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TAIGA the Tunka Advanced Instrument for cosmic ray physics and Gamma Astronomy —
present status and perspectives.
View the table of contents for this issue, or go to the journal homepage for more
2014 JINST 9 C09021
(http://iopscience.iop.org/1748-0221/9/09/C09021)
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2014 JINST 9 C09021
PUBLISHED BY IOP PUBLISHING FOR SI SS A MEDIALAB
RECEIVED:June 12, 2014
ACC EPTED :July 12, 2014
PUBLISHED:September 18, 2014
INTERNATIONAL CONFERENCE ON INSTRUMENTATION FOR COLLIDING BEAM PHYSICS
BUDKER INSTITUTE OF NUC LE AR PHYSICS, NOVOSIBIRSK, RUSSIA
FEB RUA RY 24 – M ARCH 1, 2014
TAIGA the Tunka Advanced Instrument for cosmic
ray physics and Gamma Astronomy — present status
and perspectives.
The TAIGA collaboration
N.M. Budnev,bI.I. Astapov,iA.G. Bogdanov,iV. Boreyko, jM. B ¨uker,fM. Br ¨uckner,k
A. Chiavassa,dA.V. Gafarov,bO.B. Chvalaev,bN. Gorbunov, jV. Grebenyuk,j
A. Grinyuk,jO.A. Gress,bT. Gress,bA.N. Dyachok,bS.N. Epimakhov, fD. Horns,f
A.L. Ivanova,bN.I. Karpov,aN.N. Kalmykov,aY.A. Kazarina,bV. Kindin,i
N.V. Kirichkov,bS.N. Kiryuhin,bR.P. Kokoulin,iK.G. Kompaniets,i
E.N. Konstantinov,bA.V. Korobchenko,bE.E. Korosteleva,aV.A. Kozhin,a
M. Kunnas,fL.A. Kuzmichev,aV.V. Lenok,bB.K. Lubsandorzhiev,c
N.B. Lubsandorzhiev,aR.R. Mirgazov,bR. Mirzoyan,e,bR.D. Monkhoev,b
R. Nachtigall,fA.L. Pakhorukov,bM.I. Panasyuk,aL.V. Pankov,bV.A. Poleschuk,b
E.G. Popova,aA. Porelli,hV.V. Prosin,aV.S. Ptuskin,gA.A. Petrukhin,iG.I. Rubtsov,c
M. Rueger,k,hV.S. Samoliga,bP.S. Satunin,gV.Yu. Savinov,bYu.A. Semeney,b
B.A. Shaibonov Jr.,cA.A. Silaev,aA.A. Silaev Jr.,aA.V. Skurikhin,aM. Slunecka,j
C. Spiering,hL.G. Sveshnikova,aA. Tkachenko, jL. Tkachev,jM. Tluczykont, f
R. Wischnewski,hI.I. Yashin,iA.V. Zagorodnikovband V.L Zurbanovb
aSkobeltsyn Institute of Nuclear Physics MSU, Moscow, Russia
bInstitute of Applied Physics ISU, Irkutsk, Russia
cInstitute for Nuclear Research of RAN, Moscow, Russia
dDipartimento di Fisica Generale Universiteta di Torino and INFN, Torino, Italy
eMax-Planck-Institute for Physics, Munich, Germany
fInstitut fur Experimentalphysik, University of Hamburg, Germany
gIZMIRAN, Troitsk, Moscow Region, Russia
hDESY, Zeuthen, Germany
iMoscow Research Nuclear University MEPhI, Moscow, Russia
jJoint Institute for Nuclear Research, Dubna, Russia
kInstitute for Computer Science, Humboldt-University Berlin, Germany
E-mail: budnev@api.isu.ru
c
2014 IOP Publishing Ltd and Sissa Medialab srl doi:10.1088/1748-0221/9/09/C09021
2014 JINST 9 C09021
ABS TR ACT: TAIGA stands for “Tunka Advanced Instrument for cosmic ray physics and Gamma
Astronomy” and is a project to built a complex, hybrid detector system for ground-based gamma-
ray astronomy from a few TeV to several PeV, and for cosmic ray studies from 100 TeV to 1 EeV.
TAIGA will search for “PeVatrons” (ultra-high energy gamma-ray sources) and measure the com-
position and spectrum of cosmic rays in the knee region (100 TeV–10 PeV) with good energy res-
olution and high statistics. TAIGA will include Tunka-HiSCORE — an array of wide-angle air
Cherenkov stations, an array of Imaging Atmospheric Cherenkov Telescopes, an array of particle
detectors, both on the surface and underground and the TUNKA-133 air Cherenkov array.
KEYWORDS: Large detector systems for particle and astroparticle physics; Gamma telescopes
2014 JINST 9 C09021
Contents
1 Introduction 1
2 Gamma ray observatory TAIGA 2
2.1 Tunka-133 3
2.2 Tunka-HiSCORE array 4
2.3 Tunka-IACT array 7
2.4 A net of Tunka particle detectors 8
3 Conclusion 9
1 Introduction
In recent years ground-based very high energy (VHE) gamma-ray astronomy became the most dy-
namically developing field in high energy astrophysics. Advantages of gamma-rays as information
carriers from the powerful Galactic and extragalactic sources is due to the fact that in contrary to
charged cosmic rays, they preserve their emission direction from sources and unlike neutrinos, they
are easy to detect. Gamma-radiation of ≥100 GeV was detected from more than 150 sources of
different types but gamma-rays with energies higher than 10 TeV are detected only from 10 sources.
Above 100 TeV there are no detected sources of gamma-rays. One can say that the latter energy
range remains Terra incognita for present day gamma-astronomy, albeit it is of utmost importance
for resolving many fundamental problems of astroparticle physics.
So far most of the knowledge about sources of high energy gamma-rays has been ob-
tained from arrays of Imaging Atmospheric Cherenkov Telescopes (IACT), namely: HEGRA [1],
H.E.S.S. [2], MAGIC [3], VERITAS [4] etc. An IACT consists of a tessellated mirror and a cam-
era in its focal plane, which Cherenkov pulses from showers are detected. The method of image
shape analysis proposed by A.M. Hillas [5] in 1985, allowing to select with high efficiency EAS
produced by gamma-rays from those initiated by charged cosmic rays. Effective detection area of
a single IACT is about 0.5×105sq. m. The so-called 4th generation IACT array CTA [6] will have
substantially higher sensitivity than the existing telescopes in the energy range of up to 100TeV.
Due to limited shower collection area of the arrays (0.1–1 km2) it is hardly possible to go into the
energy range considerably higher than 100 TeV. For extension to the higher gamma-ray energies
with correspondingly lower fluxes it is necessary to construct arrays with lots of telescopes, dis-
tributed over an area of a few square kilometres at least. The estimated cost of such an IACT array
can be more than 100 million US dollars per square kilometer. So, it is necessary to find other,
more cost-effective experimental solutions which would allow one to have an array of an area of
the order of 10–100 km2for studying PeVatrons.
A very interesting approach expanding possibilities of traditional gamma IACT in high energy
range could be inclusion of imaging gamma telescopes with wide-angle Cherenkov detectors in one
installation. A wide-angle Cherenkov array allows one to reconstruct the arrival direction of EAS
–1–
2014 JINST 9 C09021
with high precision (to 0.1 degrees), the axis position of the EAS, energy and shower maximum
height. Preliminary calculations show that this information, combined with the shape of the EAS
image can significantly suppresses the background from cosmic ray events. Common operation
of a telescope and a wide-angle array can allow one to abstain from simultaneous measurement
of the EAS by several closely packed telescopes and thus allows to increase the distance between
telescopes up to 600 m. Along with substantially low cost of such an array of largely spaced
telescopes combined with a wide-angle array (cost of a single station of the wide-angle array makes
only ∼1/30 of a cost of an imaging telescope), it can provide the necessary huge collection area
for preparing measurements at ∼100 TeV.
In the energy range above 100 TeV to use of large area muon detectors is a very important
method for the separation of gamma-quanta induced EAS from cosmic rays induced showers. The
possibility to separate gamma-quanta induced showers is based on the fact that the number of
muons in cosmic ray showers is 30 times more than for gamma-quanta showers. Because of the
relatively low number of muons in EAS at such energies (the number of muons EAS from protons
with energy of 100 TeV is 1000) the total area of muon detectors should be 0.2–0.3% of array total
area and be 2000–3000 m2. Thus, cost reduction per unit of muon detector area is a very important
issue. To solve many crucial tasks of high energy gamma-astronomy we propose to design and
to construct the gamma-observatory TAIGA (Tunka Advanced Instrument for cosmic ray physics
and Gamma Astronomy) — a complex, hybrid detector system. TAIGA will include four different
types of detectors:
1. The HiSCORE wide Field of View (FoV; ∼0.6sr,) integrating air Cherenkov detector sta-
tions, placed at distances of 150–200 m from each other, covering an area of initially ∼1 km2
and >10 km2at a later phase of the experiment. These detectors consist of four bialkali
PMTs of size either 20 cm or 25 cm in diameter, located next to each other. The PMTs have
Winston-cone shape light-guides, which increase their light collection area by factor of four.
The PMT outputs are summed up.
2. The TUNKA-133 wide FoV integrating air Cherenkov detector. This full size detector, op-
erational since 2009 with 175 stations, is spread over an area of ∼3 km2. The individual
stations consist of single large PMT (diameter 20 cm).
3. Up to 10 IACTs with reflector area of 10 m2and equipped with imaging cameras of 400
PMT-based pixels. The FoV of the single telescope will be 8 ×8 degree. The inter-telescope
distances will be optimized in the range of 300–600 m.
4. A net of surface and underground stations to measure the muon component of air showers. It
is desirable to cover an area of a few 1000m2with underground muon detectors for providing
good rejection of hadron showers for energies ≥100 TeV.
2 Gamma ray observatory TAIGA
The gamma-ray observatory TAIGA will be located in the Tunka valley, about 50 km from Lake
Baikal in Siberia, Russia, were since 2009 the full-size Tunka-133 detector is in operation [8–10].
The experience of TUNKA-133 on the energy spectrum and mass composition of cosmic rays
together with it’s existing infrastructure are important factors for selecting this location.
–2–
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Figure 1. Layout of the Tunka-133 array.
2.1 Tunka-133
Presently the Tunka-133 array consists of 175 Cherenkov detectors distributed over 3 km2area
(figure 1). The central dense core of Tunka-133 includes 133 wide-angle optical detectors grouped
into 19 clusters, each with 7 detectors — six hexagonally arranged detectors and one in the centre.
These clusters are installed in a circle of 500 m radius. The distance between the detectors in
each cluster is 85 m. Six of the same “outside” clusters are placed at the distance of 700–1000 m
from the centre of the Tunka-133 array. The Tunka-133 detectors consist of a metallic cylinder
of 50 cm diameter, including a PMT (EMI 9359 or Hamamatsu R1408) with hemispherical-shape
photocathode of 20 cm diameter. The container has a Plexiglas window on the top, and is heated
against frosting. The angular aperture is defined by the geometrical shadowing of the PMT by the
metallic cylinder. The detector is equipped with a remotely controlled lid protecting the PMT from
sunlight and precipitation. Apart from the PMT with its high voltage supply and pre-amplifier,
the detector box contains a light emitting diode (LED) for both amplitude and time calibration
and a controller. The controller is connected with the cluster electronics via twisted pair (RS-
485 protocol). To provide the necessary dynamic range of 3 ×104, two analog signals, one from
the anode and another one from one of the dynodes, are read out. They are amplified and then
transmitted to the central electronics hut of each cluster. The ratio of amplitudes of these signals
is about 30. It is not planned to heat the inner volume of the optical detector boxes, therefore all
the detector electronics are designed to operate over a wide temperature range (down to −40◦C).
The signals from PMTs are sent via 95 m coaxial cable RG58 from detectors to the centre of each
cluster. This leads to a broadening of the signals, the minimum value of FWHM is close to 20 ns.
The cluster electronic includes the cluster controller, 4 four-channel FADC boards, adapter unit for
connection with optical modules and a temperature controller. The 12 bit and 200 MHz sampling
FADC boards are based on AD9430 fast ADCs and FPGA XILINX Spartan XC3S300 microchip
(figure 2). The cluster controller consists of an optical transceiver (type V23826-K305-C363-C3 for
inside clusters and type SNR-SFP-LX-20 for outside clusters), a synchronization module, a local
–3–
2014 JINST 9 C09021
Figure 2. Internal structure of FADC board (shown only one channel, the second one is identical) (left).
Example of signals from anode and dynode (right).
time clock and a trigger module. The optical transceiver operating at 1000 MHz is responsible for
data transmission and formation of 100 MHz synchronization signal for cluster clocks. The cluster
trigger (the local trigger) is formed by the coincidence of at least tree pulses from optical detectors
exceeding the threshold within a time window of 0.5 µs. The time mark of the local trigger is fixed
by the cluster clock. The accuracy of the time synchronization between different clusters is of
about 10 ns. Each cluster electronics is connected to the DAQ centre with a power cable (4-copper
wires) and optical. The central DAQ room consists of several DAQ boards synchronized by a
single 100 MHz oscillator. On each board 4 optical transceivers are installed, and so one board can
operate with 4 clusters. The full number of boards may be increased to 20 (now only 7 are used),
and hence, the DAQ can work with up to 80 clusters. The boards are connected to the master PC by
100-MHz-Ethernet lines. The energy spectrum of cosmic rays in the range of 6–1000 PeV has been
reconstructed using data from 3 winter seasons of measurements with the Tunka-133 array [10].
The spectrum points out reliably the existence of a “second knee” at 300 PeV, where at this energy
the spectral power-law index changes by 0.3: from 2.977 +
/
−0.1 below the knee to 3.3+
/
−0.1 at
higher energies. These results confirm earlier findings of the KASCADE-Grande experiment [7].
The spectral steepening at 300 PeV could be interpreted as a “second knee” in the energy spectrum
related with a transfer from galactic to extragalactic origin of the cosmic rays. We find a decrease
of the mean (in logarithm) of the atomic number or composition, this lightening at energies higher
than 100 PeV (see figure 3(right)) also points out to a transfer to extragalactic cosmic rays [10].
This may hint to the fact that cosmic rays of higher energies are of extragalactic origin
2.2 Tunka-HiSCORE array
The principle of the Tunka-HiSCORE array follows the idea outlined in [11]. It is rather similar
to the one used for the Tunka-133 array. Again, this method is based on the sampling of the
Cherenkov light front of air showers. The detector stations measure the light amplitudes and its
arrival time differences over a distance of few hundred meters. This approach allows to reconstruct
the shower direction and to measure in detail the cosmic ray spectrum and its composition, as well
as it should help providing a gamma/hadron separation through reconstruction of the maximum
of the shower height. The Tunka-HiSCORE array will be a net of optical stations based on large
–4–
2014 JINST 9 C09021
Figure 3. All particle energy spectrum (left) and dependence of the mean lnA from energy (right), as
measured by the Tunka-133 array [10].
Figure 4. The concept of a HiSCORE station.
sensitive area PMTs (figure 4) distributed over an area of several square kilometres in the first
stage of the experiment and is planned to be extended up to 10 km2in future. As in the presently
operating Tunka-133 array, each station will detect Cherenkov light from EAS with wide angular
acceptance, but summing up the outputs of 4 high sensitive 8–10 inch size PMTs. These are
equipped with Winston cones, providing an increase of PMT effective area in four times. Because
of the used sum-signal of 4 PMTs with light-guides, this array has a substantially lower energy
threshold compared to Tunka-133. Signals of all PMTs of each optical station are summed up and
digitized with a step of 0.5 ns by using fast electronics, which is based on the Digital Ring Sampler
DRS-4 chip. Simulation results show that the data of the first stage of the Tunka-HiSCORE array of
∼1 km2area will allow us to reconstruct the arrival direction of EAS with an accuracy of 0.1–0.3
degrees. It should possible to detect the bright sources as, for example, the Crab Nebula, in the
energy range ≥30 TeV.
–5–
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Figure 5. Sketch of a possible layout for the 1st stage Tunka-HiSCORE optical stations: regular grid of
optical stations with 150–200 m spacing (squares), 9 station installed in 2013 (green squares) and Tunka-133
optical station (dots).
A possible layout of 64 stations of the 1st stage Tunka-HiSCORE array, embedded into
the Tunka-133 array, is shown in figure 5. Distances between the stations are proposed to be
100–200 m. The geometric arrangement of the optical detectors is currently being optimized, aim-
ing to obtain a relatively low energy threshold (few 10’s of TeV). Since autumn 2013, a prototype
setup with 9 optical stations each with four 8” size PMTs of type R5912, have been in operation.
Each PMT is equipped with a light collector (Winston cone), with of 0.4 m diameter and a half
opening angle of 30◦made of ten strips of ALANOD 4300UP foil with reflectivity of ∼80%. The
cones increase the light collection area by a factor of four in comparison to bare PMT, resulting in
a total light collection area of 0.5 m2per station [11]. A Plexiglas plate is used on top for protecting
the PMT and the light collector from dust and humidity. It is possible to tilt the stations and thus to
change the observed FoV of the stations.
A data acquisition system (DAQ) for recording the PMT pulses with GHz sampling both from
anode and from one of the dynodes is used for improving the dynamic range. Relative time syn-
chronization electronics is used for recording the light pulses from all stations. A (sub-) ns precision
is needed to reach an angular resolution of about 0.1◦for gammas with energies above 50TeV. The
DAQ, the slow control for the PMT-high voltage, environmental control and some auxiliary elec-
tronics are placed in a special temperature-controlled box (figure 6). To acquire the shower data,
the high-gain anode signals of the PMTs are split and amplified by a custom-made fast preamplifier.
One branch of the signal is processed by a clipped sum trigger. In this trigger system, each signal
is first clipped above an adjustable clipping threshold and then the clipped signals are summed up.
If there was an event, the amplitude of the signal is four times the amplitude of a single clipped
signal. The threshold for the following trigger should be slightly below this level so that it only fires
if all four channels had a signal that was high enough to be clipped. This reduces noise, especially
the noise produced by after pulses in the PMTs. The other branch of the PMT signals is summed
–6–
2014 JINST 9 C09021
Figure 6. View of the detector (left) and electronics container (right) of a HiSCORE prototype station.
up and sampled via a custom-made DRS4 chip based board. For the same purpose we also sum up
and sample the low-output-gain PMT channels, which are read out at the fifth dynode. Every event
in each station is provided with a timestamp to reconstruct the direction of the shower. For this, the
time between stations needs to be synchronized to a nanosecond level (also see ref. [12]. The setup
for this will be based on the improved Tunka time synchronization (and is based on a calibration
signal sent over the optical fibre) combined with the White Rabbit system, a Gbit-ethernet -based
system developed at CERN for synchronization of nodes with a spacing of up to 10 km. The latter
can yield a resolution of 1 ns with a phase stability of 200 ps. Different solutions for station-to-
data-acquisition centre communication (industrial mini PCs) and for slow control are available and
are currently being tested.
2.3 Tunka-IACT array
To increase the sensitivity and to decrease the threshold, to organize monitoring of certain sources
and above all to separate more reliably gamma-ray initiated showers from the background of
charged cosmic ray initiated showers, it is really important to add an array of Imaging Atmospheric
Cherenkov Telescopes to TAIGA. We have designed and start constructing the first Tunka-IACT,
the prototype of the array. A net of IACTs are planned to be implemented within TAIGA. The
reflectors of these telescopes will have an area of ∼10 m2, and a focus of 4,75 m. We are designing
imaging camera matrixes consisting of ∼400 PMT-based pixels, with observational field of view
of 8×8 degrees and an angular size of 0.36 degrees for a pixel. Preliminary simulation results show
that joint operation of the wide-angle Cherenkov array Tunka-HISCORE and a system of Tunka-
IACTs can be a very effective inexpensive and a fast way to extend the gamma-astrophysics into the
unexplored yet super high energy domain. The basic idea is that the two arrays (TAIGA-HiSCORE
and TAIGA-IACTs) will work in a coincidence mode and will complement each other: by using the
timing and the imaging methods, showers with high precision can be measured. Namely, the results
of reconstruction of core position and arrival direction of coincident EAS events from the TAIGA-
HiSCORE array (expected accuracy of core location ∼5–6 m and arrival direction ∼0.1 degree)
will help an IACT in rejecting the high energy background events on large impact parameters of
say up to 300–500 m. On its turn, the IACT observing a concrete strong source in the sky, say, for
example, the Crab Nebula, can with high efficiency tag the coincident high energy gamma event
candidates, which can be used by the HiSCORE array for developing, training and improving its
–7–
2014 JINST 9 C09021
Figure 7. Muon distribution for EAS induced by gamma-quanta and protons with energy 30 TeV.
gamma-ray selection methods. What is particularly important that the distance between IACT tele-
scopes can be made considerably larger than in the current arrays of H.E.S.S., MAGIC, VERITAS,
CTA etc. As a result for development of 1 km2area array only 4–9 (the number will be specified
in the course of further simulations) TAIGA-IACT telescopes will be needed. Due to a dedicated,
optimized design of the IACTs the total cost of such a complex array can be an order of magnitude
lower than the cost of an IACT array of the same area built with use of only classical narrow-angle
IACT detectors.
2.4 A net of Tunka particle detectors
One of the effective ways to select gamma induced EAS is to measure the muon component of
EAS because it strongly depends on the primary particle type. This should work well for the
energy range above 100 TeV. The possibility to select EAS initiated by a gamma-ray is based on
the fact that the number of muons in a cosmic ray induced EAS is on average 30 times higher
compared to gamma-ray events (figure 7).
We intend to have a net of underground muon detectors with a total area of up to 3000 m2. As
a first step in 2014 we will finish the deployment of 19 scintillation stations some on the surface
to detect electron component of EAS and others underground for detecting EAS muons (figure 8).
Each surface detector will include 12 scintillation counters with a size 80 ×80 ×4 cm3from the
KASCADE-Grande array and each underground detector will include 8 of the same type counters.
DAQ, the synchronization and control systems for the particle detectors will be the same as the one
used for the Tunka-133 array.
In future we plan to choose an optimum and cheap design for the muon detector type and for
its data acquisition system, testing a pilot version of different detectors. Various options of muon
detectors include the following types:
1. Scintillation detectors of various types.
2. Water Cherenkov detectors on the basis of long (up to 10 m) pipes.
3. Cherenkov detectors on the basis of water tanks of large area.
–8–
2014 JINST 9 C09021
Figure 8. A net of scintillation stations. Layout of 19 such stations embedded into the Tunka-133 array
(centre). Overall view of the scintillation station (left top), surface detector box for the scintillation coun-
ters (right top), inside of the underground muon detector (left down), inside of the surface detector box
(right down).
It is supposed that muon detectors will be united into clusters of 30–50 m2. Each cluster will be
connected with the data acquisition centre with a fiber-optics cable. The data acquisition system,
cluster synchronization with accuracy of 1 ns and the digital trigger system are under design. The
possibility to reconstruct the energy and the EAS height by the Tunka-133 and TAIGA-HiSCORE
data and to estimate the number of muons will allow us to start searching of local and diffusion
gamma radiation with energies higher than 1 PeV. Existence in one installation of the Cherenkov
and muon detectors of large area will allow us help clarifying the mass composition change around
the knee (3 PeV).
3 Conclusion
In 2013, the TAIGA - collaboration has started the design and construction of a new hybrid array for
the energy range of gamma-rays above 2–3TeV (high energy gamma-astronomy) to study the high
energy gamma-radiation fluxes from known sources (including observations of Tycho), searching
for new local Galactic sources of gamma-rays with energies higher than 20–30 TeV (PeVatrons),
searching for signals from nearby extragalactic sources Mrk421 and Mrk-501 aiming to study the
gamma-ray absorption on intergalactic background radiation and searching for axion-photon transi-
tions. The study of gamma radiation in the high energy range is of interest not only for astrophysics,
but also for testing the theories predicting a violation of Lorenz invariance and for searching for
super-heavy dark matter. Figure 9shows the point source sensitivity of the gamma-observatory
TAIGA in comparison with other experiment.
It is planned to upgrade in fall of 2014 the Tunka-HiSCORE array to 32 optical stations with
distances of 100 m between stations. The total area of this array will be 0.3 km2. In 2015 the area
of the array is planned to be increased to 1 km2. Joint operation of the first Cherenkov telescopes
–9–
2014 JINST 9 C09021
Figure 9. TAIGA (Tunka-HiSCORE and Tunka-IACT) point source sensitivity (very preliminary) in com-
parison with other experiments.
TAIGA-IACT and the TAIGA-HiSCORE array in the energy range of 30–100 TeV will yield a
sensitivity of ∼3×10−13 erg cm−2s−1(with detection of 50 events in 500 hours of observations).
Such sensitivity would give us a good opportunity to measure the energy spectrum of gamma-rays
from the Tycho SNRsone one main PeVatron candidate. For TAIGA placed at 53◦N/L this source
may be observed during more than 200 hours in one year, taking into account 50% of good weather
condition). This sensitivity level would allow us to observe significant signals, from the sources
observe by IceCube as neutrinos, if these sources have Galactic origin, and would allow to make a
survey for new PeVatrons.
Acknowledgments
This work was supported by the Russian Federation Ministry of Education and Science (agreement
14.B25.31.0010, gosudarstvennoe zadanie No
¯3.889.2014/K), the Russian Foundation for Basic
Research (Grants 11-02-00409, 13-02-00214, 13-02-12095, 14-02-10001) and by the Helmholtz
Association and the Russian Foundation for Basic Research (Grant HRJRG-303) and by the
Deutsche Forschungsgemeinschaft (Grant TL51-3).
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