, the Tibet-III array was further extended to 37000 m2 with 733 plastic scintillation counters byadding 200 more FT counters (i.e., 697 FT + 36 density PMTs
in total), as is shown in Fig. 2. The trigger rate has been set to1500 Hz at a few TeV energy threshold. In 2003, 56 FT detectors were further added in the outer region of the array in order
to expand the effective high-density area up to 37000 m2 . The
37000 m2 Tibet-III with 789 counters are in continuous operation. The trigger rate is 1700 H.z We are now accumulating
Tibet-III data to obtain a standard TeV γ -ray flux from the
Crab and to search for unknown constant/transient TeV γ -raysources.
A hybrid experiment of emulsion chambers (EC) and air
shower array started in 1996 to obtain the energy spectrumof the primary cosmic-ray proton flux around the knee energy
region The total area of EC, each having the size of 40 cm 
50 cm, is 80 m2 and the total thickness of lead plates is 14
r.l. High-sensitivity X-ray films were interleaved between the
14 r.l. lead plates at every 2 r.l. to detect γ -ray families. Just
below the emulsion chambers, the burst detectors (BD) withthe same area as EC were set up to locate the air shower coresof the family events to be observed in EC. This detector complex was set up near the center of the Tibet-II array (AS) to geton the information on air showers accompanied with the family events. The first EC exposure was terminated in August

23

all particles

dj/dE x E2.5 [eV1.5 m−2sr −1s−1]

1018

Fig. 2. 37000 m2 Tibet-III array operating at Yangbajing.

17

10

Tibet−III(2003)
KASCADE(2003)
EAS −TOP10BASJE−MAS(2003)
Akeno
BASJE
CASA−MIA
DICE
15
HEGRA
10
13
12
16

10

10

JACEE
MSU
protonRUNJOB
SOKOL
Sphere
TACT
Tunka

Tibet−I

1014

15

16

10

17

10

10

Primary Energy [eV]PROTON satellite
Tibet-III
AKENO

10

10

Fig. 4. Primay all-particle energy spectrum in the knee energy region.

All - (P + He)
All

1.5

E2.5 × dJ/dE (GeV /sr/m2/s)

10

4

3

1

0.9
0.8

-2.74
2

This work(QGSJET+HD)
This work(QGSJET+PD)
This work(SIBYLL+HD)
This work(SIBYLL+PD)
Tibet-B.D.(QGSJET+HD)
Tibet-B.D.(QGSJET+PD)
KASCADE(QGSJET)KASCADE(SIBYLL)
JACEE
RUNJOB

10
10

4

10

5

0.7

-3.06
-3.11
-3.13
-3.11

0.6

This work(QGSJET+HD)
This work(QGSJET+PD)
This work(SIBYLL+HD)
This work(SIBYLL+PD)
KASCADE(QGSJET)KASCADE(SIBYLL)

0.5
0.4
0.3

10

6

7

8

10
10
Energy (GeV)

Fig. 3. Primay proton energy spectrum in the knee energy region.

0.2
0.1
0
10

of 1997 and X-ray films inserted in EC were developed foranalysis. A high-energy family event of about 500 TeV wasobserved in this exposure and its primary energy is estimatedto be about 1016 eV from the size of the accompanying AS
data. This hybrid experiment continued until 1999. Using the
BD + AS data, we obtained the energy spectrum of primary
protons (820 proton-induced events during 690-day detector
live time) with its primary energies 200–1000 TeV by a neuralnetwork method. The differential energy spectrum obtainedin this energy range can be fit by a power law with the spectral index of 2.97 0.06, which is steeper than that obtainedby direct measurements at lower energies. We also obtained
the energy spectrum of helium nuclei. All of the features ofBD events are wholly compatible the heavy enriched composition in the knee region. In 2003, we finished the analysisof all the EC data by means of an automatic analysis program
based on track information read out by the image scanners.
The primary proton spectrum including the EC + BD + AS
data around the knee region, being inaccessible by any direct
observations, are being analyzed with the 3-year data obtainedby the hybrid experiment and a preliminary result is shown in
Fig. 3.
We also obtained a preliminary result on the primarycosmic-ray all-particle energy spectrum, as shown in Fig. 4,

4

10

5

10

6

7

8

10
10
Energy (GeV)

Fig. 5. Primay (all-particle-(proton+helium)/all-particle flux ratio in theknee energy region.

which is consistent with the one obtained by the Tibet-I experiment. The (all-particle - (proton+He))/all-particle flux ratio is shown in Fig. 5, indicating that the knee is composed of
nuclei heavier than helium. It should be noted that the flux ratio largely cancels out the systematic energy scale uncertainty
in the airshower energy determination.
Based on the Tibet-III data from 1999 and 2003, we succeeded in observation of the multi-TeV galactic cosmic-ray
anisotropy at solar time frame due to the terrestrial orbital
motion around the Sun, i.e. Compton-Getting (C-G) effect.
As shown in Fig. 5, we observed a clear C-G anisotropy inthe 6.2 and 12 TeV data samples, while the anisotropy observed in the 4 TeV data sample deviates from the expected
C-G anisotropy. This suggests an additional anisotropy superposed at the multi-TeV energies, e.g., the solar modulation
effect.
As discussed elsewhere, the Tibet air shower array is very
powerful to get new information on the relation between time
variation of the large-scale structure of the solar and inter-

24
4.0 TeVvariation [%]

0.04
0.02
0.00
−0.02
−0.04

6.2 TeV

12 TeV
(a)

0.04
0.00
−0.04
−0.08
0.04
0.00
−0.04
−0.08

(b)

(c)
0

6

12

18

24 0

6
12
18
24 0
local solar time [hour]

6

12

18

24

Fig. 6. Multi-TeV cosmic-ray anisotropy at solar time frame obtainedby Tibet-III.

1990 (Tibet-I)

1996 (Tibet-II)

2000 (Tibet-III)

1991 (Tibet-I)

1997 (Tibet-II)

2001 (Tibet-III)

1992 (Tibet-I)

1998 (Tibet-II)

2002 (Tibet-III)

1993 (Tibet-I)

1999 (Tibet-II)

2003 (Tibet-III)

Fig. 7. Yearly variation of the Sun’s shadow observed by the Tibet air
shower array from 1990 to 2003. Mode energies are 10 TeV.

planetary magnetic fields and the solar activities, since highstatistics data taken by the Tibet air shower array can follow up the movement of the sun’s shadow at every one-twomonths. The sun’s shadow was observed in the direction significantly away from the optically observed sun’s position during the period from 1990 through 1993. Note that this period
corresponded to the near-maximum or at the decreasing phasein the solar cycle 22. In 1996 and 1997, however, we foundthat the sun’s shadow was observed in the direction close to
the optically observed sun’s direction, since the solar cycle
was in a quiet phase then. Since 1998, the sun’s shadow began to be obscure shown in Fig. 6, as the solar activities wereback in an active phase. We expected that the sun’ shadow
would change its position again, according to the solar activities as the next solar cycle (Cycle 23) goes toward its maximum around the year of 2001 or 2002. As the maximum
showed a double-peak structure and the solar activities werestill in high states in 2002, we have not observed a clear sun’s
shadow yet, although the sun starts exhibiting a darker shadow
in cosmic rays in 2003. Thus, the results will contribute considerably to the study of solar terrestrial physics.
This group has developed and completed several automatic measuring systems that are powerful for analyzing cosmic ray tracks or air shower spots, that is, automatic microdensitometers and precise coordinate-measuring systems controlled by a computer. Enormous data recorded on nuclear

emulsion plates or X-ray films are rapidly and precisely measured by the use of these measuring systems. At present, five
workstations are introduced to improve data-processing capacity.
The emulsion-pouring facilities can meet the demands formaking any kind of nuclear emulsion plates which are used
for cosmic ray or accelerator experiments. The thermostatic
emulsion-processing facilities are operated in order to develop
nuclear emulsion plates or X-ray films. Using these facilities,
it is also possible to make and develop emulsion pellicles in
600-micron thickness each. In this way, these facilities areopen to all the qualified scientists who want to carry out joint
research program successfully.
A scintillation-fiber calorimeter has been developed to beloaded on satellites/balloons in the future for the observationof GeV - TeV γ rays/electrons.

List of Other Inter-University Research Programs
Code No.: 8Title: Sidereal daily variation of 10 TeV galactic
cosmic-ray intensity observed by the Tibet air shower array.
Spokesperson: K. Munakata (Shinshu University)
Participating Institutions: Shinshu Univ.
Code No.: 9Title: Observation of High-Energy Electrons and Gamma
Rays with a Balloon-borne Instrument
Spokesperson: Shouji Torii (Kanagawa University)
Participating Institutions: Kanagawa Univ., ISAS,
Rikkyou Univ.,
ICRR Univ. of Tokyo, and Shibaura Inst. of Technology.
Code No.: 10Title: Observation of high-energy primary electrons andatmospheric gamma rays by emulsion chambers.
Spokesperson: T. Kobayashi (Aoyama Gakuin University)
Participating Institutions: Utsunomiya Univ., Kanagawa
Univ., ICRR Univ. of Tokyo, Univ. of Tokyo.

Development of Wide Angle High Resolution Detector[Spokesperson : M. Sasaki]
ICRR, Univ. of Tokyo, Kashiwa, Chiba 277-8582
ICRR, Univ. Tokyo; IfA, Univ. Hawaii; Ibaraki Univ.; KEK;
National Astronomical Observatory; National Taiwan Univ.;
National United Univ.; Toho Univ.; Tohoku Univ.; Tokyo
Institute of Technology; Univ. Hawaii Hilo; Univ. Hawaii
Manoa.
Ashra Project Summary
The activities herein contribute to the technology development and demonstration of the All-sky Survey High
Resolution Air-shower detector (Ashra; http://www.icrr.utokyo.ac.jp/ashra) for the advancement of particle astrophysics. Ashra records images in unprecedented arc-minute25

detail of high-energy cosmic particle interactions in the atmosphere using new CMOS technology. Tracks of nitrogen fluorescence and beams of Cherenkov radiation reveal the arrival
direction, energy, and identity of cosmic rays (neutrinos, nuclei and gamma rays) over an energy range of 7 orders of magnitude, from TeV to EeV scale and beyond. Ashra sees not
only near horizontal air showers due to neutrinos, but thoseshowers exiting the nearby mountains due to conversion of
traversing tau neutrinos, and includes a view of the Galactic
Center region. Moreover, the detector can be used to seek optical flashes from Gamma Ray Bursts, providing early alert
to telescopes for more detailed observation, particularly thoselocated nearby.
Ashra consists of three high-altitude (3 km) stations to be
incrementally installed at Mauna Loa, Hualalai, and Mauna
Kea on the island of Hawaii. Each station, separated by 30–
40 km, observes the entire moonless night sky with 12 detectors. A detector utilizes four 2.2 m projected diameter
spherical reflectors for viewing a 50Æ -circle (0.5 sr) region ofsky. Detector images pass through a photoelectric pipeline fortrigger processing and recording on 4 mega-pixel CMOS sensors. Covering the sky with 48 M pixels, roughly a thousandfold improvement over presently employed technology, represents resolution measured in arc minutes rather than degrees(as achieved with photomultipliers). This unprecedented pixelization provides shower reconstruction with better pointing
accuracy (arc-minutes), improved energy determination, andan extension of the range to which distant high energy showers can be detected (because of being able to resolve shorter
tracks on the sky).
The increased pixelization will pay off in energy, particle type discrimination, origin direction and effective volume. The most unique feature of Ashra however is the ability
to serve as simultaneous detector for both the long duration(100 μ s) fluorescent events and short (few ns) air Cherenkov
events, not possible with earlier instruments. The problem ofintegrating background light over long times is solved by animage delay system to the CMOS sensors, and a moving electronic shutter driven from a parallel multi-anode photomultiplier trigger.

Fig. 1. The 2/3 scale Ashra prototype telescope operated on
Haleakala.

Fig. 2. Example of a 50-degree FOV image in which the constellation of Taurus and Orion can be clearly identified. The inset of a2-degree square window shows a close-up view of the Pleiades.

The major gain here is the ability to monitor the whole sky
continuously for extended (dark) periods, making an unprecedented search for new point (and time varying) sources suchas the numerous gamma ray bursts, soft gamma ray repeaters,
unidentified EGRET sources and so on. Note that the effective area of Ashra for such a search is much larger than water
Cherenkov based detectors.
All Ashra systems now exist in the prototype stage andfirst versions have been assembled on Haleakala since Fall
2004 (Fig. 1). Land acquisition on Mauna Loa via the University of Hawaii is well underway. It is planned to install thefirst station on Mauna Loa beginning in 2005 and completed
in 2006. Major detector components include optics, image
pipeline, trigger, imaging sensors, and data acquisition hardware. In the Ashra optics, photons from air showers are collected using modified Baker-Nunn optics (Sasaki et al., 2002).
The fine resolution (arcminutes) in the wide field of view(0.5 sr) has been already demonstrated (Fig. 2). Adding those,
we successfully made a cross-observation of the field covering the HETE-2 WXM error box of GRB041211 continuously
between the time 1h7m before GRB041211 and that 1h41m
after GRB041211 taking 2000 images covering the WXM error box every 5 s with 4 s exposure time. We detected no
new objects in the WXM error box resulting the 3-sigma limiting magnitudes are stringently derived (GCN Report #2846,
see also http://www.icrr.u-tokyo.ac.jp/ashra/GRB041211/).
A demonstration alt-az mounted air Cherenkov detector on
Haleakala is in place and has also been collecting data with
the image pipeline and the self-trigger system.

26

Chacaltaya Observatory of Cosmic Physics[Spokespersons : H. Yoshii and F. Kakimoto]
In collaboration with the members of:
Ehime Univ., Matsuyama; Tokyo Inst. of Tech., Tokyo;
The Inst. of Phys. and Chem. Res., Saitama; IIF, UMSA,
La Paz, Bolivia; Okayama Univ., Okayama; Nagoya Univ.,
Nagoya; National Astronomical Observatory, Tokyo; Kobe
Univ., Kobe; Musashi Inst. of Tech., Tokyo; MPI fur Physik,
Muncen, Germany.
The air shower experiment of BASJE (Bolivian Air
Shower Joint Experiment) group has been continued from1962 at Mt. Chacaltaya in Bolivia (Fig. 1). The main themeof the present experiment is to investigate the origin of primary cosmic rays around the knee region. For this study wemeasure the energy spectrum, the chemical composition andthe arrival distribution of primary cosmic rays using the MAS(Minimum Air Shower) array (Fig. 2).
The information of the chemical composition is also important for derivation of the energy spectrum, which is estimated from the observed air shower size spectrum by usinga Monte-Carlo simulation. In our experiment, the chemicalcomposition is determined with three different experiments.
First we measured arrival time distributions of air Cherenkov
light for cosmic rays from 1015 to 1016 5 eV. Since this method
has a better mass resolution, we could determine the mixing
ratios of p, He, CNO group, Ne-S group, and Fe as a functionof the primary energy [1]. However, since the energy range
is so limited in this method, second we observed lateral distributions of air Cherenkov light. Then we estimated the mixing ratios of p-He, medium heavy group and Fe in the energy
range from 1014 5 to 1016 eV [2]. These results are consistent
with each other at the over rapping energy range and the resultat the lowest energy range is consistent with that of the direct
observations as shown in Fig. 3. Notably these two experiments measure the longitudinal developments of air showers
at their early stages. So third we examined the observed equiintesity curves to estimate the chemical composition using theinformation of the longitudinal developments at the deeper atmosphere [3]. This result is also shown in Fig. 3 at the energy range from 1014 to 1016 eV and is also consistent with

Fig. 1. The Laboratory of Cosmic Ray Physics at Mt. Chacaltaya.

Fig. 2. The MAS array.

Fig. 3. The averaged ln A of primary cosmic rays.

the previous two results. As a result, the chemical composition becomes gradually heavy from 1014 eV to the knee region
and the main component at and above the knee is Fe. Taking
account of this result, we estimated the primary energy spectrum converting from the observed air shower size spectrum
as shown in Fig. 4. This spectrum shows the knee structurearound 1015 5 eV and the power index changes from 2.66 to
3.19 at the knee. Then we examined models of the cosmic
ray origin comparing with our results. At present, Biermann’smodel [4], in which the cosmic rays are accelerated at WolfRayet stars, is the most promising one and alternatively we
proposed a model. In our model the primary energy spectra
for each components are extrapolated linearly using the direct
observation data but assuming the existence of leakage fromthe galactic disc at a constant rigidity of 1014 5 V [5].
Anisotropy in the arrival distribution of primary cosmic
rays is very small especially in the energy range below theknee due to their gyro motion in the galactic magnetic field.
However if there is a powerful source of primary cosmic raysnear the solar system, we can expect the existence of cosmic

27

Fig. 4. The energy spectrum of primary cosmic rays.

Fig. 6. The arrangement of detectors of a new array at Mt. Chacaltaya.

Bibliography[1] Y. Shirasaki et al., Astroparticle Physics, 15, 357–381(2001)[2] H. Tokuno et al., in preparation for the submission.
[3] S. Ogio et al., Astrophys. J., 612, 268–275 (2004)[4] P. L. Biermann, Astron. Astrophys, 271, 649 (1993)[5] S. Ogio and F. Kakimoto, Proc. of 28th ICRC at
Tsukuba, HE, 315–318 (2003)[6] V. S. Ptuskin, in Astropysics of cosimic rays ed. by V. S.
Berezinskii et al., North-Holland (1990)[7] S. Ogio et al., Il Nuovo Cimento, 24C, 625–630 (2001)
Fig. 5. Nobs Nexp against the galactic latitude for each galactic longitude bin.

ray flow from the direction of the source. For example Ptuskincalculated the anisotropy amplitude as a function of the diffusion coefficient for several sources [6]. These amplitudes areabout the order of 0.1% and so that we must analyze data withextreme caution. Then we derived the ratio of Nobs (the number of observed showers) to Nexp (that of expected ones) as
a function of the galactic longitude for each galactic latitude
bin using our observed data of air showers with primary energies above 3  1013 eV [7]. This result is shown in Fig. 5and shows that there is enhancement of primary cosmic rays
at regions with the galactic longitude from 220 to 340 degrees
and with the galactic latitude around 0 degree. This result isvery interesting because this region includes the Vela SNR.
For the next step we start the construction of a new array toobserve air showers with the primary energies above 1016 eV.
The arrangement of the detectors is shown in Fig. 6. The newexperiment will be started in 2006.

28

ASTROPHYSICS and GRAVITY DIVISION
Overview

In collaboration with the members of: TAMA collaboration
NAOJ, Tokyo; KEK, Tsukuba; UEC, Tokyo; Osaka City
Univ., Osaka; Kyoto Univ., Kyoto; Osaka Univ., Osaka; Niigata Univ., Niigata)

Overview
A gravitational wave is a physical entity in space-time predicted by Einstein’s theory of general relativity. Its existence
was indirectly proven by the observation of PSR1913+16 by
Taylor and Hulse [1], who won the Nobel prize in 1993. However, nobody has succeeded to directly detect gravitational
waves. The theory of gravitation can be tested by the detection of gravitational waves. A gravitational wave detector isthe last eye of mankind to inspect the universe. In order todirectly observe gravitational waves, we have been developing a sensitive interferometric gravitational wave detector. We
succeeded in operating TAMA for 558 hours with a sensitivity by which we could make observations of the coalescenceof double neutron stars occurring within a range of 73 kpc
with a signal-to-noise ratio of 10 (DT9 in 2004). Since thefirst data-taking run (DT1), we have accumulated data of morethan 3100 hours in total. A facility for data analysis was madeat ICRR, Kashiwa in 2003, and the first approach on TAMA
data analysis was completed.
After DT9, the observational operation was stopped foralmost one year for noise hunting so as to attain the objective
sensitivity. We will be able to report the limiting factor ataround the frequency of the best sensitivity.

Term
6-Aug  7-Aug
17-Sept  20-Sept
20-Apr  23-Apr
21-Aug  4-Sept2-Mar  8-Mar
15-Aug  20-Sept
31-Aug  2-Sept
14-Feb  14-Apr
28-Nov  10-Jan

Year
1999
1999
2000
2000
2001
2001
2002
2003
2004

Length (Hour)
7
31
13
161
111
1038
25
1158
558

Status of TAMA Project
TAMA is a gravitational wave project to construct a 300 mbaseline laser interferometer at the Mitaka campus of the National Astronomical Observatory of Japan (NAOJ) and to conduct observations. TAMA is a five-year project, started in
April, 1995, and is organized by researchers belonging to universities and national laboratories. We regard the TAMA interferometer as a step toward the final scale interferometer inthe sense of technology and construction budget. We have so
far achieved nine data-taking runs that span from two to fourweeks (Table 1). In the former half runs, its sensitivity improvement was the first priority. Nonetheless, both the stability and reliability had to be improved to check the sensitivity,
itself. In the latter half runs, the operation became easier after
installing an automatic control system.
-10

10

-11

10
]

[Spokesperson : Kazuaki Kuroda]
ICRR, Univ. of Tokyo, Kashiwa, Chiba 277-8582

Run
DT1
DT2
DT3
DT4
DT5
DT6
DT7
DT8
DT9

1/2

TAMA Project

Table 1. TAMA data-taking runs including long-term observations

Displacement noise [m/Hz

Astrophysics and Gravity Division consists of Gravitational Wave Group, The Sloan Digital Sky Survey Groupand Theory Group. The Gravitational Wave Group conducts
TAMA project jointly with researchers of gravitational wave
experiment and theory in Japan. The Group also conducts
a CLIO project that aims to practically test the cryogenic
laser interferometer system underground in Kamioka as one of
R&Ds for LCGT project to detect gravitational wave events.
The Sloan Digital Sky Survey Group continues accumulating data of images and spectroscopic observation of galaxies and publishing papers in collaboration with worldwide researchers. Theory Group introduced a new professor who
conducts theoretical study of the Universe in addition to theexisting research group of astroparticle physics.

2001/08/03 DT6
2003/02/20 DT8
2003/11/04 DT9

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10

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10

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10

-15

10

-16

10

-17

10

-18

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-19

10

1

10

2

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3

10
Frequency [Hz]

10

4

10

5

Fig. 1. Before operation stop for noise-hunting in 2004 the sensitivity was improved after the power recycling mirror installation. The
sensitivity was physically limited by the photon shot noise at frequencies of more than 1 kHz. Since seismic noise in lower frequencies than 30 Hz disturbed the stable operation of the interferometer, a tighter feedback gain of the mirror alignment control
loop was adopted, which resulted in an actuator force noise.

29

FP Cavity
300m arm

FP Cavity
300m arm

Laser
Power recycling m irror

Fig. 2. Optical configuration of TAMA300 as the final stage (DT7-DT9).
The reflectance of the recycling mirror was chosen to obtain apower recycling gain of 3 at this moment.

Distance of detecting inspirals with SNR=10100
Observable Distance [kpc]During the course of this project, we established techniques necessary for constructing a large-scale and highly sensitive interferometer. From analyses of DT8 and DT9, no
traces of gravitational waves were detected. The previous result was published, and its stable operation was reported [2][3]. The sensitivity at DT9 is shown in Fig. 1. At the timeof DT6, we had analyzed several noise sources that limited
the sensitivity, and almost all sources were identified, except
for the noise that limits the sensitivity at frequencies around
600 Hz. After DT6, TAMA entered the Phase II stage, where a
power recycling mirror was installed to reduce the shot noise.
The recycling mirror was installed between the beam splitterand the mode cleaner (in Fig. 2). The role of this mirror isto recycle the laser power returning from the beam splitter bymaking the phase coincident with the input laser beam fromthe entrance. At this stage, the power-recycling gain was set
at three, but would be finally set at 10. We could observe thecoalescence of neutron-star binaries occurring within 73 kpc
with a signal-to-noise ratio of 10 at the time of DT9. The
range depends on the mass of the target, as shown in Fig. 3.
As far as the coalescence of the binary neutron stars with thenominal mass of 1.4 M is concerned, the distance range does
not become larger by improving the sensitivity at lower frequency, such as 100 Hz. A sensitivity improvement at lower
frequencies is effective for the detection of coalescences involving heavier compact stars, like black holes. After DT9,
we stopped the observational operation of TAMA and entereda work phase of noise-hunting to attain the final sensitivity.
Since noise-hunting is meaningful only when the temporal
best sensitivity is obtained, almost all adjustments must be
done at quiet times in a sense of seismic vibration. The sensitivity in a frequency range of less than 300 Hz was limitedby the alignment control noise, which arose from the actuator force noise on the main mirror. For observation datataking runs, the feedback gain was set rather high to maintainstable operation against any irregular disturbance of the seismic motions. This gain needs to be decreased to reduce thisforce noise. Electric circuits demodulating the signal from thephotodiode were checked and revised. The oscillator noises
were checked and confirmed to be safe by a simulation of themodulation-demodulation procedure.
The most effort was applied to reduce scattering lightnoise. Almost all reflections were captured by light-absorbing
cells that were made of a highly light-absorbing material.
However, some reflections were not correctly captured due tolimited space for mounting absorbers. For example, the reflection point from the outer surface of the near mirror of the
Fabry-Perot cavity is quite close to the reflection from the inner surface of the near mirror, even at the position of the photodiode, which means that the absorber cell blocks the output
beam of the interferometer. If the wedge angle of the near mirror was designed much larger, this would be avoided. Among
many reflective surfaces inside the vacuum chamber aroundthe optical system of the interferometer, the reflection fromthe surface of the photodiode was regarded as the most harmful one because the photodiode is placed on an optical benchoutside the vacuum. The reflection light is heavily contaminated by the vibration of the photodiode due to seismic noise.
In order to estimate how much light-scattering noise90

2.7Msolar - 96.3kpc

2001/06/02 DT6
2003/02/20 DT8
2003/11/04 DT98070

1.4Msolar - 72.8kpc

60
50
0.5Msolar - 32.7kpc

40
30

10Msolar - 19.8kpc

20
10
0
10

-1

2

3

4 5 6 7 89

0

2

3

10
mass [Msolar]

4 5 6 7 89

10

1

2

Fig. 3. Detectable range in DT9 for each mass of a neutron star binary
system by a signal-to-noise (SN) ratio of 10. The coalescence ofthe nominal binary system (1.4 M of each mass) can be detected
up to 73 kpc.

Fig. 4. For noise-hunting, the central part of the interferometer(Michelson part) was inspected with appropriate stray-light absorbers consisting of boxes made of a highly reflective material.

30

igata Univ., Niigata)
Recycled Michelson ~ displacement noise level(2004/12/10)

-10

10

2004/01/15 Just after DT9
2004/01/29 PO removal, Dark WFS
2004/02/05 4PD->1PD, servo filter noise improv.
2004/02/12 PD modified
2004/04/23 Scattered light in the vacuum chamber
2004/12/10 PO reinstallation

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1/2

Displacement noise [m/Hz ]

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Overview

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1

2

10

3

10

10

10

4

10

5

Frequency [Hz]Fig. 5. Improvement of the sensitivity of the Michelson part without
Fabry-Perot cavities to find the ultimate limiting factor of the sensitivity.

Expected RMI phase noise level-7

10

Current status of LCGT Project

-8

10

-9

1/2

Displacement noise [m/Hz ]

10

-10

10

2004/12/10 RMI Phase noise
2003/11/04 RFPMI DT9

-11

10

-12

10

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10

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10

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10

-16

10

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10

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10

Although the recent discovery of the binary neutron starsystem [4] has increased the birth rate of the coalescence, it
is not easy to observe such gravitational wave events during
a short observation time. This is because the birth rate of the
coalescence of double neutron stars is 10 5  10 6 a year ina galaxy as big as our Galaxy [5]. This means that we need
to wait for long time to detect a gravitational wave from thecoalescence of binary neutron stars by the present detector.
This is the reason why we have to plan LCGT (Large-scale
Cryogenic Gravitational wave Telescope) [6]. There are manyother possible gravitational wave sources in the universe otherthan the coalescence of binary neutron stars. However, thecoalescence of binary neutron stars differs completely fromother sources in the sense that its wave form is precisely predicted, and its existence has certainly been confirmed.

0

10

10

1

2

10

10

3

10

4

5

10

Frequency [Hz]Fig. 6. If the present sensitivity limit arises only from the light scattering, the improvement of the Michelson part will result in increasing
the sensitivity.

is produced in the interferometer, the central part of theinterferometer (Michelson part) was checked. Absorbers
were mounted as shown in Fig. 4 to kill stray light beams.
Figure 5 shows the history of improving the sensitivity of theMichelson part without Fabry-Perot cavities. By using straylight-beam absorbers, one and a half orders of the noise magnitude were reduced. If the sensitivity limit arises only fromthis stray light, a one and half order improvement of the sensitivity is expected by restoring Fabry-Perot cavities (Fig. 6).

LCGT Project[Spokesperson : Kazuaki Kuroda]
ICRR, Univ. of Tokyo, Kashiwa, Chiba 277-8582
In collaboration with the members of: LCGT collaboration
NAOJ, Tokyo; KEK, Tsukuba; UEC, Tokyo; Osaka City
Univ., Osaka; Kyoto Univ., Kyoto; Osaka Univ., Osaka; Ni-

The target sensitivity of LCGT is to observe binary neutron star coalescence events occurring at 252 Mpc with
S/N=10 in its optimum configuration. This is ten-times
more sensitive than that of the LIGO (I), and by two orders more than that of TAMA at their most sensitive frequencies. This will be attained by dual interferometers locatedunderground, using a three-kilometer length baseline, cooling mirrors at cryogenic temperature, and a high-power laser
source employing 150 W output. The optical configuration isa power recycled Fabry-Perot-Michelson interferometer withthe resonant-sideband-extraction (RSE) scheme (in Fig. 7).
The detailed design of the control system is under consideration.
Table 2 lists the important parameters of LCGT, whichwere revised two times from the original design. The ultimate
sensitivity of a laser interferometer is determined by seismic
noise at low frequencies (10-30 Hz) (which is reduced by improving the vibrational isolation system), and it is limited byphoton shot noise at higher frequencies (more than 300 Hz),which can be improved only by increasing the light power inthe main cavities. The sensitivity of middle frequencies (30300 Hz) is limited by the photon recoil force noise. This requires that thermal noise is reduced both by decreasing thetemperature and by decreasing the internal mechanical loss
(i.e., increasing the mechanical Q of vibration modes). The
source of thermal noise comes from both mirror internal vibration, mechanical loss of the optical coating and swing noiseof the pendulum suspending the mirror. The reduction of thermal noise is attained by cooling down both the mirror, itself,
and the suspension system that suspends the mirror.
The main effort on the research and development for
LCGT has been placed on cryogenic mirrors for the past years.
The implementation of cryogenic mirrors is one of the moststraight-forward solutions to improve the sensitivity. A conceptual design of the cryogenic mirror system is shown in
Fig. 8, where only a part of the end-mirror chamber is focused.

31

Fig. 7. Optical design of LCGT. The optical configuration is a power recycled Fabry-Perot-Michelson interferometer with the resonant-sideband-extraction (RSE) scheme.

Table 2. LCGT design parameters to detect binary neutron-star coalescence events in 252 Mpc.

Item
Baseline LengthInterferometers

Optical PowerBeam radius at End
Main Mirror

Suspension
pendulum
VacuumParameter
3 km
Two setPower recycled Fabry-Perot-Michelson with RSE
Laser:150 W; Finesse:2000
Input power at BS: 750 W
Cavity power 770 kW
3 cm
Sapphire 30 kg, 20 K
Diameter 25 cm
Mechanical Q:108
Frequency: 1 Hz; Q:1  108
10 K
10 7 Pa

The mirror is suspended by one or two loops of sapphire fibers
connected to an alignment control platform that has a heatlink to the heat anchor (4 K) inside the vacuum located justabove the platform. This platform is suspended with an insulator rod connected through the center holes of the radiation
shields to an isolation table suspended by a low-frequency vibration isolator, which is placed at room temperature. Boththe cryogenic system and the vibration isolator are put inside
a common high-vacuum chamber.
To realize this concept, the following research subjects
were considered:

1. Removal of heat produced by high-power laser illumination.
2. Holding the high Qs of the mirror internal modes andsuspension pendulum.
3. Reducing the contamination of mirror surfaces.
4. Estimating heat production by optical loss in the mirror.
5. Alignment control of mirrors in a cryogenic environment
6. Low mechanical loss of the optical coating.
We had already reported on experimental results of the firsttwo items in the annual report (1997–1998) [7]. In regard to
the third and fourth items, the results were reported in the annual report (2000–2001) and published in papers [8]. As foritem 5, we confirmed that a superconducting film could be
used for the receptor of the magnetic force in place of permanent bar magnets that are normally used in the existing detectors. The film can be easily sputtered on the mirror surface
without harmfully degrading the mechanical Q of the mirror.
The basic behaviour of this method was reported in a paper[9].
With respect to the last item, we reported on a measurement of the bulk substrate of the mirror at cryogenic temperature in the annual report (2003–2004). We can now cor-

32
5#5QTUKOKNCTU[UVGO

105

RRR

104

103

102
0.1



Fig. 8. Schematic design of the cryogenic suspension system. The
mirror is suspended by sapphire fibers connected to a platform thathas a heat link to a 4 K heat anchor inside the vacuum. The platform is also suspended with an insulator rod connected throughthe holes of radiation shields to an isolation table suspended by a
low frequency anti-vibration system placed at room temperature in
the common high vacuum.

rectly estimate the thermal vibration noise of the optical coating while considering the inhomogeneous loss that had been
neglected at an early stage of interferometer development.
The substrate of the cryogenic mirror is sapphire, which has alarge thermo-elastic thermal noise at room temperature. However, since the thermal-expansion ratio of sapphire at cryogenic temperature goes down to nearly 0 and the heat conductivity becomes greater, the thermo-elastic noise is drastically reduced at the cryogenic temperature. Thermal noise estimated from the result of coating was well below the design
sensitivity of LCGT, which means that this coating noise does
not limit the sensitivity, whereas, the sensitivity of a roomtemperature mirror is limited by this effect. This is the significant merit of the cryogenic mirror system [10].
All of the above R&D confirmed the feasibility of reducing the thermal noise of the interferometer in the middlefrequency region. This research underlines the basis of LCGT.
However, for a practical cryogenic detector, much practicalR&D is needed for the installation of cryogenic mirrors. Oneof the earliest R&D activities was the Kashiwa cryogenic interferometer system reported in the annual report (2000–2001;
2002–2003; 2003–2004). By this Kashiwa cryogenic interferometer, we learned the necessity of several practical R&D
activities and began to construct the CLIO interferometer in
Kamioka to establish techniques for the cryogenic interferometer. Some of these practical studies are described here.

Diameter [mm]1Fig. 9. Size dependence of the thermal conductivity of pure metal
fibers. RRR(see text) represents thermal conductivity in this measurement. The solid line shows the calculated upper limit of aluminum (round plots), and the broken line corresponds to thatof copper (diamond plots). This result supports that the thermal conductivity of the wire is approximately proportional to its
cross-sectional area in these diameters.

Practical R&Ds: Dependence of the thermal conductivity on the thickness of fiber
The mirror is cooled by sapphire crystalline fibers thatsuspend a main mirror at the final stage of the anti-vibration
system. The thermal conductivity of the sapphire fiber wasmeasured (and reported in the previous annual report (2001–
2002)) using 250 μ m diameter sapphire fibers. Also, the thermal conductivity of the heat links was measured. The platformthat suspends the main mirror is supported through holes ofradiation shields from a room-temperature base, which meansthat some heat links are needed to connect the platform with athermal anchor inside the cryostat system. Since the platform
is close to the final stage of the anti-vibration system, suchmechanical links might produce mechanical disturbances onto
the main mirror. To prevent such a disturbing force, the mechanical rigidity of the heat links must be as low as possible
while maintaining high heat conductivity. This is realized by
dividing the heat link to many thin and long links. Imagine
that one thick fiber is divided into N identical thin fibers whileleaving the total cross-sectional area unchanged. Since the diameter is reduced in proportion to the square root of N, thetotal rigidity is reduced according to N1 , because each rigidity
changes in proportion to the fourth power of the diameter. In
this context, the thermal conductivity is supposed to be proportional to the cross-sectional area. However, the thermal
conductivity of sapphire fiber is not subject to this law, butchanges differently in proportion to the cube of the diameter,as reported in the previous annual report (2001–2002).
The design material of the heat links used in LCGT is purealuminum. Although copper is possible, the heavier mass density introduces a demerit of low-frequency resonances in theheat links. According to the measurement, the size dependence of the thermal conductivity was not found for diameters
ranging from 0.2 mm to 2 mm, as shown in Fig. 9, where RRR
is shown in place of the conductivity, because RRR represents

33
㱗/4

㱗/2

CMP

Nd:YAGPD
Polarizer
Sapphire rod (60mm)

Polarizer

Fig. 11. Optical phase retardance caused by some imperfection iscompensated by the Babinet-Soleil compensator (CMP), which ismeasured by determining two independent angles.

Practical R&Ds: Measurement of the optical qualities of sapphire

Fig. 10. Mirror suspension system adopts SPI (see text) to reject the
introduction of mechanical vibration through heat links. The auxiliary mirror (SPI mirror) behaves as a heat anchor for the main
mirror. The SPI mirror is suspended by an insulator wire froman upper platform placed in a room-temperature vacuum chamber
through holes in the radiation shields.

the thermal conductivity in this temperature region for metals.
RRR is defined as the ratio of the relative electrical resistanceρ
at 4.2 K and 300 K, ρ300K . This confirms that the design of
4 2K
LCGT is not affected by this size effect [11].
Practical R&Ds: Isolation from refrigerator vibration
Since the mechanical vibration of a refrigerator is large
enough to excite the suspension system of the main mirror
through heat links, both efforts to reduce the vibration of therefrigerator, itself, and that to isolate the vibration from entering to the mirror suspension system were taken. The latter was accomplished by the fine and soft heat-link material
described in the previous subsection. The development of aquiet refrigerator was conducted at a cryogenic science centerof KEK at Tsukuba. Some important techniques were submitted as patents. Along with these R&D efforts, we plan to
introduce a system of suspension-point-interferometer (SPI)
in order to eliminate vibration introduced through heat links,
which was originally proposed by R. Drever a long time ago,
and was recently tested by Y. Aso in the physics department of
the School of Science, the University of Tokyo. The main mirrors are suspended from points where auxiliary mirrors form
another Fabry-Perot cavity, which will be fixed in an accuracyof a light wavelength by a feedback loop. Even if a disturbing force is applied to one or both of the auxiliary mirrors,
the feedback system tries to make the cavity length constant,resulting in high rejection of the introduction of vibration. In
order to guarantee the design isolation, the SPI is arranged asshown in Fig. 10, including the main mirror suspension system.

The substrate of the main mirrors in the cryogenic interferometer (LCGT) is sapphire, because it has ultra-low mechanical loss and high heat conductivity at cryogenic temperature. Only fluorite is a possible substitute of sapphire, withsome difficulties that should be removed if we have to utilize
fluorite in place of sapphire. However, sapphire has birefringence, which is not a defect if the laser beam axis is aligned
along its c-axis, which is the so-called optical axis. Also, sapphire has a rather high optical scattering loss compared withthat of fused silica. A measurement of the birefringence ofsapphire was scheduled both for inspecting the crystal quality of sapphire substrates that were purchased from companies producing sapphire, and for estimating the optical quality of each substrate. The measurement principle is shown in
Fig. 11. Suppose that there is some imperfection of the crystal
that causes an optical phase retardance between two orthogonal polarizations. Through a half-wave plate we introduce
linearly polarized light that coincides with one of principle refractive axes. The retardance caused when traveling inside thesapphire substrate is compensated by a compensator (CMP),
the condition of which is measured by knowing an angle position of a polarizer to extinguish the output light. Without thesubstrate the angle of this polarizer is orthogonally set with
the input polarization angle. The fluctuation of the magnitudeof the retardance is represented as the indices of the magnitude of the crystal imperfections. In this measurement, for atypical point of the substrate, the angle of the polarizer andthe compensated angle were recorded. Practically, the polarizer angle was determined by knowing the optimum point.
Figure 12 shows an example of mapping the measured optical phase shift caused by the substrate, which should be zero
if there is no birefringence. A measurement of four sapphirerods was completed, and we confirmed that there were no apparent defects inside, and that the magnitudes of the imperfection were negligibly small in the sense of the fluctuatedbirefringence.
In order to measure the optical scattering loss of thosesample rods, we brought two sapphire rods to Australia, andmade measurements in collaboration with a gravitational wave
group in the Physics department, University of Western Australia, where a measuring device of Rayleigh scattering had
been developed. We successfully measured the scattering lossof the samples. Figure 13 shows a strong scattering region on
the substrate, the origin of which is believed to have resulted
from a remnant of the polishing powder (carbon particle). By
this knowledge, we could request the procedure of manufac-

34
30
0.26

Y[mm]

20

0.24

10

0.22

0

0.2
0.18

−10

0.16
−20
0.14
−30
−30

−20

−10

0
10
X[mm]

20

30

Fig. 12. Measured birefringence fluctuation in terms of the retarded
angle. The fluctuation of the magnitude of the retardance represents the magnitude of the crystal imperfections, which is negligibly small for the mirror substrate. Note that a lighter color showsa worse retardance as in the right-hand side column. The unit ofnumbers is radians.

30
25

20

Y[mm]

10

20

0
15

−10
−20
−30
−30

10
−20

−10

0
10
X[mm]

20

30

Fig. 13. Mapping of the measured scattering amplitude in terms of
power loss. Note that lighter color shows stronger loss, as in theright-hand side column. The unit of numbers is ppm/cm.

turing sapphire.
Apart from the above R&D issues, researchers at the
National Astronomical Observatory developed control techniques for the RSE scheme. Also, researchers at KEK tested
the mechanical strength of bonded pieces made of sapphire,
besides quiet refrigerators. We could not find any correlation
between the birefringence and the scattering light loss.
We steadily advance towards the realization of LCGT bythese R&D activities. The next step in R&D for cryogenic
mirrors is to show a reduction of the noise amplitude, itself,
which needs a more realistic interferometer other than the prototype with a shorter baseline. For this sake, the CLIO interferometer at Kamioka is being constructed.

Construction of CLIO at Kamioka[Spokesperson : Kazuaki Kuroda]
ICRR, Univ. of Tokyo, Kashiwa, Chiba 277-8582
In collaboration with members of: KEK, Tsukuba; KyotoU, Kyoto; Physics Department of UT, Tokyo; ERI of UT,
Tokyo

CLIO is a project to construct a 100 m-baseline underground cryogenic interferometer at Kamioka. A tunnel housing the vacuum tubes with chambers has been constructed.
The CLIO forms a bridge connecting the CLIK(7 m prototype
cryogenic interferometer at Kashiwa campus) and the planed
LCGT(3 km cryogenic interferometer at Kamioka).
The objective of CLIO is to present the validity ofcryogenic interferometers in an intermediate baseline scale(100 m). We focus attention on the decrease of thermal noiseby lowering the temperature by this practical interferometersystem. For this reason, the interferometer optical system isdesigned to be as simple as possible in order to make sophisticated control techniques unnecessary during operation. This
involves a locked Fabry-Perot configuration with ring modecleaners (Fig. 14). The main mirrors are cooled at 20 K by refrigerators. The lowest noise level of CLIO is designed to be
10 19m Hz around 100 Hz, which would be 10 18m Hz,
which results from thermoelastic noise of sapphire mirrors if
cryogenics is not applied [12].
Once the objective is attained, the CLIO interferometeris used to observe gravitational wave events in parallel with
the TAMA interferometer until completion of the constructionof LCGT. The merits of the underground site are lower seismic noise and temperature stability. The former characteristic
makes interferometer locking easily controlled, and the latterassures long-term stable operation [13]. The site of CLIO,
near the Super-Kamiokande neutrino detector, is shown in
Fig. 15. The tunnel was dug in 2002, and a strain meter for
geophysics was installed in 2003. The construction of CLIO
began in late 2003, and installation of the mode cleaner vacuum system was reported in the annual report (2003–2004).
By early this year, three sets of cryostats had been installed,
and its cryogenic test was finished for one cryostat (Fig. 16)
and one-arm vacuum system with cryostats of both ends had
been constructed (vacuum leak test was almost finished). Vacuum chambers for the input optics in the center room were
fully installed.

Bibliography[1] J. H. Taylor and J. M. Weisberg, Astrophysical J., 345(1989) 434.
[2] N. Kanda, et al., Class. Quantum Grav. 20 (2003) S761.
[3] K. Soida, et al., Class. Quantum Grav. 20 (2003) S645.
[4] M. Burgay, et al., Nature, 426 (2003) 531.
[5] C. Kim, V. Kalogera and D. R. Lorimer, AstrophysicalJ. 584 (2003) 985.
[6] K. Kuroda, et.al., Int. J. Mod. Phys. D 8 (1999) 557;
K. Kuroda,et al., Class. Quantum Grav. 19 (2002) 1237;
T. Uchiyama, et al., Class. Quantum Grav. 21 (2004)
S1161.
[7] T. Uchiyama, et.al., Phys. Lett. A 242 (1998) 211; T.
Uchiyama, et.al., Phys. Lett. A 273 (2000) 310.
[8] S. Miyoki, et al., Cryogenics 40 (2000) 61: S. Miyoki, etal., Cryogenics 41 (2001) 415: T. Tomaru, et al., Phys.
Lett. A 283 (2001) 80.

35

Input Optics
Two Mode Cleaners(Mod Transmission Control is adopted for 2ndMC)

Gravitational
Wave
SignalFeed
Back

Main Interferometer
EM2

Locked Fabry-Perot

Feed
Back

Perpendicular
Arm Cavity

2nd MC
Length:9.5m
Finesse:1500

OptionFeed Feed
Back Back

Sapphire Mirror
Temp:20K

MMTFP

Innolight
Laser
1064nm

MMTMC

Inline Arm Cavity
Length:100m
Finesse:7500

R= - 43m
R= 30m2W
fm1
MC1
Signalfm2 fmFP
Mod.Unit
12M, 15.8M

Fused BS
Silica Beam
SplitterTemp:300K

FM1

EM1

MC2
SignalPD

Faraday Isolator

Oscillator

λ/4,λ/2Plate

4 divided
Photo Detector

Mode Matching Mirror

EOM

PBS

Mixer

Photo Detector

Sapphire Mirror

FeedAroundand MC End

MirrorFig. 14. Optical design of the CLIO interferometer. This is a locked Fabry-Perot configuration with ring mode cleaners.

Fig. 15. Location of the CLIO interferometer.

[9] N. Sato, et al., Cryogenics 43 (2003) 425.
[10] K. Yamamoto, et al., Class. Quantum Grav. 21 (2004)
S1075.
[11] K. Kasahara, et al., TEION KOHGAKU 39 (2004) 25(in Japanese).
[12] M. Ohashi, et al., Class. Quantum Grav. 20 (2003) S599.
[13] S. Sato, et al., Phys. Rev. D 69 (2004) 102005.

Fig. 16. This cryostat was installed in 2004 and its cryogenic test has
been successfully finished. One-arm of the vacuum system wasconstructed and vacuum chambers for the input optics were fully
installed.

36

Sloan Digital Sky Survey[Spokesperson : Michael Strauss]
Princeton Univ., Princeton, New Jersey 08544 USA
In collaboration with the members of:
University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 1130033, Japan; Nagoya University, Chikusa, Nagoya 464-8602,
Japan; National Astronomical Observatory of Japan, 2-211 Osawa, Mitaka, Tokyo 181-8588, Japan; Tohoku University, Aramaki, Aoba, Sendai 980-8578, Japan; Japan Women’s
University, 2-8-1, Mejirodai, Bunkyo-ku, Tokyo, 112-8681,
Japan; The University of Chicago, 5640, South Ellis Ave.,
Chicago, IL 60637, USA; Fermi National Accelerator Laboratory, P.O. Box 500, Batavia, IL 60510, USA; Institute forAdvanced Study, Einstein Drive, Princeton, NJ 08540, USA;
Johns Hopkins University, Baltimore, MD 21218, USA; Los
Alamos National Laboratory, Los Alamos, NM 87545, USA;
Max-Planck-Institute for Astronomy, Königstuhl 17, D-69117
Heidelberg, Germany; Max-Planck-Institute for Astrophysics,
Karl Schwarzschildstrasse 1, D-85748 Garching, Germany;
New Mexico State University, P.O. Box 30001, Dept 4500,
Las Cruces, NM 88003, USA; University of Pittsburgh, 3941
O’Hara St., Pittsburgh, PA 15260, USA; Princeton University, Princeton, NJ 08544, USA; United States Naval Observatory, P.O. Box 1149, Flagstaff, AZ 86002-1149; Universityof Washington, Box 351580, Seattle, WA 98195, USA
Four years of operation were completed since the commencement of the production run early in the year 2000. The
initial goal of the Sloan Digital Sky Survey (SDSS) was tocover 10000 square degrees of sky for both imaging and spectroscopic observations, but at the time the production survey
began, the Five Year Baseline was developed to provide amore realistic metric against which the progress was evaluated. This decreased the target area to 8550 square degrees(7800 square degrees in the northern sky and 750 square degrees in the southern sky). The imaging carried out over
4 years has sustained nearly the planned pace. At the end of
the fourth year 84 % of the northern-sky baseline and 100%of southern-sky baseline were completed. Spectroscopic surveys, however, are significantly behind the schedule. Only
40% of the baseline area were observed for the north (70%for the south). This is due to unusually poor weather conditions in the spring of 2003 and in the winter of 2004. The
time-consuming spectroscopic runs were severely affected.
(The overall fraction of time available for observing in January 2004, for instance, was only 16%, compared with thebaseline value of 60%.) Otherwise, all the operations havebeen working smoothly (98% is the uptime fraction), producing the data as expected. Software has undergone extensive
fine tunings and all the data were re-reduced, which are now
made public as The Second Data Release of the Sloan Digital
Sky Survey, comprising 3300 square degrees of imaging and2600 square degrees of spectroscopic surveys done to 1 July
2002. This data base contains 88 million objects including
370000 objects with spectroscopic information.
The prime scientific goals of the SDSS are focused on
extragalactic themes, including the large-scale structure ofgalaxies over very large volume of the Universe, and detailed

characterisations of the galaxy properties and those of quasars.
As the survey proceeds, these goals have gradually been satisfied. The clustering of galaxies mapped in the three dimensional space looks very similar to what is expected in themodel of the Universe dominated by cold dark matter withdensity fluctuations starting from nearly scale-invariant adiabatic perturbations, as predicted in the model of inflation.
As a quantitative measure this density field is characterisedby the statistic called the power spectrum, the squared amplitude of the Fourier modes of fluctuations as a function ofscale. One of the important results is an accurate derivation of this power spectrum from galaxy clustering over the
scale from 10 to 200 Mpc, and the demonstration that it joinssmoothly the spectrum derived from the temperature field
imprinted on the cosmic microwave background measured
by Wilkinson Microwave Anisotropy Probe (WMAP) (seeFigure 1). Combined with WMAP that explores the Universe at z 1000, this lends the most convincing support tothe standard model of structure formation in the Universe.
The combined data of WMAP and SDSS yield the cosmological parameters Ωm  0 30 0 04, ΩΛ  0 70 0 04 andH0  70 4 km s 1 Mpc 1.
Given a few hundred thousands of galaxies with accurate
multicolour photometry and spectroscopic information, manystudies have been conducted to establish statistical propertiesof galaxies. To mention a typical result, an accurate determination was made on the average stellar mass (together withthe average age and the heavy element abundance) containedin galaxies. Combined with the accurate luminosity functionof galaxies, obtained also from the SDSS, this tells us thatbaryons in stars are 6 1% of the total. Another direction
is studies of the properties of galaxies in more detail, taking the advantage of accurate photometry and spectroscopy.
One such project conducts visual classification of galaxies
into Hubble types. For this is a laborious process, the initialsample contains only 2000 nearby galaxies, notwithstanding
this is the largest homogeneous sample of morphologicallyclassified galaxies ever produced. This ‘small’ sample alreadyclarified a number of aspects that were not known or confusing. One example is a discovery of actively ‘star forming
elliptical galaxies’ (Figure 2), which contrasts with the conventional wisdom that elliptical galaxies are those that havelong lost star formation activity, consisting only of old stars.
We find the fraction of such elliptical galaxies on the orderof 0.1%, hardly discernible in samples available before theSDSS.
Gravitational lensing is also a subject to which the SDSS
is making a significant contribution in a number of ways. The
SDSS is one of the first few that observed a weak effect ofgravitational lensing of galaxy images by foreground galaxies. The lensing effect appears as distorted images of galaxies
due to the gravitational shear field, but it is only on the order of a few percent, compared with the order of unity effectof randomly oriented galaxies that have intrinsically different
shapes. Millions of galaxies are needed to extract this small
signal from noise of the order of unity. With the SDSS galaxy
sample, this distortion was unambiguously detected, showing that the mass concentration around galaxies behaves as
r 0 8 as a function of distance r, consistently with the famous

37
1 8 law of the galaxy-galaxy correlation.

r
The uniquenessof the SDSS observation over similar weak lensing projects
rests in the fact that the distances to foreground galaxies are
known, which in turn allows the estimate of the dark mass associated with galaxies. The analysis showed that galaxies are
surrounded by the dark mass which amounts to 200 times thelight in solar units. This large mass to light ratio had been inferred from a few observations in the past under some assumptions, but this time is derived without resorting to any assumptions. The large quasar sample of the SDSS also provides an
excellent platform to search for classical strong gravitational
lenses of quasar images, i.e., splitting of images. Amongthe novel cases that deserve special scientific discussion, wequote a lens of four images with their maximum separation
being 14.4 arcsec, 2.5 times larger than has ever been found(Figure 3). The creation of statistical sample of strong lenses
is eagerly awaited, since it will provide a new tool to study
dynamical states of galaxies and their evolution.
From the commissioning phase, the SDSS project has
made unrivalled contributions to our understanding of highredshift quasars. By now the project has found 6 quasars with
redshift higher than 6 out of 4600 square degrees — There are
no z  6 quasars reported from other existing projects. The
highest redshift is 6.42, which means that the light was emitted only 0.84 Gyr after the Big Bang. Our finding shows theabundance of luminous quasars declining exponentially with
redshift from z 3. This sets a limit on the model of quasars
at high redshifts. Also interesting with high redshift quasars isthat they allow a study of the ionisation state of the intergalactic medium (IGM). We know that free electron and proton recombine and neutral hydrogen is formed at z 1500, but it
was again highly reionised before z 3. When and how thisreionisation took place is a matter of significant interest from aview of galaxy formation. The spectrum shows a flux of highredshift quasars shortwards the Lyman α line nearly vanishingfor z 6 (called Gunn-Peterson trough), indicating the changeof states of the IGM at this redshift. This contrasts with thereionisation epoch of z 17 inferred from WMAP. How thesefindings will be reconciled will be followed with interest. Besides high redshift quasars, a large number of quasars are assembled and their properties are being studied, including thesearch for gravitational lenses and for quasars with peculiarfeatures. Our second catalogue of quasars, containing 17000
quasars in 1400 square degrees was published.
After data taking started, it has been recognised that theexcellent quality of multicolour photometry and enormousvolume of the data produced by SDSS are also useful to makeprogress for understanding the Galactic structure and the low
mass end of stars. The most interesting finding for the former
perhaps is the discovery of tidal debris which were left behindwhen dwarf galaxies travel through the Galactic halo. The firstevidence was a conspicuous enhancement in the number density of stars in a small area when it was studied as a function ofradius. It was soon identified as the tidal leftover of the Sagittarius dwarf galaxy orbiting around the Galaxy. Another tidal
track was then discovered for disruption of a globular cluster,
Palomar 5. The tidal tail contains a comparable mass as does
the cluster. The third case is the Monoceros stream. These
pieces of evidence indicate a view that the Galaxy is still be-

Fig. 1. The power spectrum of galaxy clustering scaled to the present
epoch (SDSS galaxies) plotted with four other independent measures, extending four decades on spatial resolution, and demonstrating remarkable consistency. These data provide fundamental
constraints on cosmological models. Figure is taken from Tegmark
et al. Phys. Rev. D69, 103501 (2004).

ing built up in the halo. The perturbations on the tidal tracks
tell us that the Galactic halo is appreciably non-spherical andlumpy. The other notable contribution from the SDSS is to thestellar science proper. A substantial number of low temperature stars found are on the extension of the M dwarf, named
the L and T dwarfs, the latter dominated by methane and water features in their spectrum just as with Jupiter. Despite thediscontinuous behaviour in colour space it was clarified thatthese stars are on a single temperature sequence. The sample
allows detailed studies of low temperature stellar atmosphere.
The mass contained in these low mass stars is no more than afew percent of the total star mass: therefore they cannot be adominant component of the Galaxy mass.
To date 400 scientific papers have been written for regular journals and conference proceedings (about 150 include
Japanese team members). The number of papers using the
SDSS data by authors outside the SDSS is also rapidly increasing. Nevertheless, we feel that the state of analyses isstill quite premature, far from exploitation of the data we produced.

38

obtained, in particular, on the chaotic inflation with higherpower monomials, the natural inflation, and the new inflation.
We study how the curvaton scenario changes (and relaxes) the
constraints on these models.

Bibliography[1] T. Moroi, T. Takahashi and Y. Toyoda, hep-ph/0501007.

Spacetime Symmetries on the Supermanifolds[Spokesperson : J. Harada]
ICRR, Univ. of Tokyo, Kashiwa, Chiba 277-8582
Fig. 2. Images and spectra of the three actively star-forming ellipticalgalaxies found in the SDSS. The fourth panel represents a bonafide elliptical galaxy. The first three show conspicuous emission
lines as those observed in late spiral or irregular galaxies, where
active star formation takes place. Figure is taken from Fukugita et
al. Astrophys. J. (Letters) 601, L127 (2004).

Harada investigated the spacetime symmetries on the supersymmetric supermanifolds as the extension of Lorentz andPoincaré invariance. It was found that the new rotational invariance on supermanifolds may be the possible generalization of Lorentz and Poincaré invariance. The new symmetry generator is the Lorentz vector, and it induces the coordinate transformation on not the ordinary spacetime coordinates
x but on the anticommuting Majorana spinor coordinates θ .
It was also found that the new rotational invariance and theLorentz invariance are naturally unified. This unified symmetry may be the correct spacetime symmetry in the very early
universe.

Bibliography[1] J. Harada, in preparation.

EWIMP Dark Matter Detections[Spokesperson : Sh. Matsumoto]
ICRR, Univ. of Tokyo, Kashiwa, Chiba 277-8582
In collaboration with the members of ICRR and YITP, Kyoto
Univ..
Fig. 3. Subaru telescope follow-up of the quadruple lens discovered
with the SDSS. The four images of the quasar are labelled A, B,
C and D. A rich cluster of galaxies is evident, centred in the field.
The dominant central galaxy of the cluster, which causes splittingof images of the quasar is labelled G. The separation between
C and D is 14.43 arcsec, the largest among lenses ever found.
Figure is after Inada et al. Nature 426, 810 (2003).

Theory Group
Implications of the Curvaton on Inflationary
Cosmology[Spokesperson : T. Takahashi]
ICRR, Univ. of Tokyo, Kashiwa, Chiba 277-8582
In collaboration with the members of ICRR and Tohoku Univ..
We consider the effects of the curvaton, late-decaying
scalar condensation, to observational constraints on inflation
models. From current observations of cosmic density fluctuations, severe constraints on some class of inflation models are

Hisano, Matsumoto and Saito investigated some detection
signals for the dark matter having electroweak SU(2)L chargeof the Standard Model (EWIMP). Such dark matter is realized in the minimal supersymmetric standard model (MSSM)
when the lightest SUSY particle is the Higgsino- or Winolike neutralino. In indirect detections, the calculation of theannihilation cross section of the dark matter is important forpredicting the signals. They found that the non-perturbative
effects significantly contribute to the cross sections due to thethreshold singularity when the dark matter is heavy comparedto the weak gauge bosons. As a result, the cross sections (andtherefore the detection signals such as gamma ray flux) may
be enhanced by several orders of magnitude compared to leading calculations. In direct detections, the dark matter-nucleon
scattering is an important process. They discussed one-loop
contributions to the scattering cross section for the EWIMP
dark matter. While the tree level contribution is suppressedby new physics scale such as the dark matter mass, the oneloop contribution generated by the gauge interaction is not
suppressed. It may significantly alter the total cross section
when the cross section is 10 45 cm2 or less.

39

Bibliography[1] J. Hisano, S. Matsumoto, M. M. Nojiri and O. Saito,
hep-ph/0412403 (to be published in Phys. Rev. D).
[2] J. Hisano, S. Matsumoto, M. M. Nojiri and O. Saito,
Phys. Rev. D 71 (2005) 015007.

Significant Effects of the Second KK Particles on LKP Dark Matter Physics[Spokesperson : M. Senami]
ICRR, Univ. of Tokyo, Kashiwa, Chiba 277-8582
In collaboration with the members of ICRR and Saitama
Univ..
Kakizaki, Matsumoto, Sato and Senami investigatedKaluza-Klein (KK) dark matter physics is drastically affectedby the second KK particles. Various intriguing phenomena
caused by the second KK modes were discussed. In particular, we reevaluated the annihilation cross section and thermal
relic density of the KK dark matter quantitatively in universal
extra dimensions, in which all the standard model particles
propagate. We demonstrate that the KK dark matter annihilation cross section can be enhanced, compared with the treelevel cross section mediated only by first KK particles. The
dark matter mass consistent with the WMAP observation isincreased.

Bibliography[1] M. Kakizaki, Sh. Matsumoto, Y. Sato and M. Senami,
hep-ph/0502059.

Hadronic Electric Dipole Moments in Supersymmetric Grand Unified Theories[Spokesperson : M. Kakizaki]
ICRR, Univ. of Tokyo, Kashiwa, Chiba 277-8582
In collaboration with the members of ICRR and Tohoku Univ..
After the discovery of successful agreement of the three
extrapolated gauge couplings at a higher energy scale, supersymmetric grand unified theories (SUSY GUTs) have been
considered as ones of the most excellent candidates for atheory beyond the standard model. Meanwhile, it was recently pointed out that null results of hadronic electric dipole
moment (EDM) experiments severely constrain CP-violatingsquark generation mixings in SUSY models. Bearing this situation in mind, we investigated hadronic EDM constraints on
SUSY GUT models which predict sizable sfermion mixings.
In the SUSY SU(5) GUT with the right-handed neutrinos,
the hadronic EDMs are sensitive to the right-handed downtype squark mixings, especially between the second and third
generations and between the first and third ones, comparedwith the other low-energy hadronic observables, and the flavor mixings are induced by the neutrino Yukawa interaction.
We find that the current experimental bound of the neutron
EDM may imply that the right-handed tau neutrino mass issmaller than about 1014 GeV given the minimal supergravity boundary condition. In orbifold GUT models where the

GUT symmetry and supersymmetry are broken by boundary conditions in extra spatial dimensions, the appropriate
choice of a field configuration yields rich flavor structure in
the fermion and sfermion sector. We point out that a marginal
chromoelectric dipole moment of the up quark is induced by
the misalignment between the CP violating left- and righthanded up-type squark mixings together with the large top
quark mass in contrast to the conventional four-dimensional
SUSY GUTs. The sensitivity reach of hadronic EDMs is expected to be improved furthermore in future experiments, suchas the deuteron EDM measurement. The interplay between future EDM and lepton flavor violation searches will probe thestructure of GUT and SUSY breaking mediation mechanism.

Bibliography[1] J. Hisano, M. Kakizaki, M. Nagai and Y. Shimizu, Phys.
Lett. B 604 (2004) 216.
[2] J. Hisano, M. Kakizaki and M. Nagai, in preparation.

Big-Bang Nucleosynthesis and Hadronic
Decay of Long-Lived Massive Particles[Spokesperson : M. Kawasaki]
ICRR, Univ. of Tokyo, Kashiwa, Chiba 277-8582
In collaboration with the members of ICRR, Osaka Univ. andTohoku Univ..
We study the big-bang nucleosynthesis (BBN) with thelong-lived exotic particle, called X. If the lifetime of X islonger than  0 1 sec, its decay may cause non-thermal nuclear reactions during or after the BBN, altering the predictions of the standard BBN scenario. We pay particular attention to its hadronic decay modes and calculate the primordial
abundances of the light elements. Using the result, we derive
constraints on the primordial abundance of X.
Compared to the previous studies, we have improved thefollowing points in our analysis: The JETSET 7.4 Monte
Carlo event generator is used to calculate the spectrum of
hadrons produced by the decay of X; The evolution of thehadronic shower is studied taking account of the details of
the energy-loss processes of the nuclei in the thermal bath;
We have used the most recent observational constraints on theprimordial abundances of the light elements; In order to estimate the uncertainties, we have performed the Monte Carlosimulation which includes the experimental errors of the crosssections and transfered energies.
We will see that the non-thermal productions of D, 3 He,
4 He and 6 Li provide stringent upper bounds on the primordial abundance of late-decaying particle, in particular when
the hadronic branching ratio of X is sizable. We apply our results to the gravitino problem, and obtain upper bound on thereheating temperature after inflation.

Bibliography[1] M. Kawasaki, K. Kohri and T. Moroi
Physical Review D in press.

40

OBSERVATORIES andA RESEARCH CENTER
Location of the Institute and the Observatories in Japan

Norikura Observatory
Location:g
Telephone (Fax):
Telephone (satellite):Telephone (car):Nyukawa-mura, Ohno-gun, Gifu Prefecture 506-2100
N 36Æ 06 , E 137Æ33 , 2770 m a.s.l.
+263-33-7456
090-7721-5674
090-7408-6224

Akeno Observatory
Location:
Telephone / Fax:

Akeno-mura, Kitakoma-gun, Yamanashi Prefecture 407-0201
N 35Æ 47 , E 138Æ30 , 900 m a.s.l.
+551-25-2301 / +551-25-2303

Kamioka Observatory
Location:
Telephone / Fax:

456 Higashi-mozumi, Kamioka-cho, Hida-shi, Gifu Prefecture 506-1205
N 36Æ 25 26 , E 137Æ19 11 , 357.5 m a.s.l.
+578-5-2116 / +578-5-2121

Research Center for Cosmic Neutrinos
Location:
Telephone / Fax:

5-1-5 Kashiwanoha, Kashiwa, Chiba Prefecture 277-8582+4-7136-3138 / +4-7136-3115

41

NORIKURA OBSERVATORY
Norikura Observatory (36.10ÆN and 137.55ÆE) was
founded in 1953 and attached to ICRR in 1976. It is locatedat 2770 m above sea level, and is the highest altitude mannedlaboratory in Japan maintained all the year. Experimental facilities of the laboratory are made available to all the qualified scientists in the field of cosmic ray research and associated subjects. The AC electric power is generated by the dynamo and supplied throughout the observatory. In 1996, two
dynamos of 70 KVA each were replaced with the new ones.
The observatory can be accessed easily by the car in summer(June–October) but an aid of snowmobile is necessary in winter time.

accelerated at least up to 50 GeV in a large solar flare considering the geomagnetic cut-off rigidity of the proton is estimated to be 11.5 GeV at Mt. Norikura. The high altitudeof the observatory is essential for detecting the flare particles
without significant attenuation.
In 1990, Nagoya group constructed a solar neutron telescope consisting of scintillators and lead plates, which measures the kinetic energies of incoming neutrons up to several hundred MeV. This telescope observed high energy neutrons associated with a large flare occurred on the 4th of June,
1991. The same event was simultaneously detected by theneutron monitor and the high counting meson telescope of the
Norikura observatory. This is the most clear observation ofsolar neutrons at the ground level in almost ten years since thefirst observation at Jungfraujoch in 1982.
A new type of large solar neutron telescope (64 m2 sensitive area) was constructed by Nagoya group in 1996. It
consists of scintillators, proportional counters and wood absorbers piled up alternately. This takes a pivotal role among aworldwide network of ground based solar neutron telescopesof the same type in Yangbajing in Tibet, Aragatz in Armenia,
Gornergrat in Switzerland, Chacaltaya in Bolivia and Mauna
Kea in Hawaii. The Sun is being watched for 24 hours using
this network.

Fig. 1. Norikura Observatory.

Present academic interests of the laboratory is focused on
the modulation of high energy cosmic rays in the interplanetary space associated with the solar activity and the generationof energetic particles by the solar flares.
For the modulation study, two small experiments havebeen operated continuously for a long time. One is a neutron
monitor operated to study the correlation of solar activity andthe cosmic ray flux. The other is a high counting meson telescope consisting of 36 m2 scintillation counters to study thetime variation of cosmic rays with energies of 10–100 TeV.
The Sun is the nearest site to the Earth capable of accelerating particles up to high energies. When the Sun becomes
active, flares are frequently observed on its surface. The flareaccelerates the proton and ion to high energy and they are detected on the Earth soon after the flare. Among the particles
generated by the flare, high energy neutrons provide the mostdirect information about the acceleration mechanism as theycome straight from the flare position to the Earth without being affected by the magnetic field.
The detection of solar flare particles has been very activein the Norikura observatory for the last 10 years. The muon
telescope of the observatory successfully detected flare particles in association with a large flare occurred on the 29th of
September, 1989. The data suggested that protons could be

Fig. 2. New Solar-Neutron Telescope of Nagoya Group.

The Sun is reaching the maximum activity in 2001 andthe active phase will continue for next few years. All the telescopes in the Norikura observatory, neutron, meson and muon
telescopes, will be operated continuously through this solar
cycle (Cycle 23) in order to obtain comprehensive informationon the solar flare phenomena. Important hints for understanding the mechanism of cosmic ray acceleration will be obtainedby this measurement.
In addition to the long-term cosmic-ray observations mentioned above, various kinds of short-dated experiments arecarried out every year taking an advantage of the high altitudeof the observatory. As a few examples, following experiments

42

have been performed; a search for super heavy particles withplastic plates, a precise measurement of atmospheric gamma
rays and muons, collection of cosmic dusts contained in thesnow and the performance study of the balloon borne cosmic
ray experiments. A part of the facility has been open for theenvironmental study at high altitude such as the aerosol removal mechanism in the atmosphere.
The 50th anniversary of the Norikura Observatory wascelebrated in 2003. The feasibility of the automatic operationof the Norikura Observatory during winter period has beentested since winter 2004.

43

AKENO OBSERVATORYThe Akeno Observatory has evolved into a complex scientific center for studying ultra-high-energy cosmic rays throughextensive air showers in the energy range 1015 eV  1020 eV,
and is being used by many universities. The observatory is in
Akeno-chou, Hokuto-shi, located about 20 km west of Kofu
and about 130 km west of central Tokyo. Its altitude is 900 m
above sea level, and the location is at longitude 138.5ÆE andlatitude 35.5ÆN. An aerial view of the 1 km2 array area isshown in Fig. 1.
The following seven experiments were performed in
Akeno Observatory.
a) Observation of Ultra-High Energy Cosmic Rays (N.
Hayashida, ICRR)b) Prototype Test of Telescope Array (TA) Detectors (M.
Fukushima, ICRR)c) Atmospheric Monitoring Experiments with LIDAR
System (M. Chikawa, Kinki Univ.)d) Observation of Galactic Cosmic Ray Intensity with
Large Area Muon Telescopes (S. Kawakami, Osaka
City Univ.)e) Observation Test for High Energy Cosmic Ray Using Very Wide Field Refractive Optics (S. Ebisuzaki,
Riken)
f) Multicolor Imaging Telescopes for Survey and Monstrous Explosions (N. Kawai, Tokyo Inst. Tech.)g) Wide Field Telescope for Hunting Gamma-Ray Burst
Optical Flashes (T. Tamagawa, Riken)

Fig. 1. Akeno central campus and 1 km array area.

C: Central laboratory

D: Dormitory

S: Sub-electronics stations

M: Muon detectors (1 GeV)

ME: Muon detectors (0.5 GeV)

44

Fig. 2. TA surface detector prototype

Fig. 5. Large area muon telescope

Fig. 3. TA fluorescence detector prototype

Fig. 6. Telescopes for gamma-ray burst observation

Fig. 4. LIDAR system for atmospheric monitoring experiment

45

KAMIOKA OBSERVATORYThe observatory operates a 50,000 ton water Cherenkov
detector, Super-Kamiokande,which is located 1000 m underground in Kamioka Mine. The purpose of Super-Kamiokande
is to study elementary particle physics and astrophysics
through neutrino detection and nucleon decay searches.
Super-Kamiokande discovered evidence for neutrino oscillations using atmospheric neutrinos in 1998. Solar neutrino oscillation was established in 2001 by comparing results fromthe SNO experiment in Canada. In 2002, neutrino oscillation
was confirmed using artificial neutrinos produced by a proton
accelerator at KEK.
There are also 100 m long laser interferometers in
Kamioka Mine which are aiming to study gravitational waves
and geophysics. Using the low background environment in
Kamioka Mine, dark matter search experiments are also being constructed.
There are research offices, a computer facility and a dormitory for researchers located outside of Kamioka Mine andeasy access to the detectors in the mine.

Fig. 2. Super-Kamiokande detector.

Fig. 1. Map of Kamioka observatory.

Fig. 3. 100 m baseline laser interferometers for gravitational wave andgeophysics in Kamioka mine.

46

Fig. 4. Mt. Ikeno-yama in which Kamioka Mine is there.

Fig. 5. Research offices and computer facility.

Fig. 6. Dormitory for researchers.

47

RESEARCH CENTER FOR COSMIC NEUTRINOS
Research Center for Cosmic Neutrinos was established in
April 1999. The main objective of this center is to study neutrinos based on data from various observations and experiments.
In order to promote the studies of neutrino physics, it isimportant to provide the opportunity for discussions on theoretical ideas and experimental results on neutrino physics.
Therefore, one of the most important, practical jobs of thiscenter is the organization of neutrino related meetings. In
the fiscal year 2004, we organized one international workshop
and two local neutrino meetings. The topic of the international workshop was on the sub-dominant oscillation effects
in atmospheric neutrino experiments. About 30 researchers
working in neutrino oscillation phenomenology, neutrino experiments, cosmic ray measurements, atmospheric neutrino
flux calculation and neutrino interaction discussed details on
the possible observation of the sub-dominant oscillation effect, such as the effect of the solar neutrino oscillations on
atmospheric neutrino experiments.
In the local meetings, we have discussed topics such as, atmospheric neutrinos, solar neutrinos, accelerator neutrino experiments, reactor neutrino experiments, high energy neutrino
astronomy, models of neutrino mass, and unified theories. In
each meeting, about 40 researchers participated.

Atmospheric neutrino data from Super-Kamiokande givethe most precise information on neutrino oscillations. With
the increased data, it is more important to have better predictions of the neutrino flux. Therefore, we work on the prediction of the atmospheric neutrino flux. In order to predict
the flux accurately, it is important to know the details of thedata on the measurements of primary and secondary cosmic
ray fluxes. For this reason, we have a close collaboration withresearchers working in the cosmic ray flux measurements.
We worked for the basic study of the possibility of theintermediate detector on J-PARC long baseline neutrino experiment. This study is also related to the future 1 Mton water Cherenkov detector (Hyper-Kamiokande), which shouldbe able to measure various important quantities such as leptonic CP violation, proton decay and supernova neutrinos.
It is important that the general public knows the achievements of the present science. Because of this reason, we had apublic lecture at Kashiwa. Two active scientists related to neutrino physics lectured on various aspects of neutrino physics.
Finally, we mention that scientific staffs in this center are
actively working in the Super-Kamiokande and K2K experiments.

Fig. 1. RCCN International Workshop on sub-dominant oscillation effects in atmospheric neutrino experiments, Kashiwa, Dec. 2004.

48

APPENDICES

A. ICRR International Workshops
B. ICRR Seminars
C. List of Publications(a) Papers Published in Journals(b) Conference Papers(c) ICCR Report

D. Doctoral Theses
E. Public Relations(a) ICCR News(b) Public Lectures(c) Visitors

F. Inter-University Researches
G. List of Committee Members(a) Board of Councillors(b) Executive Committee(c) Advisory Committee

H. List of Personnel

49

A. ICRR International Workshops

RCCN International Workshop“sub-dominant oscillation effects in atmospheric neutrino experiments”
Date: 9 11 December, 2004
Place: Institute of Cosmic Ray Research, Kashiwa, Japan.
OutlineAtmospheric neutrino experiments have been studying νμ  ντ oscillations extensively. The zenith angle and energy
dependent deficit, or more recently the LE dependent deficit, of atmospheric muon neutrinos has been used to constrain theneutrino oscillation parameters: sin2 2θ23  0 9 and 1 9  Δm223  3 0  10 3 eV2As a natural extension of the present νμ  ντ oscillation studies, atmospheric neutrino experiments should study threeflavor oscillation effects. One is the effects driven by θ13 and the other is the ones driven by the solar oscillation terms (θ12 andΔm212 ). Especially, thanks to the diameter of the Earth, atmospheric neutrino experiments are, in principle, sensitive to the solar
neutrino oscillation terms. If the solar neutrino oscillation effect is observed, it is possible to get unique information, such as adiscrimination of θ23 larger or smaller than 45 degree for non-maximal sin2 2θ23 .
It is predicted that the effect of solar neutrino oscillations on atmospheric neutrino experiments is relatively a small effect,
and improvements of the accuracy of our understanding of the flux and the neutrino interaction cross sections are required. In
addition, very careful evaluations of the systematic uncertainties are required. These topics will be discussed in detail in thisworkshop.

Participants
20 from Japan, 3 from U.S.A., 2 from Italy, 2 from Rusia, 2 from Spain, 1 from U.K., 1 from Germany, 1 from Rep. of Korea

Toward Very High Energy Particle Astronomy 5 (VHEPA-5)
Date: 7 8 March, 2005
Place: Institute of Cosmic Ray Research, Kashiwa, Japan.
OutlineThe 5th annual meeting on High Energy Particle Astronomy, “Toward Very High Energy Particle Astronomy 5” aimed to
explore and understand recent rapid process on very high energy particle astronomy/astrophysics both in observational andtheoretical sides. We had twelve invited talks on various subjects. The highlights of this year’s meeting were two topics. Firstone was the latest results of H.E.S.S. presented by Dr. S. Funk from Max-Planck Institute fuer Kernphysik. In particular the
discovery of 8 new TeV sources on the Galactic plane, among which 6 are unknown from optical, radio or X-ray’s observations,
was astonishing. This is the first time ever new astronomical sources are found by the Cherenkov telescope. Cangaroo team too
showed their results after this talk. Second excitement of the meeting was recent gigantic gamma-ray flare of SGR 1806-20.
Dr. N. Kawai presented a review on this hottest subject. There were also a couple of related talks on this giant flare event,mostly from the theoretical point of view.

Participants
40 from Japan, 1 from Germany, 1 from Swiss

50

B. ICRR Seminars
Date

Lecturer

Title

April 19, 2004

KAKIZAKI Mitsuru“Flavor structure in super-symmetric models”(ICRR, University of Tokyo)

April 22, 2004

ASAOKA Yoichi(ICRR, University of Tokyo)“Ashra Experiment : The status of thedetector development”

April 23, 2004

Alexei Yu Smirnov(Tokyo Metropolitan Univ.
International Center for
Theoretical Phys., Trieste
Institute for Nuclear Research of
Russian Academy of Sciences)“Toward precision measurements in solar
neutrino”

April 26, 2004

SENAMI Daishi“Neutrino mass and Leptogenesis via
multi-scalar field evolution”(ICRR, University of Tokyo)

May 10, 2004

TAKAHASHI Satoshi(ICRR, University of Tokyo)“Toward Understanding the Dark Sideof the Universe”

May 13, 2004

Dr. Ralph Engel(Institute of Nuclear Phys.,
Forschungszentrum Karlsruhe)“Status and Prospect of the Pierre Auger
Project”

May 17, 2004

HARADA Junpei(ICRR, University of Tokyo)“Hypercharge and baryon minus lepton numberin E-6 GUTcd ht rs”

MATSUMOTO Shigeki“Explosive Dark Matter Annihilation”

May 24, 2004

(ICRR, University of Tokyo)

June 9, 2004

KANEYUKI Kenji“New result of K2K”

(ICRR, University of Tokyo)

June 14, 2004

KITANO Ryuichiro(Princeton University)

July 5, 2004

YAMAMOTO Tokonatsu(Chicago University)

October 18, 2004

Yoshimura Motohiko(Okayama University)

October 27, 2004

TANAHASHI Seiji(Tohoku University)

November 18, 2004

J. J. Gomez Cadenas(Valencia and KEK)

November 18, 2004

NISHIMURA Jun(Emeritus Professor in the University ofTokyo and ISAS)

November 25, 2004

TANAKA Reizaburo(Okayama University)“Anomaly mediation, electroweak symmetry
breaking and baryo/lepton genesis”“Observation of urtra-high energy cosmic ray
by Pierre Auger Project”“Towards resolution of the mass hierarchy
problem in a cosmological context”“The Structure of Corrections to Electroweak
Interactions in Deconstructed Higgsless Models”“Effect of correlations and degeneracies
in future neutrino oscillation experiments”“Cosmic Ray Electrons and Supernova
Acceleration”“Prospects for TeV Scale New Physics
at LHC”

51

Date

Lecturer

Title

December 9, 2004

TAKAYAMA Humihiro“SuperWIMP dark matter and the related
topics”(University of California)

January 5, 2005

MURAKI Yasushi(Nagoya University)“Cosmic ray acceleration in thundercloudand effect on climate”

C. List of Publications — 2004 fiscal year(a) Papers Published in Journals
1. “Search for Dark Matter WIMPs using Upward Through-going Muons in Super-Kamiokande,” S. Desai et al. (SuperKamioknade collaboration), Phys. Rev. D70 (2004) 083523; D70 (2004) 109901.
2. “Atmospheric neutrinos,” T. Kajita, New Jour. Phys. 6 (2004) 194.
3. “A measurement of atmospheric neutrino oscillation parameters by Super-Kamiokande-I,” Y. Ashie et al. (SuperKamioknade collaboration), Submitted to Phy. Rev. D, hep-ex/0501064.
4. “Search for nucleon decay via modes favored by supersymmetric grand unification models in Super-Kamiokande-I,”
K. Kobayashi et al. (Super-Kamioknade collaboration), Submitted to Phy. Rev. D, hep-ex/0502026.
5. “Measurement of single π 0 production in neutral current neutrino interactions with water by a 1.3 GeV wide band muon
neutrino beam,” S. Nakayama et al. (K2K collaboration), Submitted to Phys. Lett. B, hep-ex/0408134.
6. “The K2K SciBar detector,” K. Nitta et al., Nucl. Inst. Meth. A535 (2004) 147.
7. “Evidence for muon neutrino oscillation in an accelerator-based experiment,” E. Aliu et al. (K2K collaboration), Phys.
Rev. Lett. 94 (2005) 081802.
8. “A new calculation of the atmospheric neutrino flux in a 3 dimensional scheme,” M. Honda et al., Phys. Rev. D, 70(2004) 043008.
9. “Atmospheric Neutrinos,” T. Kajita, New J. of Phys., 6 (2004) 194.
10. “Detection of Gamma Rays around 1 TeV from RX J0852.0-4622 by CANGAROO-II,” H. Katagiri, R. Enomoto.,
L. T. Ksenofontov and M. Mori et al [CANGAROO Collaboration], Astrophys. J. 619, L163 (2005) [arXiv:astroph/0412623].
11. “Detection of sub-TeV Gamma-rays from the Galactic Center Direction by CANGAROO-II,” K. Tsuchiya, R. Enomoto,
L. T. Ksenofontov, M. Mori and T. Naito, et al. [CANGAROO-II Collaboration], Astrophys. J. 606, L115 (2004)[arXiv:astro-ph/0403592].
12. “A Search for TeV Gamma-ray Emission from the PSR B1259-63/SS2883 Binary System with the CANGAROO-II
10-m Telescope,” A. Kawachi, Y. Naito, J. R. Patterson, and P. G. Edward et al. [The CANGAROO-II Collaboration],
Astrophys. J. 607, 949 (2004) [arXiv:astro-ph/0402214].
13. “The naissance of high energy particle astronomy with ASHRA,” M. Sasaki, High Energy particle News, Volume23,
Number 2 (July/August, September 2004) 63–69.
14. “Observation by an Air-Shower Array in Tibet of the Multi-TeV Cosmic-Ray Anisotropy due to Terrestrial Orbital
Motion Around the Sun,” M. Amenomori et al., Phy. Rev. Lett., 93 (2004)061101-1-4.
15. “Automatic analysis of the emulsion chamber using the image scanner applied to the Tibet hybrid experiment,” S. Ozawa
et al., NIM, A523 (2004) 193–205.
16. “Ultrastable performance of an underground-based interferometer observatory for gravitational waves,” S. Sato, et al.,Physical Review D, 69(2004) 102005-1-102005-7.
17. “Coincidence analysis to search for inspiraling compact binaries using TAMA300 and LISM data,” H. Takahashi, et al.,Physical Review D, 70(2004) 042003-1-042003-17.

52

18. “Manufacturing of a 10-km-scale radius-of-curvature surface using a thin-film coating technique,” S. Miyoki, M. Ohashi,
K. Waseda, H. Karochi, T. Tomaru, to be published in Optics Letters.

(b) Conference Papers
1. M. Nakahata, “Super-Kamiokande and Solar Neutrinos,” Carolina Neutrino Workshop 2004, University of South Carolina, USA, April 16–17, 2004. (to be published)
2. Y. Suzuki, “Future Experiments on Sub-MeV Solar Neutrinos,” Solar Neutrino Physics after 50 years, Workshop in
honor of John Bahcall, University of Milano, Milano, Italy, May 6th, 2004. (to be published)
3. M. Nakahata, “Super-Kamiokande’s Solar Neutrino Results,” XXIst International Conference on Neutrino Physics andAstrophysics (Neutrino 2004), Paris, France, June 14–19, 2004. (to be published)
4. Y. Suzuki, “Future Solar Neutrino Experiments,” XXIst International Conference on Neutrino Physics and Astrophysics(Neutrino 2004), Paris, France, June 14–19, 2004. (to be published)
5. Y. Koshio, “XMASS,” The 12th International Conference on Supersymmetry and Unification of Fundamental Interactions (SUSY04), Tsukuba, Japan, June 17–23, 2004. (to be published)
6. S. Nakayama, “K2K NC π 0 production,” 6th International Workshop on Neutrino Factories and Superbeams (NuFact04),
Osaka, Japan, July 26–August 1, 2004. (to be published)
7. Y. Takeuchi, “SuperK: solar neutrino physics,” 6th International Workshop on Neutrino Factories and Superbeams (NuFact04), Osaka, Japan, July 26–August 1, 2004. (to be published)
8. C. Saji, “Recent results from Super-Kamiokande on atmospheric neutrino measurements,” 32nd International Conference
on High Energy Physics (ICHEP 2004), Beijing, China, August 16–22, 2004.
9. K. Ishihara, “Solar Neutrino Measurement at Super-Kamiokande,” 32nd International Conference on High Energy
Physics (ICHEP 2004), Beijing, China, August 16–22, 2004. (to be published)
10. Y. Takeuchi, “Recent Status of the XMASS Project,” 32nd International Conference on High Energy Physics (ICHEP
2004), Beijing, China, August 16–22, 2004. (to be published)
11. Y. Suzuki, “Super-Kamiokande Results (Neutrino Oscillations),” Nobel Symposium on Neutrino Physics, Haga Slott,
Sweden, August 19–24, 2004. (to be published)
12. Y. Koshio, “Day/Night asymmetry in SK and in larger detectors,” International Workshop on “Neutrino Oscillation
Workshop 2004” (NOW2004), Otranto, Italy, September 11–17, 2004. (to be published)
13. M. Nakahata, “Future solar neutrino experiments,” International Workshop on “Neutrino Oscillation Workshop 2004”(NOW2004), Otranto, Italy, September 11–17, 2004. (to be published)
14. S. Moriyama, “Super-Kamiokande atmospheric neutrinos,” International Workshop on “Neutrino Oscillation Workshop
2004” (NOW2004), Otranto, Italy, September 11–17, 2004. (to be published)
15. S. Moriyama, “XMASS,” 5th International Workshop on the Identification of Dark Matter (IDM2004), Edinburgh, Scotland, September 6–10, 2004. (to be published)
16. C. Mitsuda, “Recent Status of the XMASS Project,” 5th International Heidelberg Conference on DARK MATTER IN
ASTRO AND PARTICLE PHYSICS (DARK 2004), Texas, USA, October 3–9, 2004. (to be published)
17. S. Moriyama, “XMASS,” The Future of Dark Matter Detection, Chigago, USA, December 9–10, 2004. (to be published)
18. S. Nakayama, “Effect of solar terms to θ23 determination in Super-Kamiokande and important systematic errors for future
improvements,” RCCN International Workshop sub-dominant oscillation effects in atmospheric neutrino experments,
Kashiwa, Japan, December 9–11, 2004. (to be published)
19. M. Nakahata, “Japanese and eastern facilities,” Topical Workshop in Low Radioactivity Techniques (LRT2004), Sudbury,
Canada, December 12–14, 2004. (to be published)
20. Y. Takeuchi, “Distillation purification of liquid xenon,” Topical Workshop in Low Radioactivity Techniques (LRT2004),
Sudbury, Canada, December 12–14, 2004. (to be published)

53

21. S. Moriyama, “Self-shielding effect for liquid xenon,” Topical Workshop in Low Radioactivity Techniques (LRT2004),
Sudbury, Canada, December 12–14, 2004. (to be published)
22. A. Minamino, “XMASS experiment : Dark matter search with liquid Xe detector,” KEKPH meeting, KEK, Japan, March
03–05, 2005. (to be published)
23. Y. Koshio, “The current status of XMASS using 100 kg detector,” Applications of Rare Gas Xenon to Science andTechnology (XeSAT2005), Waseda, Japan, March 8–10, 2005. (to be published)
24. Y. Takeuchi, “Distillation purification of xenon for krypton and measurement of radon contamination in liquid xenon,”Applications of Rare Gas Xenon to Science and Technology (XeSAT2005), Waseda, Japan, March 8–10, 2005. (to be
published)
25. T. Kajita et al., “Future atmospheric neutrino experiments - the case of water Cherenkov detectors,” The 5th workshop on“Neutrino Oscillations and their Origin” (NOON2004), (11–15 February 2004, Odaiba, Tokyo, Japan) to be published.
26. T. Kajita, “Atmospheric neutrinos - past, present and future,” 7th International Conference on Heavy Quarks and Leptons,(1–5 June 2004, San Juan, Puerto Rico) to be published.
27. T. Kajita, “Atmospheric neutrinos - present and future,” 6th International Workshop on Neutrino Factories and Superbeams, (26 July–1 August 2004, Osaka, Japan) to be published.
28. T. Kajita, “Super-Kamiokande atmospheric neutrino results and long baseline experiments in Japan,” Neutrino Oscillation Workshop (NOW2004), (11–17 September 2004, Otranto, Italy) to be published.
29. K. Okumura, “SK atmospheric Neutrino Results,” Second NOVE International Workshop on NEUTRINO OSCILLATION IN VENICE, (3–5 December, 2003 Venice Italy) 129133.
30. L. T. Ksenofontov, R. Enomoto, H. Katagiri and K. Tsuchiya et al. [CANGAROO Collaboration], “Search for TeV
Gamma-rays from the Remnant of SN 1987A,” 2nd VERITAS Symposium on TeV Astrophysics of Extragalactic
Sources, Chicago, Illinois, 24–26 Apr. 2003, New Astron. Rev. 48, 485 (2004).
31. H. Kubo et al. [CANGAROO Collaboration], “Status of the CANGAROO-III Project,” 2nd VERITAS Symposium on
TeV Astrophysics of Extragalactic Sources, Chicago, Illinois, 24–26 Apr. 2003, New Astron. Rev. 48, 323 (2004).
32. M. Mori, “Recent Topics from Very High Energy Gamma-ray Astrophysics,” 6th RESCEU Symposium “Frontier in
Astroparticle Physics and Cosmology,” Sanjo hall, Univ. Tokyo (November 04–07, 2003), “Frontier in Astroparticle
Physics and Cosmology” (eds. K. Sato and S. Nagataki, Universal Academy Press, Tokyo, 2004)
33. A. Kawachi, T. Naito and S. Nagataki, “High Energy Emissions from the PSR1259-63/SS2883 Binary System,” Proceeding of the International Symposium on High Energy Gamma-Ray Astronomy Heidelberg, Germany (July 26–30,
2004) (to be published)
34. M. Mori [CANGAROO Collaboration], “Recent Results from CANGAROO-II & CANGAROO-III,” ibid.
35. Ken’ichi Tsuchiya [CANGAROO Collaboration], “Detection of Sub-TeV Gamma-rays from the Galactic Center with
the CANGAROO-II Telescope,” ibid.
36. M. Ohishi, M. Mori and M. Walker, “Gamma-Ray Spectra due to Cosmic-Ray Interactions with Dense Gas Clouds,”
ibid.
37. K. Tsuchiya, “Status of an Atmospheric Cherenkov Imaging Camera for the CANGAROO-III Experiment and Perspectives of the Field,” Fourth International Conference on Physics Beyond the Standard Model “BEYOND THE DESERT”’03 (June 9–14, 2003, Castle Ringberg, Tegernsee, Germany) published as “Beyond the Desert 2003” (ed. KlapdorKleingrothaus, H.-V., Springer, Heidelberg, 2004), pp. 819–830.
38. I. Kouga et al., “The simulation for the aerosol analysis using imaging rider,” 23th laser sensing symposium (16–17
September 2004, Tsukuba, Japan) 107–108.
39. M. Sasaki et al., GRB041211: Ashra Prototype optical observation, GCN report (2846).
40. M. Takita, “Extensive Air Shower Observation for Energy  1017 eV and Related Topics,” the 28th International Cosmic
Ray Conference, (31 July–7 August 2003, Tukuba, Japan), vol. 8, 277–297.
41. S. Miyoki and LCGT Collaborations. “Large scale cryogenic gravitational wave telescope,” (5–9, Sep., Seattle, Washington, USA) Nuclear Physics B - Proceedings Supplements, Volume 138, January 2005, Pages 439–442.

54

42. T. Akutsu et al., “Burst wave analysis of TAMA300 data with the ALF filter,” The 9th annual Gravitational Wave Data
Analysis Workshop (GWDAW-9), (15–18 December 2004, Annecy, France) to be published.
43. T. Uchiyama et al., “Present status of CLIO in 2004,” 19th European Cosmic Ray Symposium (Aug. 29–Sep. 5, 2004,
Firenze, Italy) to be published.
44. M. Ohashi et al., LCGT Project, 19th European Cosmic Ray Symposium (Aug. 29–Sep. 5, 2004, Firenze, Italy) to be
published.
45. N. Sato et al., “Underground Interferometers in Japan,” 2005 Aspen Winter Conference on Gravitational Waves Gravitaional Wave Advanced Detectors (GWADW) (16–22 January 2005, Aspen, Colorad, USA).
46. N. Sato et al., “Making a data analysis processor with FPGA for a gravitational-wave event searc,” 19th European Cosmic
Ray Symposium (Aug. 29–Sep. 5, 2004, Firenze, Italy).

(c) ICRR Report
1. ICRR-Report-507-2004-5 (August 23, 2004)“Unstable state decay without exponential Law-small Q value s-wave (sQs) decay,”
T. Jittoh, S. Matsumoto, J. Sato, Y. Saito and K. Takeda
2. ICRR-Report-508-2004-6 (August, 2004)“Big-bang nucleosynthesis and harmonic decay of low-lived massive particles,”M. Kawasaki, K. Kohri and T. Moroi
3. ICRR-Report-509-2004-7 (July, 2004)“making waves on CMB power spectrum and inflation dynamics,”
M. Kawasaki, F. Takahashi and T. Takahashi
4. ICRR-Report-510-2004-8 (September, 2004)“Early reionization by decaying particles and cosmic microwave background radiation,”
S. Kasuya and M. Kawasaki
5. ICRR-Report-511-2004-9 (October 18, 2004)“Life-time entropy production due to the decay of domain walls,”
M. Kawasaki and F. Takahashi
6. ICRR-Report-512-2004-10 (October 1, 2004)“Constraining neutrino masses by CMB experiments alone,”K. Ichikawa, M. Fukugida and M. Kawasaki
7. ICRR-Report-513-2004-11 (December 29, 2004)“Non-perturbative effect on dark matter annihilation and gamma ray signature from galactic center,”J. Hisano, S. Mastumoto, M.M. Nojiri and S. Saito
8. ICRR-Report-513-2004-12“Significant effects of second KK particles on LKP dark matter physics,”
M. Kakizaki, S. Matsumoto, Y. Sato and M. Senami
9. ICRR-Report-515-2004-13 (March, 2005) “Detection of TeV gamma-rays from the supernova remnant RX J0852.04622,”
H. Katagiri
10. ICRR-Report-516-2004-14 (March, 2005) “Very high energy gamma-ray observations of the Galactic Center with the
CANGAROO-II telescope,”
K. Tsuchiya
11. ICRR-Report-517-2004-15 (March, 2005) “Very high energy gamma-ray observations of the Galactic plane with the
CANGAROO-III telescope,”
M. Ohishi

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D. Doctoral Theses
1. “Very High Energy Gamma-ray Observations of the Galactic Center with the CANGAROO-II Telescope,”
Ken-ichi Tsuchiya, Ph. D. Thesis, University of Tokyo, Mar. 2005
2. “Very High Energy Gamma-ray Observations of the Galactic Plane with the CANGAROO-III Telescopes,”
Michiko Ohishi, Ph. D. Thesis, University of Tokyo, Mar. 2005
3. “Observation of TeV Gamma-rays from the Active Radio Galaxy Centaurus A with CANGAROO-III Imaging Atmospheric Cherenkov Telescope,”
Shigeto Kabuki, Ph. D. Thesis, University of Tokyo, Jun. 2005
4. “The sub-hundred GeV γ -ray astronomy by MAGIC,”
Keiichi Mase, PhD Thesis, University of Tokyo, Feb. 2005
5. “Study of coincidence analysis to search for gravitational waves from inspiraling compact binaries using interferometer
data,”
Hiroyuki Takahashi, PhD Thesis, Niigata University, Mar. 2005
6. “Exploring Clusters of Galaxies,”
Hiroshi Ohno, PhD Thesis, University of Tokyo, Mar. 2005
7. “Cosmic Microwave Background Constraint on Neutrino Mass,”
Kazuhide Ichikawa, PhD Thesis, University of Tokyo, Mar. 2005
8. “Origin and Evolution of Large Lepton Asymmetry,”
Hiromitsu Takahashi PhD Thesis, University of Tokyo, Mar. 2005

E. Public Relations(a) ICRR NewsICRR News is a newspaper published quarterly in Japanese to inform the Institute’s activities. This year’s editors were K.
Okumura and M. Ohashi. It includes : (1) reports on investigations by the staff of the Institute or made at the facilities of theInstitute, (2) reports of international conferences on topics relevant to the Institute’s research activities, (3) topics discussed at
the Institute Committees, (4) list of publications published by the Institute [ICRR-Report, ICRR-Houkoku(in Japanese)], (5)list of seminars held at the Institute, (6) announcements, (7) and other items of relevance. The main topics in the issues in 2004
fiscal year were :No. 54 (May 10, 2004)
Statements from old and new directors (Y. Suzuki and M. Yoshimura)No. 55 (July 22, 2004)
Emperor and Empress visiting to Super-KamiokandeCompletion of CANGAROO-III telescope (M. Mori)
K2K new result (K. Kaneyuki)
Report on Neutrino 2004 conference (S. Moriyama)No. 56 (March 18, 2005)
Report on Kyodo-riyo workshop (M. Shiozawa)
ICRR open house (H. Sagawa)

56(b) Public Lectures“ICRR public lecture on Neutrino,” May 15 2004, Kashiwa.
“Oepning address,” Takaaki Kajita, ICRR, “Supernova explosion and neutrinos - how the black holes and neutron stars
are formed,” Katsuhiko Sato, University of Tokyo, “IceCube - a neutrino telescope in the south pole -,” Shigeru Yoshida,
Chiba University,“JPS public lecture on Neutrino,” November 6 2004, Chuo University.
“Measuring the neutrino mass using the earth,” Takaaki Kajita, ICRR, “Solving the solar neutrino problem by KamLAND,” Kunio Inoue, Tohoku University, “Unified theories of particle physiscs - message from neutrinos -,” Masako
Bando, Aichi University,(c) Visitors
Visitors to Kamioka Observatory. (Total: 210 groups, 4380 people)
Masayuki Hara (Vice Governor of Gifu prefecture) [May 21, 2004]
Yutaka Nakaoki (Governor of Toyama prefecture) [June 29, 2004]
Emperor Akihito and Empress Michiko; Yamato Inaba (Vice-Minister of MEXT); Masatoshi Koshiba (Professor emeritus, Univ. of Tokyo); Takeshi Sasaki (President of Univ. of Tokyo); Toshitsugu Fujii (Vice-President of Univ. of Tokyo);
Taku Kajiwara (Governor of Gifu prefecture); Toyotaro Iwai (Chairperson of Gifu prefectural assembly); Katsumi Funasaka (Mayor of Hida city); Takashi Ishida (Chairperson of Hida city council) [July 13, 2004]
Eiichi Yamashita (Upper House member) [August 24, 2004]
Financial committee members in Lower House [September 8, 2004]
Norihisa Tamura (Parliamentary Secretary of MEXT) [September 28, 2004]
Kanji Fujiki (Counselor, Personnel Division, Secretariat of MEXT) [Nobember 25, 2004]
Ikuo Shirokane (General Manager, Tokai Office, Japan Post) [December 7, 2004]
Kouichi Matsushita (President, Fujitsu Chubu Systems) [January 18, 2005]
Naoto Kan (Lower House member) [February 6, 2005]
MEXT Super Science High-school (SSH): total 8 schools
Visitors to International Astrophysical Observatory, Woomera, Australia.
Mr. Satoshi Shinki (Ministry of Education, Science and Culture); Mr. Toshihito Suzuki, Mr. Naoki Saito, Mr. Toshiharu
Oki (Administration section, University of Tokyo) [January 18–19, 2004]
Mr. Keikichi Koike, Mr. Toru Nishimura, Mr. Ryoji Shigeta (Shingei Film Co. Ltd.) [December 14–15, 2004]

57

F.

Inter-University Researches

Facility Uses
Kamioka Observatory
Norikura Observatory
Akeno Observatory
Research Center for Cosmic neutrinos
Emulsion and Air Shower Facilities in Kashiwa
Low-level Radio-isotope Measurement Facilities
in Kashiwa
OthersCollaborative Researches
High Energy Cosmic Ray Researches
in the Underground and Deep Sea
High Energy Cosmic Ray Researches
in Flyers and at High Altitude and Ground
High Energy Gamma Ray Source Researches
Chemical Composition
and Isotope Measurement
Development of Observational Methods
and Instruments
Theoretical Researches or
Rudimental Researches
OthersConferences
Special Activity on Abroad

Applications

Adoptions

Researchers

20
11
11
2
3
4

20
11
11
2
3
4

446
91
193
13
66
28

1

1

22

3

3

12

11

9

71

8
2

8
2

83
13

3

3

80

2

1

8

5
4

4
4

29
79

58

G. List of Committee Members(a) Board of Councillors
SUZUKI, Yoichiro
KURODA, Kazuaki
KAJITA, Takaaki
FUKUSHIMA, Masaki
OKAMURA, Sadanori
KOMIYAMA, Hiroshi
KOBAYASHI, Makoto
KUGO, Taichi
KAIFU, Norio
SATO, Humitaka
TOKI, Hiroshi
SUZUKI, Atsuto
OHTA, Itaru
INOUE, Hajime
ITO, Nobuo

ICRR, The University of Tokyo
ICRR, The University of Tokyo
ICRR, The University of Tokyo
ICRR, The University of Tokyo
The University of Tokyo
The University of Tokyo
KEK
Kyoto University
National Astronomical Observatory
Kounan University
Osaka University
Tohoku University
Utsunomiya University
Institute of Space and Astronautical Science
Osaka City University(b) Executive Committee
SUZUKI, Yoichiro
TORII, Shoji
MURAKI, Yasushi
TANIMORI, Toru
KAJINO, Fumiyoshi
SAKAI, Norisuke
ENYO, Hideto
WATANABE, Yasushi
NAKAMURA, Takashi
MINOWA, Makoto
KURODA, Kazuaki
FUKUGITA, Masataka
FUKUSHIMA, Masaki
NAKAHATA, Masayuki
TAKITA, Masato

ICRR, The University of Tokyo
Kanagawa University
Nagoya University
Kyoto University
Konan University
Tokyo Institute for Technology
RIKEN
Tokyo Institute for Technology
Kyoto University
The University of Tokyo
ICRR, The University of Tokyo
ICRR, The University of Tokyo
ICRR, The University of Tokyo
ICRR, The University of Tokyo
ICRR, The University of Tokyo(c) Advisory Committee
KAJINO, Fumiyoshi
MURAKI, Yasushi
TORII, ShojiTANIMORI, Toru
KAWAKAMI, Saburo
YANAGITA, Shohei
MUNAKATA, Kazuki
SHIBATA, Makio
MATSUBARA, Yutaka
HOTTA, Naoki
OGIO, Shoichi
SAKURAI, Takahisa
SHIOZAWA, Makoto
MORI, Masaki
OHASHI, Masataka
KANEYUKI, Kenji
TAKITA, Masato
HISANO, Junji

Konan University
Nagoya University
Kanagawa University
Kyoto University
Osaka City University
Ibaraki University
Sinshu University
Yokohama National University
Nagoya University
Utsunomiya University
Tokyo Institute for Technology
Yamagata University
ICRR, The University of Tokyo
ICRR, The University of Tokyo
ICRR, The University of Tokyo
ICRR, The University of Tokyo
ICRR, The University of Tokyo
ICRR, The University of Tokyo

59

H. List of Personnel
Director

SUZUKI Yoichiro

Vice Director

KURODA Kazuaki

Neutrino and Astroparticle Division
Kamioka Observatory
Professor
Associate Professor
Research Associate
Technical Staff
Research Fellows
Administrative Chief
Secretary

SUZUKI Yoichiro,
TAKEUCHI Yasuo,
KOSHIO Yusuke,
KAMEDA Jun,MIZUHATA Minoru,
NAMBA Toshio,
AKIMOTO, Masatoshi,
OKURA Yoko,
OKADA Eri

Research Center for Cosmic Neutrinos
Professor
KAJITA Takaaki
Associate Professor
KANEYUKI Kenji
Research Associate
OKUMURA Kimihiro
Technical Staff
SHINOHARA Masanobu
Research Fellows
SAJI Choji,
HONDA Morihiro,
Secretary
FUKUDA Yoko,

NAKAHATA Masayuki
SHIOZAWA Masato,
OBAYASHI Yoshihisa,
TAKEDA Atsushi,
FURUTA Takashi
MITSUDA Chikaori,
GOTO Shohachiro
MURAKI Yasuko,

NAKAYAMA Shoei,
ISHITSUKA Masaki
KITSUGI Atsuko

MORIYAMA Shigetaka
MIURA Makoto,
ISHIHARA Kenji
OGAWA Hiroshi
MAEDA Yukari,

HIGUCHI Itaru,

High Energy Cosmic Ray Division
Professor
Associate Professor
Research Associate
Technical Staff
Research Fellows

Secretary
AKENO Observatory
Associate Professor
Technical Staff

FUKUSHIMA Masaki,
ENOMOTO Ryoji,
SASAKI Makoto
HAYASHIDA Naoaki,
OHNISHI Munehiro,
AOKI Toshifumi,
OZAWA Shunsuke,
MIYAMOTO Tomoko,
HARA Satoshi,
KAWATA Kazumasa,
SAWANO Sumiko,
SOMENO Hidemasa,
KOBAYASHI Noriko,

SAGAWA Hiroyuki
TORII Reiko,
SUDA Takayuki

NORIKURA Observatory
Technical Staff
YAMAMOTO Kuniyuki,
YOKOYAMA Chiaki,

MORI Masaki
YOSHIKOSHI Takanori,

TAKITA Masato,

HATANO Yoshikazu,
ASAOKA Yoichi,
TOYODA Setsuko,
UDO Shigeharu,
KOBAYASHI Noriko,
SHIOMI Atsushi,
HUANG Jing,
HATSUMI Ryoichi,
FUKUDA Hisahe
TANAKA Akiko,

KAWACHI Akiko,
TAKEDA Masahiro
KOBAYASHI Takahide
TOKUNOU Hisao,
TSUCHIYA Kenichi,
TSUCHIYA Harufumi,
TANAKA Akiko,
KANAZAWA Shuji,

OHOKA Hideyuki,

KAWAGUCHI Masami

AGEMATSU Yoshiaki,
SHIMODAIRA Hideaki,

USHIMARU Tsukasa,
YOSHIDA Takeji

SAWANO Sumiko

60

Astrophysics and Gravity Division
Professor
Associate Professor
Research Associate
Technical Satff
Research Fellows

FUKUGITA Masataka,
OHASHI Masatake,
MIYOKI Shinji,
ISHITSUKA Hideki
YAMAMOTO Kazuhiro,
SENAMI Masato,
HARADA Junpei,

KURODA Kazuaki,
YASUDA Naoki,
UCHIYAMA Takahasi

KAWASAKI Masahiro
HISANO Junji

HAYAKAWA Hideaki,
MATSUMOTO Shigeki,
KAKIZAKI Mitsuru,

OKADA Atsushi,
TAKAHASHI TomoYAMAGUCHI Yoshiyuki,

TASHIRO Megumi,

TAKAHASHI Junko,

KITA Aiko

Administration Division
Project Division
Head
General Affairs Section
Head
Staff

NAKATSUKA Kazuo

MATSUZAKI Hiroshi
HAMANO Teruko,
AKIYAMA Makiko

Public Relations and Project Section
Head
FUJIEDA Yuichi
Secretary
KOKUBUN Yayoi,
Library Office
Staff

SAITO Akiko

Graduate Students
MASE Keiichi(D3),OHISHI, Michiko(D3),KIUCHI Ryuta(M2),YUKAWA Ryohei(M1),
MANAGO, Naohiro(D3),
AITA Yuichi(M2),
MINAMINO Akihiro (D1),MITSUKA Gaku (M1),
KONDO, Kazuhiro(D3),TOKUNARI Masao(D1),
KAMAGASAKO Shogo(M1),
TAKAHASHI Huminobu(D3),KONYA Kenichiro(M2),
KANZAKI Toru(M1),

TAKETA Akimichi(M1),KABUKI, Shigeto(D3),KAWASAKI Sho(M1),
SAITO Takayuki(M1),JOBASHI Masashi(D3),OKUMURA Akira(M2),
HOSAKA Junya (M2),NISHINO Haruki (M1),
KASAHARA Kunihiko(D2),AKUTSU, Tomomi (M2),
OGURO Kei (M2),
ICHIKAWA Kazuhide(D3),
TAKAHASHI Hiroyuki(M1),NISHIHARA Tatsuo(M1).

ADACHI, Yuki(M2),
YUASA Midori(M1),
MASUDA, Masataka(D3),NODA Koji(M1),
TAKENAGA Yumiko(M2),
OKUTOMI Akira(D2),KAMAGASAKO Shogo(M1),
SAITO Osamu(M2),
NAGAI Minoru(M2),
TAKAYAMA Tsutomu(M1),