Future Perspectives on the UHECR Observatories
and the FAST Project

Toshihiro Fujii, ICRR Seminar, 2016/May/24

1Energy Cosmic Rays with Telescope Array Fluorescence Detectors in Monocular Mode”
Education2012 Ph.D.
in Science,
University
Supervisor:
Prof. Osaka
ShoichiCity
Ogioof theUniversity
Energy Spectrum and the Mass Composition of Ultra-High2010 Thesis:
M.S. in “Measurements
Science, Osakafujii@icrr.u-tokyo.ac.jp
City
Cosmic Rays
Array Fluorescence Detectors in Monocular Mode”
2008 Energy
B.S. in Science,
Osakawith
CityTelescope
University
Office: 303, Hometown:
Nara,Prof.
Ikaruga-cho,Supervisor:
Shoichi Ogio Horyuji (法隆寺)Contact
InstituteforCosmicRay
Research+81-4-7136-3182
2010
M.S. in Science,
Osaka City University
Professional
2016
present
JSPS
Postdoctoral
Fellow (PD),Institute for Cosmic Ray Research,
Information
UniversityofTokyofujii@icrr.u-tokyo.ac.jp
2008 B.S. in Science,
Osaka
City
University
Experience
UniversityofTokyo
5-1-5 Kashiwanoha Kashiwa Rm.303
http://kicp.uchicago.edu/∼fujii
2015 - 2016
Postdoctoral Scholar, Kavli Institutefor Cosmological
Physics,
Chiba,
277-8582 Japan
Citizenship:
Japan
Universityof Chicago
Professional
2016 - present JSPS
Postdoctoral
Fellow (PD), Institute for Cosmic Ray Research,2013 - 2016
Collaborative
Researcher, Institute for Cosmic Ray
Research,
Experience
Universityof Tokyo
Curriculum
Vitae (May 2016)
Position
JSPS Postdoctoral Fellow (PD), Institute for Cosmic Ray Research, Universityof Tokyo
University of Scholar,
Tokyo Kavli Institute for Cosmological Physics,
2015 - 2016
Postdoctoral
2013 - 2015
Japan Society
for the Promotion of Science (JSPS) Postdoctoral Fellow
Universityof Chicago
Research
CosmicRay Physics,
EnergyCosmicRays
Fluorescence
Detector,
KavliUltra-High
InstituteResearcher,for
Cosmological
Physics,
Universityof Chicago
2013- 2016
Collaborative
Institutefor(UHECR),
CosmicRay
Research,
Interests
EnergySpectrum,University
MassPostdoctoral
Composition,
Ultra-High
Neutrinoand γ-ray,
Particle Astronomy
2012 - 2013
JSPS(PD),Energy
Institutefor Cosmic
Ray Research,of TokyoFellow
Universityof Tokyo2013 - 2015
Japan
Societyfor the Promotion of Science (JSPS) Postdoctoral Fellow
Education2012 Ph.D. in Science,
Osaka for
CityCosmological
University Physics, University of Chicago
Kavli Institute
Thesis: “Measurements
theFellow
EnergySpectrum
andfortheCosmic
Massof Ultra-High
2012
JSPS
(PD),
InstituteRay
Research,Research Grants
2016 -- 2013present
JSPS Postdoctoral
ResearchofFellowship
for
Young
Scientists(PD),Composition
Grant
Number
16J04564,
Energy Cosmic
Raysofwith
TelescopeTitle:
Array“Observation
Fluorescence ofDetectorsin Monocular
Mode”
University
Tokyo
and Fellowships
Principal
Investigator,
Ultra-High
EnergySupervisor:
Prof.Rays
ShoichiCosmicandOgio
Mass Composition with the New-Type Fluorescence Detector”,
2010 - present
M.S. in Science,
Osaka
City University
Universityof Tokyo,
3,400,000
JPY (Total
expense)Research Grants
2016
JSPS
Grant-in-Aid
for
JSPS Fellows,
Grantexpected
Number direct
16J04564,
2008
B.S. in Science,
Osaka
City University
2015 - present
JSPS Grant-in-Aid
for Young(A), Grant
Number 15H05443,
and Fellowships
Principal
Investigator,
Title:Scientist“Observationof Ultra-High
EnergyPrincipal
Investigator,
Title: of“Establishmentof New Technique to Observe
Rays
and Measurement
Their Mass Composition
Contact
Institute for CosmicCosmic
Ray Research+81-4-7136-3182
Professional
2016 - present
JSPS
Postdoctoral(PD),
for Cosmic
Ray Research,
Ultra-High
Energy Fellow
CosmicRays Instituteby
the New-Type
Fluorescence
Detector”,
Information
the
New-Type
Fluorescence
Detector”,
Universityof Tokyowith
fujii@icrr.u-tokyo.ac.jp
Experienceof
Universityof TokyoTokyo, 3,400,000
18,600,000
JPY(Total(Totalexpected
expecteddirect
directexpense)
expense)Tokyo,
JPY
5-1-5 Kashiwanoha University
Kashiwa
Rm.303
http://kicp.uchicago.edu/∼fujii
2015
2016
Postdoctoral
Scholar,InstituteforCosmological
Physics,
Chiba,
Citizenship:Japan
2013 -277-8582
2015 Japan
JSPS Postdoctoral
Fellowshipfor Research
Abroad,H25-339,
present
Grant-in-Aid
forKavli
Young
Scientist(A),
Grant Number
15H05443,
University
ChicagoUniversity ofofChicago, 10,512,000
JPY
Principal
Investigator,
Title: “Establishmentof New Technique to Observe2013- 2016
Collaborative
Researcher,
Institute
CosmicRay Fluorescence
Research,
2011Postdoctoral
2013
JSPS Grant-in-Aid
for
Fellows,
Grant
Number
11J02728,
Ultra-High
Cosmic
Rays
byfor
the
New-Type
Position
JSPS
Fellow(PD),Energy
InstituteforJSPS
CosmicRay
Research,
Universityof Tokyo Detector”,
University
Tokyo
Principal of
Investigator,
Title: JPY“EnergySpectrum
Ultra-High
EnergyTokyo, 18,600,000(Totalexpectedofdirect
expense)2013 -Ray
2015Physics,
Japan
Societyfor thePromotionofResearch
Science(JSPS)
Postdoctoral
Fellow
CosmicRaysEnergy
Measuredby Rays
thefor
Telescope
Array
Fluorescence
Detector”,
JSPS
Postdoctoral
Fellowship
Abroad,H25-339,
Research
CosmicUltra-High
Cosmic(UHECR),
Fluorescence
Detector,
Kavli
InstituteforUltra-High
Cosmological
Physics,
University
ChicagoUniversityof Chicago,
10,512,000
JPYInterests
Energy Spectrum, Mass
Composition,
EnergyNeutrinoand γ-ray,ofParticle
Astronomy

Toshihiro
FujiiFujii
Toshihiro

Toshihiro Fujii

2

This is how scientists see the world.
credit: http://abstrusegoose.com

3This is how scientists see the world.
credit: http://abstrusegoose.com

4This is how scientists see the world.
credit: http://abstrusegoose.com

Cosmic Rays

5

What are Cosmic Rays?

In

Proton(90%), Helium(8%) or
heavier nuclei (1%)

se!

Discovered by V. F. Hess (1912)

creaRadiations

Energetic particles injected from the universe.

E > 1019 eV, ultra-high energy
cosmic rays (UHECR)

Altitude [km]Anniversary on Aug. 7th 2012

Landing at Bad saarow, Germany on
Aug. 7th, 1912

Grandson of Hess

Memorial Stone

6

Cosmic Ray Anniversary on Aug. 7th 20127

Physics Goal of UHECR Astrophysics
Origin and Nature of Ultra-High Energy
Cosmic Rays (UHECRs) and Particle
Interactions at the Highest Energies
How frequent?
What kind of particle?
Where come from?
8How frequent?: Energy spectrum

2.5
E J(E)Knee-2-1
1 km yr

Larger energy
than the largest
accelerator
(LHC)
-2
-1
1 km century-2-1
1 m yrAnkle

Very infrequent,a large sensitive
area needed9he LAT instrument, the bulk of the
500 MeV. IC 443 is located in they emission produced by the pool of
weak relative to the region aroundeminga pulsar (2FGL J0633.9+1746)
2.5
map, W44 is the dominant source

synchrotron radiation seen in a high-resolution
radio continuum map of W44 (26) support thispicture. High-energy particles in the “crushedcloud” can be explained by reacceleration of thepreexisting galactic cosmic rays (25) and/or freshly accelerated particles that have entered thedense region (20). The mass of the shocked gas

20
Acceleration Scenario toward 10 eV
E

J(E)

Hillas&plot

B

C 443

W44
E2 dN/dE (erg cm-2 s-1)

10 -10Extragalactic origin?

10 -111

Galactic origin
Supernova remnants
IC443

10 -1281012

as measured
W44
ote thebesthaded bands
eling of thepoints for IC
ote the bestSCIENCE

Bottom-up model

10

Best-fit broken power law
Fermi-LAT
AGILE (19)π0-decay
Bremsstrahlung
Bremsstrahlung with Break
9

10

10

10

Energy (eV)

1011

1012

Neutron stars, Active
galactic nuclei,
Gamma ray bursts,
Radio galaxies,
Galactic clusters

Top-down model
fit pion-decay gamma-ray spectra, dashed lines denote the best-fit bremsstrahlung spectra, and dash-dotted lines denote the best-fit bremsstrahlung spectra
when including an ad hoc low-energy break at 300 MeV c−1 in the electron
Annihilation/decayofspectrum. These fits were done to the Fermi LAT data alone (not taking the TeV
super heavy relicdata points into account). Magenta stars denote measurements from the AGILE
satellite for these two SNRs, taken from (31) and (19), respectively.
particles, TopologicalVOL 339

15 FEBRUARY 2013

defect, Z-burst model809

γ : Lorentz factor of shock
Z : atomic number
B : magnetic field strength
R : size

Greisen-Zatsepin-Kuzmin (GZK) Cutoff
Cosmic'microwave'
background'radia5on'(CMBR)'
Cosmic'Ray

Earth
Cosmic'Ray

Interaction between UHE protons with energies above 1019.75 eV and CMBR via apion production. Heaver nuclei also interact with CMBR via photo-disintegration.
Mean free path : 50-100 Mpc (Nearby sources, compared to the universe size)

1 pc = 3.26 l.y. ~ 3×1016 m

Flux × E3/1024 (eV2 m-2 sr-1 s-1)

19.7
Expect suppression of flux above 10 eV.

Planck 201310
Preliminary

1
TA ICRC 2015
Auger ICRC 2015
-11017.5

18

18.5
19
19.5
log (E (eV))
10

20

20.5

11

How to detect very infrequentUHECRs?
= Extensive Air Shower (EAS)

Image credit: ASPERA/Novapix/L.Bret

12

How to observe Extensive Air Shower (EAS)
Longitudinal Development
Fluorescence
detector (FD)duty cycle: 15%

Xmax

Surface detector
array (SD)duty cycle: 100%
Lateral Density Distribution
CORSIKA software

https://web.ikp.kit.edu/corsika/movies13ᖺ ࡟ ࢫ ࢱࢺࡋࡓࠋࡑ ࡢ タ ィࡢ ㅖ ࣃ ࣛ ࣓ࢱࡣ

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History of Fluorescence Technique⾲ ࡟ ♧ ࡋ࡚ ࠶ ࡿ᰾◊ ✲ᡤ 㸦,16ࠋࠋࣇ ࣞ ࢿ ࣝ ࣞ ࣥ ࢬ ࡣ ᮾி኱ Ꮫཎ㸧ࡢ ࣐ ࢩ ࣥ ࢩ ࣙ ࢵ ࣉ ࡛ 〇 స ࡋᮦᩱ ࡣ ࣏ ࣓ࣜ ࢱ ࢡ ࣜࣝ 㓟࣓ ࢳ ࣝ

㸦300$

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㏱ 㐣 ᛶ ⬟ ࡣ ➨  ᅗ ࡟ ♧ ࡍ ㏻ ࡾ࡛ ࠶ ࡿ ᅗ ࡟ ♧ࡍ

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ࡉࢀ ࡓ ゅ ㏿ ᗘ ࢆ ➨  ᅗ ࡟ࠋ

✦ In 1958, proposal of fluorescence technique (Suga,

Oda＠Norikura symposium)ural University Miyako
College
Iwate Prefectural University Miyako College✦ Many photomultiplier tubes on the focal plane

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of Fresnel lens/mirror to observe fluorescencedÂę)«ŃŎƚfơƩ light.

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+ PMTs '1969|ŸdÂę)«ŃŎ(TOKYOZ1)ƙƧƪøÂǀǣ
Ǭǯƚ4ģÏŸ(´»ǳ[{ģÏ)ࠉ ➨  ᅗࠉ ࣉ ࣛ ࢫ ࢳ ࢵ ࢡ ࣞ ࣥ ࢬ ࡢ ග ㏱ 㐣ᛶ ⬟72.<2 Ѹ㸪
72.<2 Ѹ㸪
72.<2 Ѹ 㸧ࠉ ࠉ ࠉ㸦

Iwate Prefectural University Miyako College

Iwate Prefectural University Miyako College✦ Observe longitudinal profile including Xmax tobe sensitive to the mass composition of cosmicray.
✦ In 1969, first detection of fluorescencelight by TOKYO-1 (Tanahashi et al.
@Doudaira Observatory, Japan)୍ ୍
Reference: 空気シャワー観測による宇宙線研究の歴史
永野元彦, 大気の蛍光観測による宇宙線実験の始まり 棚橋五郎

14P ࠊPPࣆ ࢵ ࢳ ࡛ ษ ๐ ࡋ ࡓ ࠋ ࢸ ࣞ ࢫ ࢥ ࣮┤ ᚄ 㸪

㸬
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ග㏱ 㐣 ᛶ ⬟ ࡣ ➨  ᅗ ࡟ ♧ ࡍ ㏻ ࡾ࡛ ࠶ ࡿ

ࠋ

ࣞ ࣥࢬࡣ

First Detection of Shower by Fluorescence TechniqueѸ➨  ᅗ ࠉ 72.<2

ࡢ タ ⨨෗┿ࠋᖺ⛅ ࠊᮾ ி ࠿ ࡽ⣙ NPࡢ ᮾ

Candidates observed by TOKYO-1 (1969) ✦ Long signal duration and the

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similar amount of light (No.
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fluorescence-dominated shower ᆶwith 5×1018 eV, 680 g/cm2 (B.
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ࠉࠉࠉࠉࠉࠉࠉࠉࠉ '(/$< ࠉ7,0( 㸦SV 㸧✦ In the upgrade detector ofTOKYO-3, the 4 m2 lens was

Re-analysis by B. Dawson et al. (2011)➨  ᅗ ࠉ㑅 ࡤ ࢀ ࡓ ஦ ㇟ ࡢ ග ゅ ㏿ ᗘ

ࠋ

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୍

✦ Fly’s Eye experiment,

1NII-Electronicᕤ ᕤ (OHRWURQLR
Library ࠉService
ࠉ/LEUDU\

Fluorescence
dominated

Telescope Array experiment,
Pierre Auger Observatory
established fluorescence
technique and reported physics
results.

͂Ȣ 㹻

M. Ave et al. / Astroparticle Physics 28 (2007) 41–57

unfortunately UV protectedone.

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ࠉ ➨  ᅗࠉ ࣉ ࣛ ࢫ ࢳ ࢵ ࢡ ࣞ ࣥ ࢬ ࡢ ග ㏱ 㐣ᛶ ⬟72.<2 Ѹ㸪
72.<2 Ѹ㸪
72.<2 Ѹ 㸧ࠉ ࠉ ࠉ㸦

Airfly (2007)

x 10 2
1500

Counts

2

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1000

500

0
290

300

310

320

330

340

350

360

370

380

390

400

410

420

Wavelength (nm)Fig. 4. Measured fluorescence spectrum in dry air at 800 hPa and 293 K.

Table 1
Measured fluorescence band intensities in dry air at 800 hPa pressure and 293 K temperature

15

ࡉUHECR Observatories
Telescope Array
Experiment (TA)Utah, USA2
2
700 km (→ 3,000 km )~7 events/year (→ 30)

Pierre Auger
Observatory (Auger)
Malargue, Argentina
23,000 km

Google Earth~30 events/year

16Telescope Array Experiment (TA)

Largest cosmic ray detector in the Northern hemisphere2~ 700 km at Utah, USA
Fluorescence detector + Surface detector array

Fluorescence Detector at BRM and LR stations
Surface Detector Array2) + 256 Photomultiplier
Spherical
segment
mirror(6.8
m
507 Scintillator, 1.2 km spacing
tube(PMTs)/camera, 12 newly designed telescopes

Fluorescence detector
at MD station
Refurbishedfrom HiRes experiment,
2
Spherical mirror 5.2 m ,256 PMTs/camera,
14 telescopes

16×16
PMTs
PMT

17

Delta, Utah, USA

18

Pierre Auger Observatory (Auger)

The Pierre Auger Observatory
The Pierre Auger Observatory

2
The world’s largest UHECR observatory 3000 km
The Pierre Auger Observatory

13

array of 750 m,
io antennaInfill
array
array of 750 m,
Radio antenna array- )andcompletedin 2008
The Pierre(2004
Auger
Observatory
LIDARslaser
facilities
LIDARs and laser facilities
The Pierre Auger Observatory Surface Detector (SD)
Water Cherenkov Tank
Infill array of 750 m,
1.5km
spacing,
1600
LIDARs and laser facilities
Radio antenna array
ay of 750 m,
1665stations
surface detectors:
LIDARsand laser facilities
ntennaarray
water-Cherenkov
ion
1665tanks
surface detectors:(grid of
1.5 km,3000 km )es High elevation
Fluorescence
Detector (FD)water-Cherenkov
tanks
telescopes(grid of 1.5 km, 3000 km )3.4 m spherical mirror, 440 PMT, 30° × 30° FOV
light guide + collector ring, 4 × 6 telescope
(a)

2

2

High elevation
telescopes

1665 surface detectors:
water-Cherenkov tanks
4 fluorescence detectors(grid of 1.5 km, 3000 km2)
(24 telescopes
in total)
4 fluorescence
detectors(24 telescopes in total)

Figure 3.2: (a) Schematic depiction of a surface detector station [28]; (b) a surface
detector station deployed in the field.

tubes (PMTs) are optically coupled to the water and symmetrically positioned on
top of the tank with a distance of 1.2 m between each other. Each detector is de1665vised
surfaceto work detectors:
completely stand-alone, thus, every tank is equipped with a battery
box and a solar powersystem providing the 10 W average power required for thewater-Cherenkov
tanks
electronics [29]. A GPS (Global
Positioning System) unit is installed at each2)
(grid tank
of
1.5
km,
3000
km
tank as a basis for time synchronization between the detector and the Central Data
Acquisition System (CDAS) as well as for providing precise information about the
tank’s position. The communication between the detector and the CDAS is achieved
wirelessly via one of the four communication beacons located near the FD sites atthe perimeter of the array.

Southern hemisphere:
Province Mendoza,
Argentina
20
Southernhemisphere:
Province Mendoza, Argentina 20

4 fluorescence detectors(b)

4 fluorescence detectors(24 telescopes in total)

To detect charged particles from extensive air showers, the Cherenkov e↵ect is exploited [30]. When the velocity of a charged particle traversing a medium is greaterthan the speed of light in this medium, Cherenkov light is emitted by this particle ina cone along its trajectory. The Cherenkov light produced in the tank by secondary
particles from extensive air showers, mostly muons and electrons, is detected by the
PMTs mounted on top of the tank and converted into a current pulse. To increase
the amount of Cherenkov light collected, a reflective layer of high-density polyethylene fabric covers the inside of the tank. The signals from the PMTs are read out
and digitized by Flash Analog-to-Digital Converters (FADCs) at a rate of 40 million

19

Malargüe, Mendoza,
Argentina

20

“Scale” of UHECR Observatory

21

“Scale” of UHECR Observatory

22

“Scale” of UHECR Observatory

23

My Contributions to TA
FD Camera Assembling @Akeno
Absolute calibration of FD PMT @ICRR
Detector response simulation with
GEANT4
Operation and maintenance of TA FD
@Utah

TA

TA FD Analysis software development

AGASA

!

,
(x, y, z)=(0 m, 0 m, 10 m)

10 m

,
.

3.6: Geant4

TA

Auger
3.2: Geant4

AGASA

34

24

Energy Spectrum by TA FD
Reconstruction by one FD station to be sensitive to lower energies.
R.U. Abbasi et al. / Astroparticle Physics 80 (2016) 131–140
17.2
Energy spectrum covering three orders of magnitude above 10 eV

25
20
15
10
5
0

Number of Photo-Electrons

350
300
250

40

50
60
70
Azimuth angle [degree]

Data
Fluorescence
Direct Cherenkov
Rayleigh Scatterd C.
Mie Scattered C.

80

χ2/ndf = 235.0/329 (0.7)

2

#p.e.
3610
3240
2890
2560
2250
1960
1690
1440
1210
1000
810
640
490
360
250
160
90
40
10
0

1

24

30

Time(µs)
48
45
43
40
38
35
33
30
28
25
22
20
17
15
12
10
7
5
2
0

TA FD (this work)

IceTop-73

3

Elevation angle [degree]

35

Flux × E /10 (eV m-2 sr-1 s-1)

10

Systematic Uncert.

KASCADE-Grande

TA MD

HiRes-I

TA SD

HiRes-II
Auger ICRC 2015

10-1
16.5

200
150
100
50
0
500

600
700
800
2
Slant Depth (g/cm )

900

17

17.5

18

18.5
19
log10(E (eV))

19.5

20

20.5

T. Fujii et al., ICRC2015, R. Abbasi, T. Fujii (Corresponding author) et al.,Astroparticle Physics 80 (2016) 131-140

Fig. 14. Energy spectrum compared with results reported by IceTop-73 [36], KASCADE-Grande [37], HiRes [27], Auger [38] and other detectors within

25

10

8.5% difference on energy
scale between Auger and TA.

3
E J(E)

Preliminary

Auger

TA

Eankle
(EeV)

4.8

5.2

Es
(EeV)

42.1

63.0

1

3

24

2

Flux × E /10 (eV m-2 sr-1 s-1)

How frequent?: Energy Spectrum

TA ICRC 2015
Auger ICRC 2015
-11017.5

18

18.5
19
19.5
log (E (eV))

20

20.5

10

The energy spectrum around ankle are in good agreement,
but an energy of suppression is different.

26

Pioneering measurements by TASD and Auger water tank

10
Preliminary

2014/Nov/28 15:43:01.453549
Event display with SD waveforms

FADC count

Flux × E3/1024 (eV2 m-2 sr-1 s-1)

Detector Exchange to Understand the Discrepancy
Energy: 3.04 EeV, zenith: 46.8 deg (Preliminary)
1
DET
1219

TA Combined 2015
Auger ICRC 2013 +8.5%-11017.5
18.5Observatory
19 and19.5
20 Array20.5
Surface Detectors of the18
Pierre Auger
the Telescopelog (E (eV))

DET1218
Upper
Lower

Showerdiretion

Bin number (each 20ns)

Water
tank

DET
1317

DET1318

10

DET
1118

Practical implementation
4.2. THE SCINTILLATOR DETECTOR

69

60

R.
Takeishietal.,
ICRC2015,
he two Auger SD stations deployed at the TA Central Laser Facility.

AugerPrime

Figure 4.12: 3D view of the SSD module with the support bars. The bars are connected to the tankusing lifting lugs present in the tank structure.

4.2.7

Calibration and control system

CLF

CHAPTER 4. THE SURFACE DETECTOR

Figure 4.1: 3D view of a water-Cherenkov detector with a scintillator unit on top.

Simple and robust
construction of
detector moduleand mounting frame,
double roof forthermal insulation

Theta=46.763 E=3.037EeV Nhit=5 X=0.110km Y=1.333km

1527

AugerWhat kind of particle?: Mass Composition
down to 1017 eV using HEATIs the suppression GZKprocess or acceleration
limit?
Proton
Iron

15"

18.3
proton dominated at 10 eV.
19
σ(Xmax) is smaller than proton simulation above 10 eV.
18.5
Suggest a change of composition above ~10 eV with increasing mass number andsmall mixing.

Pierre Auger Collaboration PRD 90 122005 (2014), A. Porcelli in ICRC 2015 28

Interpretation of Xmax distribution
N fraction
He fraction
p fraction

The same <Xmax> andσ(Xmax) but different
mixture.

Fe fraction

2.2. OPEN QUESTIONS AND GOALS OF UPGRADING THE OBSERVATORY
1
0.8
0.6
0.4
0.2
0
1
0.8
0.6
0.4
0.2
0
1
0.8
0.6
0.4
0.2
0
1
0.8
0.6
0.4
0.2
0

Sibyll 2.1
QGSJET II-4
EPOS-LHC

10 18

11

Pierre Auger Collaboration

E [eV]

10 19

10 20

Intermediate composition, such as helium or nitrogen

Figure 2.9: Estimate of the composition of ultra-high energy cosmic rays at the top of the atmoIronThedisfavoredin all energy
range
sphere [23].
Xmax distributions
measuredwith the Auger Observatory have been fitted by asuperposition of four mass groups accounting for detector resolution
and acceptance effects. The
19.7
No
mass
composition
information
above
10
eV
error bars show the combined statistical and systematic uncertainties of the mass estimates, except
those related
to the choice
the hadronicinteractionmodels
models.
or problem
in ofhadronicinteraction

The Pierre Auger Collaboration, Phys. Rev. D 90, 122006 (2014) 29

N

improved, the quality of the fits is still poor over mostare shown in the lower panel of the figures, with error bars
Feof variation obtained withinP, N,
Feenergy range for theP,otherHe,Fe p-values forof thetwoN,models.
corresponding to the P,
range
all models are good for events with energy above 1019.2 eV.
the systematic uncertainties. Where p-values are less thanFe
Sibyll 2.1
Sibyll 2.1
Sibyll 2.1−4
10 , they
When helium nuclei are also included, we findthat theare indicated with downward arrows.
50
NHe
For the simple mixture of protons and iron nuclei (Fig. 2),data are well described by all models withinp systematicAuger
40 second-highest energy bin (E ¼ 1019.4–19.5 eV)
only theuncertainties over most of the energy range(Fig. 4).
log(E/eV)==log(E/eV) =To aid in the discussion,
the X max distributions of theyields goodlog(E/eV)
fit qualities
for all three hadronic log(E/eV)
interaction
19.0-19.1
19.0-19.1
19.0-19.1
30=19.0-19.1
models. However, the fit qualities for-4all three models arefits are displayed for the energy bins E ¼ 1017.8–17.9 eV
p = 0.004
p = 0.605
p < 1019.0–19.1
generally poor throughout the energy range, even when the(Fig. 5), E ¼ 10
eV (Fig. 6), and E > 1019.5
20
systematic uncertainties are taken into consideration.
(Fig. 7), respectively. Each figure contains nine panels
In order to determine whether there is any composition
that
cover
the
speciescombinationandhadronic
interaction
10
mixture where the models result in an adequate
models used. The contributions of all species are stacked
0
500
20
50

NN

log(E/eV)
>19.5

15
40

600

700

800

900

1000

500

600

700

800

900

1000

500

600

700

Sibyll2.1 II-04
QGSJET

Sibyll2.1II-04
QGSJET

Sibyll 2.1II-04
QGSJETlog(E/eV) > 19.5
log(E/eV) =log(E/eV) > 19.5
log(E/eV) =log(E/eV) > 19.5
log(E/eV) =30
1019.0-19.1

19.0-19.1

p = 0.001
p < 10-4

800

1000

Fe
N
He
p
Auger

19.0-19.1

p = 0.504
p = 0.001

900

p = 0.592
p = 0.064

20
5
10

0
0
500
500
20

600
600

700
700

800
800

900
900

1000
1000

500
500

600
600

700
700

800
800

900
900

1000
1000

500
500

600
600

700
700

800
800

900
900

1000
1000

The Pierre Auger Collaboration, Phys. Rev. D 90, 122006 (2014) 30

proton (red solid line), pure iron (blue dashed line), and a equal mixture of both (pink dash-dotted line).

Mass composition analysis in TA

800

Data (Jan/2008-Dec/2014)
19 g/cm2 sys. uncert.

SIBYLL 2.1
QGSJetII-04

Proton

EPOS-LHC
20

750
700

36 1171
8046952 72115

6
751 314 2
3

16

HiRes PRL’10

39
74

10

97

800
750
700

Iron

650

650

Telescope Array ICRC15 Preliminary

18

TA (this work)
Auger PRD’14

850<Xmax> (g/cm )<Xmax> (g/cm2)

850

QGSJetII-03
QGSJet01

900

2

900

18.5

19
log10(E (eV))

19.5

20

Telescope Array ICRC15 Preliminary

18

18.5

19
log10(E (eV))

19.5

20

Apply
the
Xmax
detectionbias-free
cut
like
Auger
improved
method.
(a) Obtained <X > compared with MC predictions(b) <X > compared with other results
max

max

Use
the
monocular
analysistomaximizethestatistics.
Figure 6: (a) Evolution of the average X
with energy from data (black points), plotted with the MCmax

T. Fujii et al., ICRC 201518
simulation
results(lines)fortwoparticle
species(proton,Consistent with the proton prediction above 10 iron)eV. and five hadronic interaction models. The
box region shows the systematic uncertainty on Xmax , 19 g/cm2 , for the monocular analysis. (b) Average
31
Auger, HiRes
results
showbya HiRes
good [18]agreement
withinXTA,
with results
reportedand Auger[19]. the systematic uncertainty.
max comparedSearch for Ultra-relativistic Magnetic Monopoles with the Pierre Auger Observatory

1010

Energy Loss (GeV/(g/cm2))

Supermassive magnetic monopoles10(M ≈ 1026 eV/c2) could be
10
produced in the early Universe. 10
10

9

Total

8

Ionization Loss

7

Pair Production

Photo Nuclear Effect6

10Intermediate mass monopoles (IMMs,
M∼1011-1020 eV/c2) can10
25
reach kinetic energies of ~10 eV by
acceleration in the galactic
10
and intergalactic magnetic fields. 10
5
4

3

2

10

1

Search for the signal produced by ultra-relativistic magnetic
monopole with 10-years Auger FD data. (a) Energy loss
4

6
8
Lorentz Factor, log 10(γ )

10

25

800

Monopole 10

700
600

Proton 1020 eVeV, γ =1011

500
400
300
200
100
0
0

200 400 600 800 10001200 1400 1600

12

Flux Upper Bound [(cm 2 s sr)-1]

2

Energy deposit (PeV/(g/cm2))

Search for Ultra-relativistic Magnetic Monopole

Slant depth (g/cm2)

10-11
-12

10

PARKER(b) IMM profile
IceCubeANITA-II

SLIM

RICEfactor γ . (b)Auger
Figure 1: (a) Energy loss of a magnetic monopole inLongitudinal
10air as a function of its Lorentz
MACRO
profile of the energy deposited by an ultra-relativistic10-14IMM of Emon = 1025 eV, γ = 1011 and zenith angle of2070◦ (red solid line). The profile of an UHECR proton
10-15 shower of energy 10 eV is shown as a black solid
line.
10-16

41

-13

10-17
10-18
10-19

-20
10
become the main cause of energy loss. Bremsstrahlung is highly suppressed by the large monopole

10

-21

Auger 2015 Preliminary

mass. An ultra-relativistic IMM would deposit a-22large amount of energy in its passage through the
10 6
7 example,
89 IMM
10 with11γ = 101211 loses
1343
Earth’s atmosphere, comparable to that of an UHECR. Foranlog(a )
FIG. 3:Collaboration,
Reconstructed signals for a simulated
magnetic
monopole of energy 10 eV and γ = 10 .
T. Fujii for the Pierre Auger
ICRC2015322 ) (cf.
44≈ 400 PeV/(g/cm
Figure 1(a)), which sums up to ≈ 1020.5 eV when
integrated over an atmo42

25

11

Where come from?: Arrival direction
COMPOSITION / MAGNETIC FIELDS

Magnetic Field (Gauss)

SOURCESGZK Cutoff → Nearby Universe → Large scale structureDeflections

High-energy astrophysics~ Z/E

UHECR → Smaller deflection in galactic/extragalactic magnetic field> Three messengers are available to study the non-thermal universe.

Anisotropyof
Size
Charged particles:
Hillas plot
Photonsp, Z, e±UHECR expected
Extra-galactic?
AGNs? GRBs? Neutrinos
↓
…?
UHECR
(Cronin TAUP 2013)
Interaction with CMB → Sources within ~ 100 Mpc
Astronomy
Propagation distance (Mpc)
Expect flux suppression and anisotropy of light component
Mean Energy (eV)GZK HORIZON

Image credit : M. Ackermann @ TAUP2013

33

Where come from?: Arrival direction
All Sky Survey with TA&PAO

TA Hotspot: 5.1σ (pre-trial)
Oversampling with 20°-radius circle
3.4σ(post-trial)
Oversampling with 20°-radius circle

Northern TA : 7 years 109 events (>57EeV)
19.75
Northern
years157
109events
events(>10eV)
Southern
AugerTA:: 107years(>57EeV)No correction forAuger Warmspot:
3.4σ(pre-trial)
E scaledifference
19.75
b/w TA and PAO !!
Southern
Auger:
10
years
157
events(>10eV)
Southern
hotspot
is
seen
at
Cen
A(Pre-trial~3.6σ)
K. Kawata et al., ICRC 2015

34

80
3
2

60

1
0
-1

40
20

-2
-3
10

0

15
Azimuth Number-3

Flux × E3/1024 (eV2 m-2 sr-1 s-1)

Auger

-2

UHECR Astronomy?

10

down to 1017 eV using HEAT

Preliminary
Aab et al.

Mass
Composition

1
TA ICRC 2015
-1

0

1

Auger ICRC 2015

2

3
alm/σlm-11017.518the right 18.5
19line is a 19.5
estimated multipole moments (in equatorial coordinates). In
panel, the blacklog (E (eV))

No anisotropy seen

10

20

20.5

Hotspot/Warmspot
All Sky
Survey with TA&PAO
Oversampling with 20°-radius circle

Equatorial Coordinates - 15° smoothing
90°

0.4

3

60°

2
30°

0.38

1
0.36

180°

120°

0.34

0.32

ApJ 794 172 (2014)

60°

0°

-30°

-60°
-90°

300°

240°

0

-1

-2

-3

Northern TA : 7 years 109 events (>57EeV)

35

Origin and Nature of Ultra-High Energy
Cosmic Rays (UHECRs) and Particle
Interactions at the Highest Energies

Flux × E3/1024 (eV2 m-2 sr-1 s-1)

Physics Goal of UHECR Astrophysics
10
Preliminary

1
TA ICRC 2015

A.Porcelli,2015 420
AugerAuger ICRC
-11017.518101718.5
19 HEAT
19.5
down toeV usinglog (E (eV))

20

20.5

10

How frequent?: Energy Spectrum
What kind of particle?: Mass Composition

All Sky Survey with TA&PAO
Oversampling with 20°-radius circle

Where come from?: Arrival Direction
36

Recent Results and New Puzzle
19.5
Precise observation of the flux suppression above 10 eV.

Gradually increase heavier composition above ankle.
Hotspot/Warmspot of UHECRs
Flux suppression due to GZK process or maximum energy ofaccelerator?
Heavier composition or hadron interaction model? Proton fraction?
19.7
Mass composition above 10 eV?
Anisotropy as indication of additional light component?
Particle physics extrapolation at the highest energies?
37

TA×4

Detailed measurement of the hotspot

Expected Result in 2020

Enlarge the fourfold coverage to TA×4 = Auger.

1 cluster (MC)2014/10/13

SD part was
funded inJapan, 2015.

H. Sagawa@UHECR201430

2 cluster (MC)H. Sagawa ICRC2015

38

µ
N max

AugerPrime

vs Xmax

Figure 4.12: 3D view of the SSD module with the support bars. The bars are connected to the tankusing lifting lugs present in the tank structure.

8.6 QGSJet II.04

E = 5·10 eV
p QGSjetII.04
19

lg (N )µ

8.5

Primary comic ray identification for next 10 years.
The SSD calibration is based on the signal of a minimum ionizing particle going through the
The scintillator units have to be precisely calibrated with a technique similar to the caldetector, a MIP. Since this is a thin detector, the MIP will not necessarily be well separated
ibration
the WCD
from the low energy background but, being installed
on top procedureof the WCD,ofa cross
trigger(cf. section 4.2.7). The size of the detector and its intrinsic
accuracy
shouldnot be the dominant limitations for the measurement. The
can be used to remove all of the background. Aboutmeasurement
40% of the calibration
triggersof the
WCD produce a MIP in the SSD. The statistics of calibration
eventsrecordedin aunits
minute,theto be adequate to guarantee the physics goals of the prodynamic
rangeof thehas
LeftModule
RightModule
normal WCD calibration period,
are therefore enoughposed
to obtaina precise measurementof theupgrade.
2
MIP. Figure 4.13 shows the MIP calibration histogram from a 2 m test module, obtained in
The detector will be assembled and tested in parallel in multiple assembly facilities to
one minute of acquisition. The MIP is clearly defined, and will allow an absolute calibration
reducethe production time and, therefore, has to be easily transportable. The mechanicalof the SSD to better than 5%.
robustness
the scintillator
units must be ensured. The units will be shipped after assemThe performance requirements for the SSD come mainly from ofcalibration
requirements:
10 cm
10cm
bly, andvalidated
Malargin shower measurement mode, the dominant measurement
errorsare due atto thePoissonfluc-üe facilities of the Pierre Auger Observatory before being
tuations of the number of particles detected, and thetransported
overall calibration
constant
to their
final determidestination on top of a WCD in the Pampa. They will then havenation. Detector non-uniformity contributes a smalltoerror
when forcomparedto theoperate10
years
inaPoissonhostile environment, with strong winds and daily temperature2
error, as long as non-uniformities are below 20%. While the FWHM of the WCD calibration
variations of up to 30 C.
histogram will be clearly smaller than that of the SSD (the calibration unit for the WCD, theVEM, is at about 100 pe), the fact that the SSD can be cross-triggered by the WCD meansthat the MIP is clearly visible against very little background. The width of the MIP distri4.2.2 Detector design
bution is mostly determined by Poisson statistics of the number of photoelectrons per MIP,
the non-uniformity of the detector, and the intrinsic The
fluctuationof theresponse
to onlya singlebaseline
designrelieson existing technology for which performance measurements
particle, mainly due to different track lengths in the scintillator. The latter factor was deterhave
been
made.
The
Surface
Scintillator
Detectors (SSD) basic unit consists of two modules
mined from simulations to be around 18%. The baseline design chosen for the SSD producesof to⇡ 8 photoelectrons
2 m2 extruded
plastic
scintillator
which are read out by wavelength-shifting (WLS)
12 photoelectrons per MIP [146], which would degrade
after
10 yearsoffibers coupleda single photo-detector.
Extruded scintillator bars read by wavelengthoperation due to aging. This amounts to a 35% contribution
to the MIPtodistribution
width.

N QGSjetII.04
Fe QGSjetII.04

N

Both W
conne
electroThe key is enhancement electromagnetic/muon
separation of SD to measure the mass composition.
Two modules in one box per station,
readout by one PMT, area ~4 m

He
p

8.4

Read-out of scintillatorswith WLS fibers

1cm

log10(Nµmax)

Fe

187cm187

124cmHe QGSjetII.04

Figure 4.1: 3D view of a water-Cherenkov detector with a scintillator unit on top.

4.2.7 Calibration and control systemshifting fibers have already been employed in the MINOS detector [143]. The active part ofeach module is a scintillator plane made by 12 bars 1.6 m long of extruded polystyrene scintillator. Each bar is 1 cm thick and 10 cm wide. The scintillator chosen for the baseline designis produced by the extrusion line of the Fermi National Accelerator Laboratory (FNAL) [144].
The bars are co-extruded with a TiO2 outer layer for reflectivity and have four holes inwhich the wavelength-shifting fibers can be inserted. The fibers are positioned following thegrooves of the routers at both ends, in a “U” configuration that maximizes light yield andallows the use of a single photomultiplier (at the cost of a widening of the time responseof the detector by 5 ns, which has a totally negligible impact). The fibers are therefore read

Boost in statistics by a factor of ~ 10.

Proton and iron separation at 100 EeV for QGSJetII-03 simulated events using simulated true
Xmax (g/cm2)
Sensitive
to a proton
contribution
as low as 10%.
), reconstructed values (center), and taking
into accountenergyresolution (right).
600

700

800
Xmax [g/cm2]

900

1000

Muons may even outperform
Xmax
2
Scintillator
detector to measure the mass composition by SD.
Install
4m
Scintillator
at highest energies !
70

Fibers routing– University Wuppertal

15

UHECR2014, Springdale (Utah), Oct. 2014

PMT/SiPM

Proton QGSJet
Iron QGSJet

60

Iron

50
40

30 EeV, MF = 1.28
30

Proton

20
10
0
-370

Sunroof

-360

-350

-340

-330

Discriminant [a.u.]

-320

-310

39

-300

Extended
FD operation
ExtendedFD
Operation
Clear sky, no moonlight

40 times higher NSB (90% moon)

15% duty cycle

Increase by 50%
by measurement
during high nightsky backgroundE = 7x1019 eV

10x reduced PMT gain by
reducing supplied HV.
Successful test has been
done last year.

Radomír Šmída – AugerPrime

11

40

41

Lateral segmentation and number of phototubes

is extending beyond the m-metal tube by 6 ð1Þ cm
at the front (rear) end.
Tests were conducted with such assembly in amagnetic field [18]. The phototube was kept at alow relative gain (See Section 5.1.1 for thedefinition) of 5 % 104 to be sensitive to themagnetic field. It was found that there was littlechange ðo1%Þ in the response up to 170 ð220Þ G
for the field parallel (perpendicular) to the axis of
the phototube. Therefore, it was concluded thatthe magnetic shield would be quite adequate forthe measured level of the stray field.

Reuse CDF PMTs for AugerPrime
Lateral segmentation
NZ % Nf

Z%2:11 ðDf ¼ 7:5" ÞZ^2:11 ðDf ¼ 15" ÞEM
1:1%Z%3:49
8 % 484 % 24

HAD
1:32%Z%3:49
7 % 484 % 24

No. of Tubes
East or West
East + West

480
960

432
864

1900 CDF PMTs are available for AugerPrime.
Check the performance of 20 years-old PMT.
HAMAMATSU R4125-04 Phototube Assembly

Nevents

B a se1.4.3. Spectral matching of optical components
Magnetic Shield
Light Mixer
The scintillator for the tiles, Kuraray SCSN-38,(Iron Pipe)(Plexiglass)is a fast and efficient blue polystyrene scintillator
Anode
Signalwith a peak emission at 428 nm [14]. Green-light
Hamamatsu
DynodeR4125-04
emitting Y-11 was chosen for its excellent wave-H.V.
length shifting efficiency with its absorption band
Clear
Magnetic Shield
Fiber Mount
Fiberswell matched to the emission band of SCSN-38.
(µ Metal)
Also important is its emission band peaking at
Fig. 3. Sketch of the phototube assembly (not to scale).
476 nm with an extended tail beyond 500 nm
which gives a good transmission through rather
PMT R4125 Ser.No. # CX2267 at 1700V, DAC = 100
lowTable 5long, up to B3 m; clear fibers. The green emission
spec2
800
Entries49998 band is also important in ensuring tolerance
Dimensions of the components in PMT assembly
Mean
214.8
700
Iron pipe
3:6 mm wall; 35RMS
mm
ID; 214 mm L 223.9 against long-term exposure to radiation at these2
χ / ndf
365.7 / 147
forward angles. However, the green emission band
m-Metal
tube
1
mm
wall;
25:4
mm
ID;
140
mm
L
600
p0
1.34e+04±
2.71e+03
Light mixer
ð10:3 mmÞ2 ; 58 mm
p1 L−0.09481 ± 0.00527 of the WLS fiber severely constrains the type of
364.7 ± 4.5
PMT glass envelop
18:6 mm OD; 88a1mm L
500
photocathode of the phototubes. The response ofg1
263.9 ± 1.3
106.7 ± 1.3 an ordinary bialkali photocathode falls off rapidly
w1
400
a2
107.6 ± 2.6
g2555.8 ± 4.2 beyond B450 nm and the only practical choice is a300
1.4.2. Magnetic shieldso-called extended bialkali photocathode whichhas quantum efficiency of B12% at 480 nm: This
200 Since the steel plates of the HAD compartment
is a part of the flux return path of the solenoid,
was one of the major considerations for choosing
100
stray magnetic field on the rear surface of thethe phototubes for our purpose.
Endplug
structure where the phototube housings
0
0
100
200
300
400
500
600
700
800
900
1000
Chargeintegral[mV*ns]were to be mounted needed some attention.
1.5. Selection process
A measurement showed that the maximum of the

Muon
Tomography
at KIT

1 p.e.
=528 fC

1 MIP=4.93 pC
=9.33 p.e.

42

JEM-EUSO
R&D
on
radio
detection
Auger

Extreme Universe Space Observatory onboard Japanese Experiment Module

Pioneer detection of UHECRs from space
Auger Engineering Radio Arrayand other MHz and GHz
detectorsof muon detection
5 year Mission 1000 EventsUpgrade
E>60 EeV
capabilities(Auger 20/yr)

Pioneering UHECR Space MissionTA
Low energy
extension

A. Olinto, ICRC2015

43

Radar detectioPhysics Goal and Future Perspectives
Origin and Nature of Ultra-high Energy Cosmic Rays (UHECRs) andParticle Interactions at the Highest Energies
5 - 10 years
Detector
R&D
Exposure and Full Sky Coverage
Radio, SiPM,
TA×4 + Auger
Low-cost
JEM-EUSO : pioneer detection from Fluorescencespace and sizable increase of exposure Detector (FD)

“Precision” Measurements
AugerPrime
Low energy enhancement(Auger infill+HEAT+AMIGA,
TALE+TA-muon+NICHE)

10 - 20 years
Next Generation Observatories
In space (100×exposure): Super-EUSO
Ground (10×exposure with high quality events):44

Fluorescence detector Array of Single-pixel Telescopes

http://www.fast-project.org

45

Fluorescence detector Array of Single-pixel Telescopes✦ Target : > 1019.5 eV, ultra-high energy cosmic rays (UHECR) and neutral particles✦ Huge target volume ⇒ Fluorescence detector arrayFine pixelated camera

Too expensive to cover a huge area

Segmented mirror telescope
Variable angles of elevation – steps.

Single or few pixels and smaller optics
15 deg45 degLow-cost and simplified/optimized FD

46

Fluorescence detector Array of Single-pixel Telescopes✦ Each telescope: 4 PMTs, 30°×30°

20 km

field of view (FoV).

UHECRs✦ Reference design: 1 m2 aperture,

15°×15° FoV per PMT✦ Each station: 12 telescopes, 48 PMTs,
Photons at diaphragm

30°×360° FoV.
1

3

1

✦ Deploy on a triangle grid with 20 km

spacing, like “Surface Detector
Array”.

2

3

Photons at diaphragm

Photons at diaphragm

56 EeV✦ If 500 stations are installed, a groundzenith 500
22coverage is ~ 150,000 km .

✦ Geometry: Radio, SD, coincidence of
16

three stations being investigated.

47

FAST Expected Exposure
1.E+8✦ Conventional operation of FD

under 10~15% duty cycle

Exposures (L=km^2*sr*yr)

1.E+7

FAST

1.E+6

✦ Observation in moon night to

JEM-EUSO

achieve 30% duty cycle,

tilt

JEM-EUSO
nadir

1.E+5
Auger

✦ Target: >1019.5 eV✦ Target: >1019.8 eV = SuperGZK events

TAx4

TA

1.E+4

✦ Test operation by Auger FD
HiRes✦ Ground area of 150,000 km2 with

AGASA

1.E+3

Preliminary

Fly's Eye

1.E+2
1990

2000

2010

2020
Year

2030

2040

2
30% duty cycle = 45,000 km(15×Auger, cost ~75 Million USD)
✦ 450 events/year

48

Physics Target

All Sky Survey with TA&PAO
Oversampling with 20°-radius circle

Flux × E3/1024 (eV2 m-2 sr-1 s-1)

GZK Recovery

UHECR Anisotropy
with ~10x statistics

FAST

10
Preliminary

Upper detection
limits to the integrated
flux and neutrinos
Firstof UHEphoton
photons
No correction for
E scale difference
b/w TA and PAO !!

1

Auger ICRC 2015

18

18.5
19
19.5
log (E (eV))

20

20.5

10

-1

log10J(E)E ,m sec sr eV

2

4

24

3

-2

-1

1,2,3

E-2 spectrum-weighted
average exposure for Eg > E0
E0
[EeV]

·ꜫÒ
[km2 sr yr]

Fg (95% CL)[km-2 yr-1 sr-1]

10

5200

1.9 x 10-3

20

6800

1.0 x 10-3

40

6300

4.9 x 10-4

FAST

23

1
17

n

Nγ
F γ (E γ >E 0)=⟨ℇ⟩→ 01/01/04–20/06/13 no n candidate→ search not limited by background→ limit below the WB bound→ top-down (exotic) models strongly constrained→ cosmogenic
model
with
pure
p
composition4
at the sourceandstrong
FRII
evolution
disfavoured
3

25

18

19

20

2
21

→
01/01/04–15/05/13
log E,eV
V. Berezinsky et al., (2001)
10

12

PRD 91 (2015) 092008-11017.5Southern hotspot is seen at Cen A(Pre-trial ~3.6σ)

Feldman-Cousins limit
to the number of photons

Conclusions

TA ICRC 2015

Northern TA : 7 years 109 events (>57EeV)
Southern Auger : 10 years 157 events (>57EeV)

17

FAST

Coincidence of UHECR hotspot andgneutral particles

49

R&D for the FAST Project
Telescope Array site Black Rock Mesa station EUSO-TA telescope

FAST camera

EUSO prototype✦ Temporally use the EUSO-TA optics at the TA site.
✦ Two Fresnel lenses (+ 1 UV acrylic plate in front for protection)✦ 1 m2 aperture, 14°×14° FoV ≒ FAST reference design.
✦ Install FAST camera and DAQ system at EUSO-TA telescope.
✦ Milestones: Stable observation under large night sky backgrounds,

UHECR detection with external trigger from TAFD.

✦ 8 inch PMT(R5912-03,
Hamamtsu)✦ PMT base (E7694-01,

Hamamatsu, AC
coupling)✦ Ultra-violet band passfilter (MUG6, Schott)
50

1

2

3

4

5

6

7

8

9

Start observation

51

Np.e. /(100 ns)

Laser Signal to Check PerformanceHit
clouds

Efficiency

FAST
FOV

TAFD

Directional sensitivity
by RayTrace of
EUSO-TA telescope

Np.e. / (100 ns)✦ Vertical Ultra-Violet laser at 6 km from FAST ≒ ~1019.2 eV
350
300

FAST

Data

✦ Expected signal TAFD/FAST: (7 m2 aperture × 0.7 shadow

Simulation

× 0.9 mirror) / (1 m2 aperture × 0.43 optics efficiency) ~10

250
200

✦ TAFD Peak signal : ~3000 p.e. / 100 ns

150
100

✦ FAST Peak signal : ~300 p.e. / 100 ns. All shots are50

detected significantly.

0
0

100

200

300

400

500

600

700

800

Time (100 ns)✦ Agreement of signal shape with simulation.

52

Distance vs Energy (from TAFD) for Candidates
Np.e. / (100 ns)

102
Impact parameter [km]Preliminary
te
e
D

e
l
b
a
ct

log(E/eV)=18.0

80

60

40

10

20

0

-20

1
17

0

17.5

18

Almost! log(E/eV)=19.1

18.5
19log10(E (eV))

19.5

20

Figure 14: Distribution of the impact parameter as a function of the primary energy reconstructed by TA for shower candidates detected by the FAST prototype. The line indicatesthe maximum detectable distance by the FAST prototype (not fitted).

100

200

300

400

500

600

700

800

Time (100 ns)

FAST
FoV

53

R&D for the FAST Project✦ FAST prototype measurements at Utah
✦ Stable operation under high night sky backgrounds.
✦ UHECR detection.

EUSO-TA+ FAST camera
telescope

Accepted for publicationin Astroparticle Physics✦ Published in Astroparticle Physics 74 (2016) 64-72✦ Next milestones by new full-scale FAST prototype✦ Establish the FAST sensitivity.

FAST meeting in December 2015
(Olomouc, Czech Republic)✦ Detect a shower profile including Xmax with FAST54

Design of Full-scale FAST Prototype

Building height – elevation 15°
required dimensions

Building - ground plan – required dimensions
Cca 3500 mm✦ Adjustable elevation 15° or 45° to

Cca 3500 mm

enlarge the FoV of the TA FD.

FOV

Cca 3000 mm

Cca 3000 mm

Cca 1000 mm

FOV✦ Robust design for maintenance freeand stand-alone
observation.
FAST components

shutter – like sectional garrage doors

600 mmPossible solution of building

roof „window“
5

DUST and STRAY LIGHT protection

4

UV PMMA „window“in octagonal aperture

cover = black shroud

closedCca 3000 mm

open4000 mm

mirrorsopen

4 PMTscamera8 inchUV filter
glass
cablingclosedelectronics

5515

Full-scale FAST Prototype

We will install the full-scale FAST prototype at Utah in June/July 2016

56

0

2000

4000

6000 8000 10000 12000

−200
−100
0 800010000
100 12000200−2000
0 20004000 6000
1400016000−200

PMT Calibration System
Integrated counts

Integrated
countsIntegrated
countsPMT ZS0022

1200

1000

23504

Mean

2111

RMS

2691

r2 / ndf

175.6 / 85

Constant

184.8 ± 2.3

Mean4582 ± 30.6

Sigma

1933 ± 21.9

1200

Entries

1400

Entries

EfficiencyEntries[%]

Department of Physics, University of Chicago

160
22

Detection
Entries Efficiency: FAST
43157 PMTs

Entries
Mean

Mean

RMS

RMS

r2 / ndf

r2 / ndf

18. HV = 2169V, Disc = 38mV, x20 Amp

20
1000

22. HV = 2252V, Disc = 50mV, x20 Amp

140
18
16

24. HV = 2266V, Disc = 85mV, x20 Amp

Constant Constant
25. HV = 2000V, Disc = 44mV, x20 Amp
Mean
Mean

120
14800

Sigma

Sigma

23504
4210
0.857

6062
55.09
263.4
/ 104
0 / −3
161.5
± 2.1
184.8
± 1.4
9949
± 34.54582± 1.4

Entries

PMT ZS0018

Entries

Hamamatsu (Scaled)
PMT ZS0025 300
PMT ZS0024
PMT ZS0022
PMT ZS0018

Me

250

Co

RM

r2 /

Me

± 42.91933 ±2982
8350.8

12

Ent

Sig
200

100
10
800

8600

80
6
600

150

4

60
2400

Department of Physics, University of Chicago

0
0

400

40

100
100

200

300

400

500

600
700
Wavelength [nm]Non-Linearity: FAST PMTs
We measure NL as a function of anode current (IA ):
200
50
200
we do this by using two LEDs
that occupy both channels
FIG.2013.
Detection efficiency results with
Hamamatsu
mea
PMT
ZS00180of the FG. By adjusting pulse amplitude, width, and desurement for comparison.
PMT ZS0022
lay, we create two identical signals, S1 and S2 . We use
PMT ZS0024 0
0-2
0 −1005000−5010000
20000150
25000
−200
0
5000 10000 15000 20000 25000−200 −150
0 15000
50 100PMT ZS0025
internal triggering on the oscilloscope,and by Integratedchanging
Integrated
countscountsIntegrated
counts-4
the delay on the FG, we can add the individual signals
the powermeter and monochromator initialized, any reto create a combined signal denoted PMT
by S
ZS0024PMT ZS0025-6
12 . In theory,
maining lights
in 2*10^5the lab are switched o↵. The compute
HV
for
Gain
S12 should correspond to the sum of S1 and S2 . Howin the-8 lab is accessed remotely to begin data acquis
ever, as we increase the amplitude of the signals and the
ZS0018:program
1113 Vtion. ThePMT
DAQ
controls the monochromato
PMT ZS0022: 950 V
PMT becomes saturated, we observe a non-linear e↵ect:and -10powermeter.
It obtains and averages 10,000 read
PMT ZS0024: 1042 V
S12
hassmaller
amplitude
than
the
sum
of
S
and
S
.
1
2
ings -12from thepowermeterover 10 s for a given step; th
PMT ZS0025:
1068 V
FIG. 12. SPE peak height distribution used to set discrimiSample
signals canbe Theseenpedestal
in Fig. ends
10. NL
is definedas:nator thresholdvalue.
at arounda height
error, P , is calculated in quadrature from Poisson statisof 350 ADC counts. Dividing this by the 4095 dynamic range
tics on bothlamp
signalback
20 powermeter40
60 readings,
80
100
120 and140
S12 (S1threshold+ S2 ) of ⇡ 85 mV.
Anode Current57[mA]of the FADC gives a discriminatorground.
The
lamp
backgroundcorresponds
to
when
th
NL =
(3)
Non-Linearity [%]umber. We test thening a single photont. We place a singleof the dual-channel
the PMT. The LED
typical LED amplind ⇡ 100 ns.
is connected to ther; the two resulting
ADC input and first
ely. The PMT anodeonverts the signal to4095 (12-bit range).

Preparing for fullautomatic PMT
FIG. 7. Integrated
count
distributionofSPE
signals,
inclu
calibration system
ing pedestal (left peak) and SPE peak fitted to a Gaussi

Possible Application of FAST Prototype
Efficiency

1
Proton✦ Profilereconstruction with geometry given by SD (smearing

gaussian width of 1° in direction, 100 m in core location).
Iron

0.8
0.6

1. Introduction2
19.5✦ Energy: 10%, Xmax : 35 g/cm

at
10
eV
The hybrid detector of the Pierre Auger Observatory [1] consists of 1600

0.4
0.2

18
18.2 18.4 18.6cross-check18.8
1919.2 of Energy and Xmax scale between✦ IndependentAuger and TA

X max Resolution [g/cm2]

Energy Resolution [%]log(E(eV))

Energy20

Pierre Auger Observatory

0

Iron

18

18.2

88

1

T. Abu-Zayyad et al. / Nuclear Instruments and Methods in Physics Research A 689 (2012) 87–97

a r t i c l e i n f o

a b s t r a c t

Article history:
Received 25 December 2011
Received in revised form
25 May 2012
Accepted 25 May 2012
Available online 2 June 2012

The Telescope Array (TA) experiment, located in the western desert of Utah, USA,observation of extensive air showers from extremely high energy cosmic rays. The
surface detector array surrounded by three fluorescence detectors to enable simulta
shower particles at ground level and fluorescence photons along the shower trac
detectors and fluorescence detectors started full hybrid observation in March, 2008
describe the design and technical features of the TA surface detector.
& 2012 Elsevier B.V.

Keywords:
Ultra-high energy cosmic rays
Telescope Array experiment
Extensive air shower array

TA ICRC 2015

Auger ICRC 2015

1. Introduction

The main aim of the Telescope Array (TA) experiment [1] is to
explore the origin of ultra high energy cosmic rays (UHECR) using
their energy spectrum, composition and anisotropy. There are two
major methods of observation for detecting cosmic rays in theenergy region above 1017.5 eV. One method which was used at theHigh Resolution Fly’s Eye (HiRes) experiment is to detect air
fluorescence light along air shower track using fluorescence
detectors. The other method, adopted by the AGASA experiment,
is to detect air shower particles at ground level using surface
2
detectors deployed over a wide area ( ! 100 km ).
The AGASA experiment reported that there were 11 events
above 1020 eV in the energy spectrum [2,3]. However, theexistence of the GZK cutoff [4,5] was reported by the HiRes

18

18.5
19
19.5
log (E (eV))

Identical
simplified FD

Proton5Preliminary

10

15
10

10

surface stations – water Cherenkov tanks and their associated electronics – and24 air19.6
fluorescence telescopes. The Observatory is located outside the city of
19.4
Malargüe, Argentina (69◦ W, 35◦ S, 1400 m a.s.l.) and the detector layout isshown in Fig. 1. Details of the construction, deployment and maintenance of
the array of surface detectors are described elsewhere [2]. In this paper we will -1
1017.5concentrate on details of the fluorescence detector and its performance.

025

Flux × E3/1024 (eV2 m-2 sr-1 s-1)

✦ Install FAST at Auger and TA for a cross calibration.

18.4

18.6
18.8
log(E(eV))

19

100

19.2

19.4

experiment [6]. The Pierre Auger experimen
suppression on the cosmic ray flux at energy a[7] using an energy scale obtained by fluores
scopes (FD). The contradiction between results f
detectors and those from surface detector arrays
be investigated by having independent ener
both techniques. Hybrid observations with SD
us to compare both energy scales. Information aband impact timing from SD observation impro
reconstruction of FD observations. Observatio
detectors have a nearly 100% duty cycle, which
especially for studies of anisotropy. Correlations
directions of cosmic rays and astronomical objec
region should give a key to exploring the origin otheir propagation in the galactic magnetic field.

20

20.5

Telescope Array
Experiment

19.6

Xmax

80
60
40

Proton

20
0

Iron

18

18.2

18.4

18.6
18.8
log(E(eV))

19

19.2

19.4

19.6

Figure 1: Status of the Pierre Auger Observatory as of March 2009. Gray dots show the
Pierre Auger Collaboration, NIM-A (2010)
positions of surface detector stations, lighter gray shades indicate deployed detectors, while

Fig. 1. Layout of the Telescope Array in Utah, USA. Squares denote 507 SDs. There are three subarrays controlled by three communication towers den
three star symbols denote the FD stations.

Telescope Array Collaboration NIM-A (2012) 58

Energy spectrum: Precise observation of the flux suppression
19.5
above 10 eV.
Mass composition: Gradually increase heavier composition
above ankle.

Flux × E3/1024 (eV2 m-2 sr-1 s-1)

Summary

10
Preliminary

1

Auger

TA ICRCA.Porcelli,
2015
420
Auger ICRC 2015

-11017.5down
to181017 18.5
eV using
19 HEAT
19.5
20

20.5

log (E (eV))
10

Arrival direction: Hotspot/Warmspot of UHECRs

Flux suppression due to GZK process or maximum energyof accelerator?
Heavier composition or hadron interaction model? Proton
19.7
fraction? Mass composition above 10 eV?

All Sky Survey with TA&PAO
Oversampling with 20°-radius circle

Anisotropy as indication of additional light component?
Particle physics extrapolation at the highest energies?
Northern TA : 7 years 109 events (>57EeV)
Southern Auger : 10 years 157 events (>57EeV)

Future observatories: TA×4, AugerPrime, JEM-EUSO and FAST
No correction for
E scale difference
b/w TA and PAO !!

59

Southern hotspot is seen at Cen A(Pre-trial ~3.6σ)
12