ISSN 1882-3432

CONTENTS
Proceedings of the 17th NICT TDC Symposium (Kashima, June 27, 2019)
VLBI at 230 GHz using SPART toward EHT and EAVN -hiTakahiro Aoki, Hiroyuki Maezawa, Sho Masui, Daiki Tsukiyama, Ayaka Tomihara, Nozomi Okada, Atsushi Nishimura, Hideo Ogawa, Kotaro Niinuma, Ryoko
Amari, Toshihisa Tsutsumi, Kenta Fujisawa, Yoshinori Yonekura, Yasuhito
Kamichi, and Makoto Inoue...... 2

Water vapor detection under rainy weather condition
Noriyuki Kawaguchi...... 6

Current Status of TEC Estimation in the Wideband Bandwidth Synthesis . . . . . 11
SoftwareTetsuro Kondo
VLBI-GNSS Collocation Survey at the Ishioka Geodetic Observing Station . . . . . 15
Saho Matsumoto, Michiko Umei, Haruka Ueshiba, Shinobu Kurihara, Takahiro
Wakasugi, Tadao Kikkawa, and Hirosi Munekane
The VLBI survey of Extremely Compact HII regions associated with a high . . . . . 19
mass star just after nuclear ignition using Kashima - Hitachi baseline
Tatsuya Ogura, Kazuhito Motogi, Takahiro Aoki, Kotaro Niinuma, Kenta Fujisawa, Yoshinori Yonekura, Mamoru Sekido, and Kazuhiro Takefuji
Pulsar Observations in Japan: Leading Role of Kashima Observatory in 2009- . . . . . 23
2019
Toshio Terasawa
Development of Wideband Antenna. . . . . 27
Hideki Ujihara
Ten years of Ibaraki Station
Yoshinori Yonekura. . . . . 28

- News - News - News - News Kashima 34-m Antenna –Damaged by Typhoon No.15 (Faxai), 2019–
Mamoru Sekido. . . . . 30

IVS NICT-TDC News No.38

2

VLBI at 230 GHz using SPART the observation in Section 4. Note that this paper focuses the SPART and does not describe thetoward EHT and EAVN -hiother stations. The observation data are under corTakahiro Aoki1 (t-aoki@yamaguchi-u.ac.jp),
Hiroyuki Maezawa2 , Sho Masui2 , Daiki
Tsukiyama2 , Ayaka Tomihara2 , Nozomi
Okada2 , Atsushi Nishimura2 , Hideo
Ogawa2 , Kotaro Niinuma1 , Ryoko Amari1 ,Toshihisa Tsutsumi1 , Kenta Fujisawa1 ,Yoshinori Yonekura3 , Yasuhito Kamichi3 ,Makoto Inoue4,5

relation and the result will be reported at another
time.
2.

The ProjectThis project is to test the millimeter-VLBI (mmVLBI) at 230 GHz using three antennas of the theGLT 12-m in Greenland, the SRAO 6-m in Korea,
and the SPART 10-m in Japan. The correlationof the observation data is supposed to be accomplished by the Shanghai Astronomical Observatory
1. Yamaguchi University, Japan; 2. Osaka(SHAO) in China. A VLBI station in Japan will be
Prefecture University, Japan; 3. Ibaraki
important to enhance the EHT capability. Aiming
University, Japan; 4. Academia Sinica Institute of
to join the EHT and extend the frequency coverage
Astronomy and Astrophysics, Taiwan; 5. Nationalof the EAVN, we planned the project and deﬁned
Astronomical Observatory of Japan, Japan
its purpose and immediate goal as follows.
Abstract: We conducted international VLBI observations at the 230 GHz frequency using the SPART 2.1 Purpose
10-m antenna of the Osaka Prefecture University
The purpose of this project is to reveal strucat the Nobeyama Radio Observatory in Japan. tural variation of active galactic nuclei (AGNs). If
Stations that participated in the VLBI were the an AGN erupts a jet with high speed, the strucGLT 12-m in Greenland, the SRAO 6-m in Ko- ture of the central region should vary in a shortrea, and the SPART. This paper brieﬂy summa- time. Observations with high cadence and angularrizes the temporal VLBI system of the SPART and resolution are required to detect such short timethe observations. The data are under correlation scale structural variation. That is why we plannedat the Shanghai Astronomical Observatory and the VLBI observations twice a short period between
results will be reported later. The observations are March 17 and 21, 2019. The actual observations
expected to lead to joining the EHT and developing were executed on March 18 and 19 considering thethe EAVN at high frequency.
weather and the status of all the stations. Targets
1.

Introduction

The Event Horizon Telescope (EHT) took an image of the shadow of the black hole in the center of M87, which was published on April 10,
2019. Seeing implementation of such global millimeter/submillimeter very long baseline interferometry (VLBI) observations, we successfully tested
interferometry at the 230 GHz frequency or the
1.3 mm wavelength using the 10-m Solar Planetary Atmosphere Research Telescope (SPART) andthe 1.85-m antenna of Osaka Prefecture University.
Advancing the millimeter-wave interferometry, we
conducted international VLBI observations using
the SPART on March 18 and 19, 2019, one month
before the press release of the M87 black hole image, and gained a foothold in joining the EHT. Stations in the VLBI observations are the SPART 10m in Japan, the GLT 12-m in Greenland, and the
SRAO 6-m in Korea. The 230 GHz VLBI observations are expected to lead to the EHT and also
the EAVN -hi- (East-Asian VLBI Network at highfrequency).
This paper elaborates on the project in Section 2,
the system setup of the SPART in Section 3, andare seven AGNs of 3C345, 1633+382, 1928+738,
3C371, M87, NGC6251, and 1653+3945.
2.2

Success Criteria

The immediate goal for the above science is toacquire VLBI fringes. Then we, the SPART team,deﬁne the success criteria are as follows:
Minimum Success To participate in the international mm-VLBI using the SPART and try todetect fringes.
Full Success To detect fringes of baselines including the SPART.
Extra Success To acquire VLBI images to which
the SPART contributes.
As of the writing of this paper, we have alreadyachieved the minimum success, that is to say, we
have completed the observations and transported
the data to the SHAO for correlation.
3.

System Setup

The conﬁguration of the SPART for the observations is described in Table 1. The start frequency

December 2019

3

4

Table 1. Conﬁguration of SPART
Position(−3871061, 3428327, 3723784) m
Slew speed
0.1 deg/s
Elevation
30–80 deg
Aperture eﬀ. 0.4 at 230 GHz
Frequency
230.200–230.712 GHz
Polarization
single linear
Sampling
1024 × 106 sample/s
Quantization 2 bit/sample
was 230.200 GHz among all the stations; however,there are diﬀerences in bandwidth and polarization
between the stations. The bandwidths of the GLT,SRAO, and SPART are 2048 MHz, 1024 MHz, and
512 MHz, respectively, and also the polarizations
are dual circular, single circular, and single linear.
The recording system used at the GLT and SRAO
was the R2DBE/Mark6 and that at the SPART
was the ADS3000+/OCTAVE. The position of the
SPART was estimated using the Google map and ahandy GPS receiver, which may cause a positional
error of tens of meters.
3.1

Relative power (dB)

0
-2
-4
-6
-8
-10
512

576

640

704

768
832
IF (MHz)

896

960

1024

Figure 1. Band-characteristics of the system.

System Layout

Figure 4 shows the SPART and two sheds used
for the observations and Figure 5 shows the system
block diagram. VLBI equipments were installed in
the Shed 1 and a GPS antenna was tied on a laundry pole on its wall. We laid two coaxial cables
between the SPART and the Shed 1 for transporting the inter-mediate frequency (IF; 0–1 GHz) and
10 MHz reference signals. The 100 V power is distributed from the Shed 1 to the Shed 2 and the
maximum current capacity is 30 A; therefore, we
had to switch oﬀ some unnecessary instruments.
3.2

2

ReceiverThe receiver block diagram and frequency conversion are indicated in Figure 6. Considering thespeciﬁcation of signal-generators used as local oscillators (LOs), we set frequencies of the ﬁrst LO,
f1 , and the second LO, f2 , to
f1 = 224 187 999 978 Hz,
f2 = 5 500 000 022 Hz.
This setting enabled us to acquire an IF signal of 512–1024 MHz corresponding to 230.200–
230.712 GHz. The band-characteristics of the receiver is shown in Figure 1. Note that no frequencyﬁlters existed before the ﬁrst LO and hence the acquired IF signal had a double sideband as shownin Figure 6. We conﬁrmed that the frequency
setting was correct through an observation of the
CO(J = 2 − 1) line of DR21 as shown in Figure 2.

Figure 2. CO(J = 2 − 1) line of DR21 observed on
March 17.
3.3

Analogue-to-Digital Conversion

We input the 512–1024 MHz IF signal to theanalogue-to-digital converter (ADC) ADS3000+with parameters listed in Table 1. Note that thehigh-order sampling of the 512–1024 MHz analogue
signal by the 1024 MHz sampling rate inverted theband of the digital data. The data handling requires attention to the bit sequence because the
OCTAVE system, the interface between the VLBIStandard Interface and the 10 Gbps Ethernet, reverses the least signiﬁcant bit and the most significant bit.
3.4

Clock

The 10 MHz reference signal is supplied to the
SPART as shown in Figure 5. Then, we transported the signal from the SPART to the VLBI instruments in the Shed 1 via the 30 m coaxial cable.
The phase stability of the reference system was notmeasured but expected enough for the VLBI observations.
The one pulse per second (1PPS) signal is not
supplied to the SPART. Therefore, we should havecreated the signal using the GPS receiver. The
time lag between the 1PPS signals of the ADC andthe GPS was displayed on a time-interval counter

IVS NICT-TDC News No.38

4


Table 2. VLBI Schedule
Start UTC Duration [s]Target11:00:00
1200
M87
11:23:00
1200
M87
11:46:00
1200
M87
12:09:00
1200
M87
12:32:00
1200
M87
12:59:00
1200
NGC 6251
13:22:00
1200
NGC 6251
13:49:00
1080
3C 371
14:10:00
1080
3C 371
14:35:00
1080
1928+738
14:56:00
1080
1928+738
15:17:00
1080
1928+738
15:42:00
1200
3C 345
16:05:00
1200
3C 345
16:29:00
1200
1633+382
16:52:00
1200
1633+382
17:16:00
1200
J1653+3945
17:39:00
1280
J1653+3945
but not monitored/recorded by a computer. We
can not know the drift of the 1PPS oﬀset but think
the oﬀset was maintained within 1.3 μs by our visual management.
4.

Observation

The observations were performed on March 18
and 19, 2019. The VLBI schedule is indicated inTable 2. Although the weather was not ﬁne, there
were no problems with the source tracking, 1PPS
oﬀset, and lock status of the LOs. There was a concern about a power level; we noticed that an oﬀset
level of power measured by a power meter occasionally and suddenly increased or decreased due
to unidentiﬁed cause. However, the power jump
did not occur during the observations and thus theobservations were thought to be succeeded.
Figure 3 shows system noise temperatures in∗, deﬁned ascluding atmospheric contribution, Tsys
∗
Tsys
=

TR
,
PSky /PR − 1

where TR indicates a temperature of a black bodyof a chopper wheel, PSky a power received whenobserving the sky, and PR a power received whenobserving the black body. The system noise increased when the weather was cloudy but decreased
to about 200 K on the latter half of the observation
on March 19 under a clear sky.
5.

Summary

We conducted international VLBI observationsat 230 GHz or 1.3 mm using the GLT 12-m, SRAO

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Figure 3. Time series variation of the system noise
temperatures.
6-m, and SPART 10-m antennas with the aims of
joining the EHT and developing the EAVN at highfrequency. For the observations, we temporally installed the VLBI equipments in the SPART as detailed in Section 3. The observation data are undercorrelation at the SHAO and the results will be reported later.
6.

Acknowledgments

We are grateful to the members of the Nobeyama
Radio Observatory, especially Dr. Minamidani, forconsidering our implementation plan and providing
advice for us.

December 2019

5

Figure 4. System layout.
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IVS NICT-TDC News No.38

6

Water vapor detection under cloud and heavy raining. Examples of comparison
results from the VLBI campaign of EURO-76 andrainy weather condition
measurement results at the Shanghai Observatory
Noriyuki Kawaguchi (kawagu.nori@nao.ac.jp),
National Astronomical Observatory of Japan,
2-21-1 Osawa, Mitaka Tokyo, Japan
Abstract: Water vapor radiometers are used tocompensate for atmospheric absorption and delaythat have a signiﬁcant eﬀect on radio astronomy
and VLBI observations. The atmospheric absorption is important to know the absolute brightnessof an observing target, and the atmospheric delay
is important to correct a delay time to be measured by VLBI for true geometrical delays to getbetter vertical solutions. Correction of the delay
time is indispensable not only for precise geodetic
measurement by VLBI and GPS but also for accurate astrometric positioning of celestial targets by
VLBI. The water vapor radiometer, however does
not work well under a bad weather condition. This
article shows a new water vapor radiometer which
introduces advanced digital technology to detect awide band spectrum of atmospheric emission. The
new water vapor radiometer makes it possible toseparate the emission spectrum of water vapor fromthe absorption spectrum of liquid water, making it
possible to measure only water vapor in the atmosphere, even under dense clouds or during rainfall.
1.

Introduction

A water vapor radiometer, WVR, is an essential device that corrects unwanted eﬀects such as
absorption and delay caused by liquid water andwater vapor in the atmosphere during radio astronomy and VLBI observations. The absorption isimportant to know absolute brightness of celestial
targets in accurate radio photometry. The delay isimportant to reduce geometrical delays which are
key observables in precise geodetic measurements
by VLBI and GPS, and astrometric measurements
by VLBI as was noted in Kawaguchi(2017)) [1] .
A new water vapor radiometer reported in thispaper, which has been developed by the National
Astronomical Observatory of Japan, NAOJ, since2017. The new radiometer was successfully demonstrated and conﬁrmed for its high performance,
ability to detect minute amount of water vapor
even in bad weather conditions.
Section 2 states that currently widely used water vapor radiometers have large measurement errors under bad weather conditions such as dense

are shown. The large error is caused by insuﬃcient
separation of brightness temperatures of water vapor and liquid water.
Section 3 shows the diﬀerence between the emission spectra of water vapor and liquid water. The
diﬀerence in the spectra is very eﬀective for separating water vapor and water droplets.
Section 4 presents the results of water vapor
radiometer measurements operated under dense
clouds and heavy raining, and shows the successful
detection of a very small amount of water vapor in
the sky.
Section 5 concludes that ultra-wide band digitalspectral measurements are important to separate
water vapor from liquid water.

2.

A large error of WVR under cloudy andrainy weather

The bad behavior of WVR responses in a geodetic VLBI campaing, EURO-72, was reported byJungho Cho(2006)[2]. He operated the Wettzel
WVR simultaneously with geodetic VLBI observations. The records of VLBI delay and the WVR
delay are shown in Figure 1. Under a clear sky the
WVR delay shows good agreement with VLBI delay but under a rainy sky the WVR delay is muchlarger than the VLBI delay.The large diﬀerence
may have caused by the liquid water, which raises
the sky brightness temperature much higher thanthe expected radiation of water vapor by absorption. It should be noted that the liquid water raise
the noise temperature by absorption, but without
delay. We need to carefully remove the liquid absorption temperature from the total sky brightness
temperature to derive the true water vapor delay.
The SHAO WVR is a WVR operated by Shanghai Astronomical Observatory for working with the
Tianma 65-m radio telescope. As was noted in
Kawaguchi(2016)) [3] , the WVR does not work wellunder the bad weather condition. In Figure 2 theExcess Path Length, EPL, measured independently
by three diﬀerent ways, GPS delays expressed by
red dots, ray tracing delays by open squares and
SHAO WVR by blue dots.
When the weather is good until around 9 o’clock,the results of the three measurements agree well,
but when the clouds begin to appear at around 10
o’clock, the delay measurements by WVR gradually get larger and larger apart from the other two

December 2019

Figure 1. Delay comparﬁson Wettzel WVR EPL
and VLBI delay

7

Figure 3. Zenith wet delay comparison during
CONT05 VLBI campaign
indicated by a blue curve and liquid absorption
spectrum is plotted with a red line. True brightness temperature of a water vapor is a height above
the red straight line of a liquid water spectrum.

Figure 2. SHAO WVR under cloudy and rainy
weather
results. The diﬀerence becomes extremely large
when it started raining at night. The delay measurements by WVR were no longer incredible. Although liquid water raises the sky noise temperature very high but does not cause any delay,
Similar extraordinary measurements by WVR
was reported in Ichikawa(2006)[4] in CONT05
VLBI campaign. Zenith wet delays measured by
WVR, GPS, VLBI and radiosondes were comparedeach other. As shown in Figure 3, the WVR delay
also shows extraordinary large values at under bad
weather.
3.

Liquid and Vapor Emission

As noted in the previous section, separation ofwater vapor and liquid water is a key to derive
the true EPL of water vapor alone exactly. Only
way to separate the liquid water from the water vapor is to measure the broad band spectrum. The
liquid absorption spectrum is almost linear withrespect to frequency. On the other hand, the emission spectrum of water vapor shows a characteristic
resonance spectrum centered at a frequency around22.2 GHz as shown in Figure 4.
In the ﬁgure, water vapor emission spectrum is

Figure 4. Spectra of water vapor and liquid water
emissions
The liquid water in a cloud and a rain significantly raise sky noise temperature. To get truebrightness of water vapor emission, contribution of
the liquid water needs to be subtracted. The cloudabsorption and the noise temperature is shown in
Figure 5 for a frequency range from 16GHz to
24GHz in case of a cloud named ’Cumulus-cumulus
Congestus’, one of nine cloud types classiﬁed by
Slobin(1982)[5]. The cloud contains liquid water
0.57g/m3 in the density distributed over the heightfrom 660m and 3400m above sea level. This type ofa cloud causes the largest attenuation and biggest
raise of sky noise temperature.
In Figure 6, liquid water emission temperatures at three frequencies around 22GHz, a resonance frequency of water vapor are calculated
in the rain rate up to 40mm/hour in the maximum according to a theoretical estimation madeby Shambayati(2008)[6].
In Figure 7, liquid water spectrum under raining

8

IVS NICT-TDC News No.38

Figure 5. Liquid absorption temperature under Figure 7. Increase of a sky noise temperature
thick cloudcaused by liquid water at a rain rate of 5mm/h
the radio astronomical beam. The lower eﬃciency
is not fatal to monitor the water vapor emission,
the strength of which is strong enough, an order ofa few tens kelvin.

Figure 6. Liquid water emission temperature underraining
at the rain rate of 5mm/hour is shown. The rainfall is not so unusual, and radio astronomy observations are often performed even under such rainfall. Figure 8. A new WVR on the VERA 20m radio
The noise temperature reaches up to 40K. Also we telescope
can see in Figure 7 that the liquid noise temperature is almost linearly increasing with frequencies
While the WVR was tested on the VERA teleat a frequency range from 16GHz to 24GHz.
scope, the weather was very bad. A thick cloud
covers over the telescope as shown in Figure 9, and
4. A new WVR operated under bad it started raining and at occasionally strong raining
weather conditionof 6mm/hour was recorded by a rain meter shown
In June, 2018 a new WVR was set right next to in Figure 10.
the radio astronomy receiver on the VERA 20-m
radio telescope as shown in Figure 8. Since the installation position is oﬀset from the radio astronomy receiver, the WVR beam is not exactly directed to the celestial target but ignored in thisexperiment. The water vapor radiometer measures
the amount of water vapor very close to the radio astronomical beam toward the target. Since
the horn for the water vapor radiometer is slightly
smaller, it irradiates an area larger than the size ofthe sub-reﬂector, causing a large spillover loss and Figure 9. Weather condition, cloudy and raining
lowering the aperture eﬃciency, but the beam size
formed by the main reﬂector is almost the same as
In Figure 11 the sky noise spectrum is shown,December 2019

Figure 10. Rainmeter record
which was observed during heavy rain. This spectrum is almost a straight line, and the spectrum
peculiar to water vapor seen around 22 GHz ishardly seen. However, detailed analysis conﬁrmed
the presence of a small 3K water vapor.

9elevation angle decreases (as the path length of theatmosphere increases), but is saturated in very low
elevations. The noise temperature of liquid water isindicated by red squares. The attenuation is evaluTliq, where Tliq is a noise temperature
ated by α = Tatmof liquid water and Tatm is atmosphere temperature. If Tliq equals to the atmosphere temperature1(around 300K), the correction coeﬃcient, 1−α, becomes inﬁnity, the sky is completely opaque. The
liquid noise temperature of around 200K in Figure
12 indicates that the heavy attenuation of 67% exists. The water vapor radiation is saturated in low
elevations, which is thought to be due to the attenuation by liquid water. The emission intensityof water vapor corrected for the absorption is indicated by green triangles. Although the absorption
was corrected, the emission intensity of water vapor was found not to be proportional to secZ but
to follow log(secZ). The logarithmic dependence
was conﬁrmed in various azimuth angles and different weather conditions. It is very strange thatthe water vapor emission has been believed to be
proportional to the path length of the atmosphere,
secZ. This should be studied more in future.
5.

Cocluding Remarks

In this paper I presented the water vapor detection under bad weather conditions, the success ofa weak water vapor emission about 3K in the noise
temperature. So far almost all WVRs in the worldshows extraordinary increase of wet delays in bad
weather conditions. It was concluded that this was
Figure 11. Tipping results under raining
caused by incomplete separation of the emission
The correctness of this small amount of water va- spectrum of water vapor and the absorption specpor was conﬁrmed by lowering the elevation angle. trum of liquid water.
In Figure 12, the horizontal axis indicates the radio The logarithmic dependence of the water vapor rawave path length (secZ, Z = 1/sin(El)) traveling diation intensity on the atmospheric path length
through atmosphere, and the vertical axis indicates has not been elucidated yet and needs to be investhe noise temperature radiated by water vapor and tigated further in the future.
liquid water.
References[1] Noriyuki Kawaguchi, Water Vapor Radiometer for Radio Astronomy and Volcano Monitoring, Technology Development Center News,
National Institute of Information and Communications Technology, Serial No. 37, page5, June 21, 2017

Figure 12. Water vapor spectrum under heavy raining
The measured noise temperature of water vapor
is indicated by blue diamonds. The radiation intensity of very little water vapor increases as the[2] Jung-ho Cho, Axel Nothnagel, Alan Roy, andRuediger Haas, A generalized scheme to retrieve wet path delays f rom water vapor
radiometer measurements applied to European geodetic VLBI network, 4th IVS General
Meeting, Concepcion, Chile, Jan. 9-11, 2006
[3] Noriyuki Kawaguchi, Water Vapor Radiometer: Challenge Again, Technology Develop-

IVS NICT-TDC News No.38

10ment Center News, National Institute of Information and Communications Technology, Serial No. 36, page 4, November 2016[4] Ryuichi Ichikawa,
Hiromitsyu Kuboki,
Masanori Tsusumi, Yasuhiro Koyama, Junichi Fujisaku, Kensuke Kokado, Masayoshi
Ishimoto and Kazuhiro Takashima, Zenith
wet delay comparisons at Tsukuba andKashima VLBI stations during the CONT05
VLBI campaign, Japan Geoscience Union
Meeting, August 2006

[5] Stephen D. Slobin, Microwave noise temperature and attenuation of clouds: Statisticsof these eﬀect at various sites in the UnitedStates, Alaska, and Hawaii, Radio Science,
Volume 17, Number 6, pages 144 3- 1454,
November-December 1982
[6] Shervin Shambayati, Atmosphere Attenuation
and Noise Temperature at Microwave Frequencies, DESCANSO Book Series, volume
10, chapter 6, p. 255, February 2008

December 201911

Current Status of TEC Estimation in the Wideband Bandwidth Synthesis Software
Testuro Kondo1,2 (kondo@nict.go.jp,
kondo@shao.ac.cn)
and Kazuhiro Takefuji1,3
1Kashima Space Technology Center,
National Institute of Information
and Communications Technology,893-1 Hirai, Kashima, Ibaraki 314-8501, Japan
2
Shanghai Astronomical Observatory,
80 Nandan Road, Shanghai, 200030, China
3
ISAS/JAXA after Oct.1, 20191.

Introduction

An algorithm using the least-squares estimation
based on a phase model was developed for determining TEC in the wide-band bandwidth synthesis(WBWS) processing [Kondo and Takefuji, 2016].
The algorithm works well. However, it is diﬃcult toapply this algorithm for data of which SNR lower
than 10. In order to improve a TEC estimation
at lower SNRs, we have developed a TEC searchfunction method in addition to the least-squares
estimation method.
2.

TECs obtained by the least-squares(LSQ) estimation using a phase model

Table 1 shows a VLBI experiment used for thisreport and Fig. 1 shows the location of baseline.

Table 1. VLBI experiment used here
Baseline
Kashima - Medicina(about 8700 km)
Date (UT)
2018/10/05(278) 07:01– 10/06(279) 14:57
Polarization H (Horizon)
Receiving Frequency (GHz)
Band#15.490 – 6.514
Band#27.990 – 9.014
Band#39.890 – 10.914
Band#412.790 – 13.814# of scans
592
Δτion (f ) is a delay caused by the ionosphere which
is given as
Δτion (f ) = αf −2 ΔT EC(2)where ΔT EC is the diﬀerence of TECs in the propagation paths leading to each station, and α is aconstant (= 1.34 × 10−7 m2 /s). Hence Eq.(1) canbe written as follows.
θ(f ) = 2π(f Δτ + αf −1 ΔT EC) + φ0

(3)Δτ , ΔT EC and φ0 are estimated by using a leastsquares method [Kondo and Takefuji, 2016].
Fig. 2 shows a comparison of TECs obtainedby the LSQ method using a phase model andthose obtained from a TEC global map provided
by IGS [IGS CDDIS DAAC, 2018]. One sigma
error after a linear ﬁtting is 3.24 TECU (TECU= 1.0 × 1016 electron/m2 ). Fig. 3 shows a histogram of the number of scans succeeded in TEC
estimation and that of failed scans for SNRs. As
shown in Fig. 3, almost all scans of which SNR
lower than 10 are failed in TEC estimation. The
success ratio is only 66 %.

MEDICINA
8700kmKASHIMA

Figure 1. Location of baseline.

In the LSQ method, phase characteristics ofcross-spectrum with respect to frequency after thewide-band bandwidth synthesis (WBWS) processing are modeled as follows.
θ(f ) = 2πf {Δτ + Δτion (f )} + φ0(1)

where f , Δτ and φ0 are an RF frequency, a residual delay and a constant phase, respectively, andFigure 2. A comparison of TECs (WBWS TEC)
obtained by the LSQ method and those (GNSS
TEC) obtained from a TEC global map.

IVS NICT-TDC News No.38

12

lows.
F (Δτ, ΔT EC)

 f1


−i2πf (Δτ +αf −2 ΔT EC) 
Swc (f )e
df 
=
f0

 f1




−i2πf Δτ

= 
SΔT
(f
)e
df
EC


(6)

f0

where

 t1
Swc (f ) =t0

Swc (f, t)e−i2πf Δτ̇ t dtand−i2παf
SΔT
EC (f ) = Swc (f )e

−1

ΔT EC

(7)

The integral portion of Eq. (6) is an expression in
the form of a Fourier transform regarding Δτ , sothat we can use FFT for calculation of this portionby calculating SΔT
EC (f ) at ﬁrst.
Figure 3. A histogram of the number of scans succeeded in TEC estimation (upper panel) and thatof failed scans (lower panel) for SNRs.

^dZd     ∆䠄 ૚ ∼ ૛ ,  䠅

In order to improve the success ratio at lower
SNRs, we have developed a new method basedon a search function similar to that for delay
and delay rate search (maybe it is a similar idea
adopted in the HOPS post-processing software[Cappallo,2014]).
3.

∆  ૚ᇱĐĂůĐƵůĂƚĞઢࢀࡱ࡯
ĐĂůĐƵůĂƚĞᇱ ,    ઢࢀࡱ࡯ ି࢏૛࣊ࢌઢ࣎ ƵƐŝŶŐ&&d

A TEC search function

A search function like that used for delay anddelay rate search is introduced here for a TEC estimation as follows.
F (Δτ, Δτ̇ , ΔT EC)
 t1  f1
−iθ(f,t)=
Swc (f, t)e
df dt
t0∆  ૚ǇĞƐ∆ ! ૛

(4)

ŶŽ

f0

where Δτ̇ is a residual delay rate, Swc (f, t) is a
wideband cross-spectrum at time t after the correction of inter-band delay diﬀerences and inner-band
phase characteristics, and θ(f, t) is a phase term att that is given asθ(f, t) = 2πf {Δτ + Δτ̇ t + Δτion (f )} + φ0

(5)where φ0 is a constant phase, and Δτion (f ) is a delay caused by the ionosphere and given by Eq.(2).
Assuming that Δτ̇ is determined well with a sufﬁcient accuracy by a WBWS process, we can integrate over time at ﬁrst. Moreover |e−iφ0 | = 1,
hence, a search function can be re-written as fol-



^dKW

Figure 4. A ﬂow of calculation of TEC search function.
The speciﬁc search procedure is shown in Fig. 4.
First, decide the range of ΔT EC as D1 ∼ D2
and step ΔD. Set ΔT EC = D1 , then calculate
SΔTEC (f ). Using this SΔT EC (f ), calculate Eq. (6)by using FFT. Increase ΔT EC by ΔD, then do
the same calculation. Repeat this process until
ΔT EC = D2 .
Values of F (Δτ, ΔT EC) calculated this way arethose on grid points shown in Fig. 5. An example

December 201913of an actual search function is shown in Fig. 6. Estimation of TEC is given as ΔT EC that shows thelargest value of search function. The position of
actual maximum value is usually determined by aparabolic approximation, because a true peak position is shifted from the grid point.
ed
uit
lp
m
A

Figure 5. Values of search function are obtainedfor grid points regarding ΔT EC and Δτ .

Figure 7. A comparison of TECs (WBWS TEC)
obtained by a TEC search function method andthose (GNSS TEC) obtained from a TEC global
map.

Figure 6. An example of TEC search function.
ΔT EC covers from −20 TECU to +20 TECU. The
value of a search function is calculated by using
FFT at each grid point using a given ΔT EC.

Fig. 7 shows a comparison of TECs obtained by
the TEC search function method and those obtained from a TEC global map. One sigma error
after a linear ﬁtting is 2.95 TECU. It’s slightly improved compared to that of the LSQ method. Fig. 8
shows histogram as same as Fig. 3. As shown inthe ﬁgure, we can see the remarkable improvementof the success ratio at lower SNRs. Total successratio has increased to 98 %.
4.

Conclusions

We have been improving a TEC estimation in
the WBWS software. The success ratio of TEC
estimation by introducing a TEC search function.
It has prominently improved the success ratio of
TEC estimation at lower SNRs.

Figure 8. A histogram of the number of scans succeeded in TEC estimation (upper panel) and thatof failed scans (lower panel) for SNRs. The success ratio at lower SNRs is remarkably improved.

One sigma error when compared with GNSS isabout 3 TECU. We were expecting a smaller one
sigma error considering from the results of theshort baseline experiments that were carried out inJapan. The diﬀerence in spatial resolution of theionosphere is considered as one of the causes. The
spatial resolution of a global TEC map obtainedby GNSS observations is only a few hundred kilometers. On the other hand, TEC estimated from

14

IVS NICT-TDC News No.38

bandwidth synthesis for geodetic VLBI, Radio Sci.,
Vol. 51, doi:10.1002/2016RS006070, 2016.

Figure 9. A comparison of TECs obtained by the
GM method (horizontal axis) and the PM method(vertical axis).

VLBI data thought to be reﬂected a smaller spatial
scale.
Another possibility is that this represents an actual estimation error. Fig. 9 shows a comparisonof TECs obtained by the GM method and the PM
method. It’s one sigma error after linear ﬁtting
is about 3 TECU. If the√ error of both is about thesame, each error is 3/ 2 = 2.1 TECU. In addition,a clear band-split structure can be seen in the scatter plot, which is also seen in Fig. 2. Whether it
reﬂects the actual condition of the ionosphere or
not is being considered.
Acknowledgement: A wide band VLBI experimentused in this report was carried out under the cooperation between NICT, INAF (Italian National
Institute for Astrophysics) and INRIM (Italian National Metrology Institute). We would like to thankstaﬀ members of INAF and INRIM engaged in theexperiment. A part of this work has been fundedby Chinese Academy of Sciences (CAS) President’s International Fellowship Initiative, Grant
No. 2018VEA00017.
References
Cappallo, R., Correlating and fringe-ﬁtting broadband
VGOS data, IVS 2014 General Meeting Proceedings,
pp.91–96, 2014.
Crustal Dynamics Data Information System (CDDISDAAC), 2018. International GNSS Service, Ionosphere vertical total electron data. Available online[http://cddis.gsfc.nasa.gov/Data and Derived
Products/GNSS/atmospheric products.html#iono]
from NASA EOSDIS CDDIS DAAC, Greenbelt,
MD, USA., Subset obtained: time period: 2018-278
to 2018-279.
Kondo, T., and K. Takefuji, An algorithm of wideband

December 201915

VLBI-GNSS Collocation Sur- tion, “ISHI”, is registered as one of the InternaGNSS Service (IGS) sites. To measure thevey at the Ishioka Geodetic Ob- tional
local-tie vector between the VLBI and the GNSS
serving Station
station, which is deﬁned as the vector between the
Saho Matsumoto(matsumoto-s96n2@mlit.go.jp),
Michiko Umei, Haruka Ueshiba,
Shinobu Kurihara, Takahiro Wakasugi,
Tadao Kikkawa and Hirosi Munekane

VLBI antenna invariant point (hereafter referred
to as the IVP) and the reference point of the IGS
station ISHI, we conducted the collocation survey
in November 2018 at the Ishioka station.
2.

Purpose of this study

Figure 1 shows a panoramic view of the Ishioka
station. Four pillars and a VLBI ground marker are
installed there. “ISK2” is an experimental CORS(Continuously Operating Reference Station). The
Abstract: The International Terrestrial Reference collocation survey was performed using these pilFrame (ITRF) is realized by a combination of so- lars and the VLBI ground marker for determining
lutions of four diﬀerent space geodetic observing the local tie vector. The IVP of the VLBI antenna
techniques. To perform this combination, it is nec- is deﬁned as the intersection of the azimuth (hereessary that relative position between these space after referred to as AZ) and elevation (hereafter
geodetic observing instruments, namely local-tie referred to as EL) axes. It must be determined invectors, are precisely known at collocation sites directly since it is not directly observable. There
where observing instruments of multiple space are two methods to estimate the IVP. One is “Outgeodetic techniques are installed. The Ishioka side method” where we observe targets attached on
Geodetic Observing Station (hereafter referred to the moving parts (near the counterweight) of theas the Ishioka station), is one of such collocation VLBI antenna with “Total Station” (hereafter resites that has a VLBI and two GNSS stations. To ferred to as TS) on the reference pillars shown inmeasure the local tie vector between them, which Figure 1. The other is “Inside method” where weis deﬁned as the vector between the invariant point observe the targets on a cabin wall from a cabin(IVP) of the VLBI antenna and the GNSS ref- pillar that is detached from the antenna moving
erence point, we conducted the collocation sur- parts. We had only experiences of the “Outside
vey in November 2018 there. In this survey, we method” so far, but the Ishioka VLBI antenna aladopted two methods to determine the IVP: “Out- lows us to carry out the “Inside method” thanks toside method” and “Inside method”. We compared its cabin pillar structure. In the “Inside method”,
the results of them, and conﬁrmed that “Inside it is expected that the observation error would be
method” produces the local-tie vector which is of less than the “Outside method” because TS is set
the similar quality as that of “Outside method” up closer to the targets. In addition, with the “Inwhich is well established. Since “Inside method” is side method”, it might be easier to estimate themore practical than “Outside method”, we decided oﬀsets and axes orthogonality between the AZ andto adopt it for future collocation surveys.
EL axes than the “Outside method”, since TS observations may be done evenly for every AZ/EL
1. Introduction
direction. We conducted the collocation survey usThe International Terrestrial Reference Frame ing both methods in November 2018 to examine(ITRF) is realized by a combination of solutions whether the “Inside method” may be adopted inof four space geodetic observing techniques (VLBI, stead of “Outside method”.
SLR, GNSS and DORIS), where VLBI solutions
3. Observations
contribute in determining the scale and EOP. In
order to achieve combination of diﬀerent solutions,
The collocation survey was performed from 8th
it is necessary to prepare relative positions between to 22nd in November 2018. We conducted it by
the instruments at a collocation site which has two following three processes.
or more space geodetic observing instruments. The
Ishioka station operates a VLBI and two GNSS sta- 3.1 Surveys between pillars and VLBI
ground marker
tions. The VLBI station is one of the VLBI Global
Observing System (VGOS) stations which is proWe measured the horizontal and vertical angles,
moted by International VLBI Service for Geodesy
and distances between the pillars and the VLBI
and Astrometry (IVS) as the next-generation stanground marker with a TS. Then we performed levdards (Ishimoto et al., 2016), and one GNSS staGeospatial Information Authority of Japan, 1Kitasato, Tsukuba Ibaraki, JapanIVS NICT-TDC News No.38

16

Figure 1. Panoramic view of the Ishioka station.
Four pillars and a VLBI ground marker are arranged around the VLBI antenna and GNSS stations.
eling surveys among them to determine relative
heights. In addition, we conducted GNSS surveys
at the pillar 2 and 3, and a reference point at the
top of Mt. Tsukuba in order to determine the orientation angle of the local frame so that we canconvert the survey results in a geocentric reference
frame.
3.2

Surveys between pillars and IGS station

We measured the horizontal angles and distances
between the pillars and the IGS station. Since theantenna reference point (hereafter referred to as
ARP) of the IGS station is not directly observable
due to the redome, it must be determined indirectly. Therefore we measured the horizontal angles at the left and the right end of the choke ring
antenna, and used their mean value. Regarding
the distance from the pillars to the IGS station, we
ﬁrst measured the shortest distance from each pillar to the bottom of the choke ring antenna, andthen calculated the distance from the pillars to the
ARP using the design value of the antenna. In addition, we measured the relative height from pillars
to the GNSS reference point and from the GNSS
reference point to the bottom of the antenna by
the leveling surveys, respectively. The sum of these
may be considered as the height of the ARP. Figure 2 shows the survey diagram about above twoprocesses (3.1, 3.2).

Figure 2. Observation diagram of the collocation
survey in 2018 (except the determination of theIVP).
of the antenna (right in Figure 3). We pointed
the VLBI antenna to 8 direction around the AZ
axis and 5 direction around the EL axis for each
AZ direction. Observed targets should be on theimaginary spherical surface (lower left in Figure 3),
and we can consider the center of it would be the
IVP.
3.3.2

Inside method

We observed the targets from the TS installedinside cabin pillar of the VLBI antenna (right inFigure 4). First we observed the targets on surrounding pillars to determine the position of theTS. Then we installed the targets on the cabin wallof the AZ cabin and carried out measurements fordetermining the AZ axis. We set the antenna EL tozero, and observed four targets while changing theantenna AZ by 30 degree from 0 to 330 degree (leftin Figure 5). Then we carried out measurements
for determining the EL axes. In order to investigate the dependency of the estimated EL axes on
diﬀerent azimuthal direction, we tentatively estimated the EL axes at the AZ of 125 degree and215 degree. For each antenna AZ, we observed twotargets while changing the antenna EL by 10 de3.3 Determination of the IVP of VLBI an- gree from 0 to 90 degree (right in Figure 5). In
tenna
total, we obtained 48 target positions for the determination of the AZ axis and 40 target positions
3.3.1 Outside method
for the determination of the EL axes. We ﬁt cirWe installed the target outside the moving parts cles to trajectories of the target and the AZ axis isof the VLBI antenna (upper left in Figure 3). Then estimated as a straight line connecting the center
we observed the target from the TS installed on of these circles. The EL axis is estimated in thethe pillars while changing the AZ and EL angles similar way as that of the AZ axis. Then the IVP

December 201917

Figure 5. Observations of the targets on the cabinwall. We observed the targets from the TS whilerotating the antenna AZ(left) and EL(right).
Table 1. IVP coordinate (※ According to the “local
frame” whose origin lies on Pillar 2 and X axis
Figure 3. Observation of the “Outside method” for is aligned with the line from the Pillar2 to Pillar3
estimating the IVP. The TS was put on the pil- (Figure 2))XYZ
lar and moved to other pillar according to the AZ
Outsidemethod(m)
31.482430.7626
9.2577rotation.
Inside method (m)31.4841 30.7623 9.2570
Δ(mm)
-1.7
0.3
0.7

IGS reference point with the “Inside method” inthe geocentric frame (Table 2). To conﬁrm our calculation by the “Inside method”, we independently
calculated the local-tie vector with the adjustment
software pyaxis (LINZ, 2015). We found that thediﬀerences of the local tie vectors between those by
our calculations and those by pyaxis were negligibly small in practical use. The standard deviationsof the local-tie vector for the “Inside method” estimated using pyaxis are as small as 0.7 mm, 0.7mm, and 0.6 mm for x, y, and z component, respecFigure 4. Observations inside the antenna cabin tively. In terms of the number of staﬀ and working
for estimating the azimuth and elevation axes. The hour, the “Inside method” is more eﬃcient thanTS was put on the cabin base, which does not follow the “Outside method” (Table 3) because we obthe antenna movement.
serve the target only from the cabin base for the“Inside method” instead from four pillars as done
is deﬁned as the closest point between the AZ and for the “Outside method”. From the results, we
conﬁrmed that the “Inside method” has suﬃcientthe EL axes.
accuracy and it is more practical. In this sense, the“Inside method” can be adopted as an alternative
4. Resultto the “Outside method”.
Table 1 shows the IVP coordinates determinedby both methods and their diﬀerences. Note that 5. Summarythese coordinates are expressed in the “local frame”
We conducted the collocation survey at the Ishwhose origin lies on Pillar 2 and X axis is aligned
ioka
station in November 2018. We determinedwith the line from the Pillar2 to Pillar3 (Figure 2).
IVP
positionusing two methods independently, theThe diﬀerences are negligibly small. Then we de“Insidemethod”
and the “Outside method”. We
termined the oﬀsets and orthogonality between thecomparedtheresultsof two methods and conﬁrmed
axes with the “Inside method” (Figure 6). Oﬀsets
that
the“Insidemethod”
could be a practical albetween the AZ axis and the EL axes for AZ of 125
ternativetothe“Outside
method”
in terms of its
degree and 215 degree were 0.2 mm and 0.1 mm, reaccuracy
andeﬃciency.
With
the“Insidemethod”,
spectively. The angles between the AZ axis and two
it
may
be
possibleto
conductacollocation
surveyEL axes were 90.0006 degree and 89.9955 degree,
more
frequently
since
it
takes
less
time
to
survey.
respectively. Finally, we calculated the local-tie
vector between the IVP of VLBI antenna and the We will submit the estimated local-tie vector to theIVS NICT-TDC News No.38

18[2] Land Information New Zealand, 2015 - Warkworth Observatory Local Tie Survey.

Figure 6. Oﬀsets and the orthogonality of the axis.
Table 2. The estimated local tie vector between the
IVP of the VLBI antenna and ISHI in the geocentric reference frame
X
YZ
Baseline
Inside(m) -12.7427 10.8579 -37.0948 40.6976
Inside(m) -12.7410 10.8611 -37.0962 40.6992(pyaxis)
Δ(mm)
-1.7
-3.2
1.4
-1.6
Table 3. The comparison of personnel and working
hour of between two methods.
day the number of staﬀ
Outside method
76∼9Inside method
15International Earth Rotation and Reference Systems Service (IERS) to be used in the constructionof the next ITRF.
6.

Acknowledgments

We would like to thank the Land Information New Zealand (LINZ) for making the pyaxis software publicly available(https://github.com/linz/python-linz-pyaxis).
We acknowledge Dr. Chris Crook at LINZ for his
assistance for the use of the software.
References[1] Ishimoto M., M. Umei, T. Wakasugi, R. Kawabata, T. Toyoda, B. Miyahara, T. Fukuzaki,
Status on the Ishioka Geodetic Observing Station, IVS NICT-TDC News, vol.36, pp.9-12,2016.

December 201919

The VLBI survey of Extremely sources both inside and outside Galactic plane.
are three research topics : star-formation,
Compact HII regions associ- There
active galactic nucleus, Galactic black hole. This
ated with a high mass star paper focuses on a topic of the star-forming.
just after nuclear ignition using
Kashima - Hitachi baseline
Tatsuya Ogura, Kazuhito Motogi, Takahiro
Aoki, Kotaro Niinuma, Kenta Fujisawa(Yamaguchi Univ.) Yoshinori Yonekura(Ibaraki Univ.) Mamoru Sekido, Kazuhiro
Takefuji (National Institute of Information
and Communications Technology)
Abstract: A size of infant HII region createdby a high mass star just after the nuclear ignition is decided by a gravitational radius forionized gas. This size is expected to be about
100 astronomical unit (au). Physical parameters(mass, luminosity, etc) of host stars exciting suchan Extremely Compact (EC) HII region give strong
observational constraints on a theoretical evolutionary model of a high mass protostar. But,
the angular size of such a newborn HII region
is very small (∼ 100 mas), because high mass
star-forming regions are mostly located on beyond
1 kpc from the Earth. Therefore, a high-sensitivity
single-baseline VLBI is a useful tool to survey such
EC HII regions, because one can selectively detect HII regions excluding other extended radiosources. For the above reasons, we started a newfringe survey towards known VLA radio sources using Kashima34m-Hitachi32m short-baseline VLBI
having the angular resolution of 89 mas at X-band.
The targets are 662 point-sources which detectedby the 5GHz survey on Galactic plane using VLA(The CORNISH survey, 1”.5, 0.4 mJy at 1σ). We
selected infrared quiet sources out of the CORNISH
sources. The most of selected targets are very radio quiet. The brightness-temperature sensitivityof the Kashima-Hitachi baseline is ∼ 103 K (600 sec
at 1σ). This is enough sensitive to detect thermal
free-free emission (∼ 104 K). We have already observed 270 targets on August 2019 and successfully
detected fringes of 122 targets. They are plausible
candidates of the EC HII region and we will conduct detailed follow-up studies by using J-VLA and
ALMA.
1.

Introduction

Figure 1. Array conﬁguration and photos of each
JVN station.
2.

Infant HII regions associated with a
newly-born high mass star

Ultraviolet radiation produced by a high mass
star generates a region of ionized gas called the HII
region. This section presents the size of infant HII
regions associated with a high mass star just after
nuclear ignition and its observational requirements.
2.1

A size of very young HII regions

A size of HII region formed by a high mass star
just after the nuclear ignition depends on a relative size between the gravitational radius (rg ) andStrömgren radius (rS ). Equating the gravitationalenergy and the kinetic energy of ionized gas, rg isshown as follows,
rg ∼

GM
ve2

(8)

where ve is sound speed of the ionized gas (WhenTionized gas = 104 [K], ve = 13 [km/s].). rg is aradius in which the ionized gas can be conﬁned by
the stellar gravity. If rg is larger than rS , the radiusof a HII region is given by rg (Figre 2). Assuming
M = 20M and ve = 13 [km/s], rg is ∼ 100 au.
Such an EC HII region is a plausible signature ofsuch a very young star and has not been detected
yet. If one can be detected the EC HII region, it ispossible to give a strong limit on a theory of highmass stellar evolution based on several physical parameters (stellar mass, luminosity, etc...).

We are conducting a fringe survey using
Japanese VLBI Network (JVN) sub-array
2.2 Obsercational requirements
which consists of Kashima, Hitachi, TakaRequired array speciﬁcations for detectiong thehagi, Yamaguchi-1st, and Yamaguchi-2nd stations(Figure 1). Our survey targets on a compact radio EC HII region are as follows. The ﬁrst request

IVS NICT-TDC News No.38

20

Figure 2. A size of HII region.
is high angular resolution : θ ∼ 100 [mas]. Highmass star-forming regions are mostly located on
1 − 10 kpc from Earth. Therefore, the angular size
Figure 4. Flux of targets .
of its HII region (rg = 100 [au]) is under 100 mas.
The second request is high brightness-sensitivity: Tb ∼ 104 [K]. One must detect thermal free- 3.2 Observational method
free emision (∼ 104 K). Brightness-sensitivity ofThis survey is performed in the single-scan∼ 103 K/1σ is required.
snapshot-mode in which the integration-time is
3. Overview of the new fringe survey
10 min/source. The performance of KashimaHitachi array is shown in Table 1. Flux densities
This project is the ﬁrst stage of EC HII region
are calibrated by ﬂux-calibrator NRAO530, adoptsurvey. To discover them, We are searching foring a single-dish ﬂux taken by Hitachi telescope.
compact radio objects (∼ 100 mas) on Galactic
plane using Kashima-Hitachi VLBI.
3.1

Table 1. The array performance

Targets

Kashima-Hitachi VLBI
Band
8192 − 8704 MHz
Targets were selected from the CORNISH surBaseline-length82.3337 km
vey (Hoare et al. 2012 [1]). This survey is a
Angular resolution∼ 89 mas (at max baseline-length)
Coherence-timeover 600 sec
blind survey coordinated with the northern Spitzer
Correlation
Software-correlator GICO3 created by NICT
◦
◦
◦
From
obsercational
results, When Δt = 600 [sec],
GLIMPSE I region (10 < l < 65 , |b| < 1 ) using
1σ = 0.3 − 0.5 [mJy]
Sensibility
VLA (5 GHz, θ = 1”.5, ΔS = 0.4 mJy/1σ). The
However, these value depend on projected baseline-lengths.
selected targets satisfy the following two characteristics : Infrared-quiet1 and Point source against
1”.5 beam. The distribution of targets is shown in 4. Progress
Figure 3. Figure 4 shows that the most of selected
We have already observed 255 targets (Table 2).
sources have a low ﬂux density of 0 − 10 mJy at
Thischapter reports the progress to date (Septem5 GHz.
ber 2019).

Table 2. Observational scheduleDate
2019/2/21
2019/7/10
2019/7/11
2019/7/25
2019/7/26
2019/7/27
2019/8/7
2019/8/8
2019/8/9

Observational code
U19052A
U19191A
U19192A
U19206A
U19207A
U19208A
U19219A
U19220A
U19221A

Number of targets
7
31
31
31
31
31
31
31
31

Figure 3. The distribution of targets.
4.1
1 Thereare no mid-infrared Spitzer GLIMPSE I, λ =3.6, 4.5, 8.0 μm point source within radius of 2”. It is anindex of young stars.

Detection statistics

Targets were recognized as a detection, when we
could obtain a valid delay solution. Our detection criterion is that the resultant SNR ≥ 6 andresidual-delay is ≤ 1 nsec. We detected 122 sources

December 201921

out of 255 targets (48 %). However, these detec- some sources at 104 K < TBmin < 105 K. Some oftions contain several known species such as back- them may be an optically-thick thermal free-freeground AGN, and protostellar jets, young stellar source.
object ﬂares, etc. The spatial distribution in Galactic plane and histogram of Galactic longitude isshown in Figure 5 and Figure 6, respectively.

Figure 5. The distribution of detections.

Figure 6. The histogram of detections.

4.2

Figure 7. The histogram of TBmin .

Figure 8. The spectral index between 5 GHz (TheCORNISH survey) and 8 GHz (This work).

Physical properties

Brightness-temperature TBmin and Spectral in- 4.3 Next plandex (αmin ) were estimated by observed ﬂux as aLong-baseline Yamaguchi-Hitachi VLBI: As
lower limit.
a next step, we identify compact non-thermal obBrightness-temperature: Figure 7 shows the jects in detected samples by Yamaguchi-Hitachi
There are many low- VLBI. Any thermal sources should be resolvedhistogram of TBmin .
brightness sources of 104 − 105 K. A temperature out by 10 mas resolution. We already includedof an HII region could reaches a few 104 K in case Yamaguchi-Hitachi baseline in U19052A. We obof extremely optically thick condition. This might tained 4 detections and 3 non-detections. The
be a candidate of thermal free-free source (under former showed high brightness temperatures of
TBmin ∼ 105 K. They are clearly background AGN.
investigation now).
The latter have a source size of 10 − 100 mas, andSpectral index: Figure 8 shows the relation of can be a Galactic object. Other detected sources
αmin and TBmin . Most of the sources showed nega- will be observed within this year.
tive αmin , suggesting a non-thermal emission from
AGN. On the other hand, positive αmin is seen in

IVS NICT-TDC News No.38

22

Short-baseline Usuda-Hitachi VLBI: Unfortunately, Kashima telescope was destroyed by the typhoon No.15 in September 2019. However, we hope
to complete this survey. Now, we consider to use
another short baseline sub-array such as UsudaHitachi baseline. Figure 9 presents a projectedbaseline in diﬀerent hour-angles. In some hourangles, we can obtain similar baseline to that ofKashima-Hitachi VLBI.

5.

Summary and Future work

The purpose of this survey is to discover candidates of the EC HII region. We have already observed 255 targets by using the Kashima-Hitachi
VLBI. The detection rate in the short-baseline is
48 % and the detected samples could include candidates of thermal Galactic objects. The next plan
is observing them using long-baseline VLBI to identify non-thermal objects. In addition, we will investigate details such as spectral index, spatial structure, etc for low brightness-temperature objects using VLA. On the other hand, we will also observe
circumstellar structures using ALMA, studying thecurrent stage of star-formation.
References[1] M. G. Hoare, et al., 2012, PASP, 124, 939

Figure 9.
VLBI.

Projected baseline of Usuda-Hitachi

December 201923

Pulsar Observations in Japan: 2. Giant radio pulses from the Crab pulsar
Leading Role of Kashima Ob- A few pulsars, including the Crab pulsar, emit
enigmatic sporadic intense radio pulses called ‘Giservatory in 2009-2019
ant Radio Pulses’ (GRP) (Fig. 1, 2) [5]. The ﬂux
Toshio Terasawa (terasawa@icrr.u-tokyo.ac.jp),
ICRR, University Tokyo,277-8582 Kashiwa-no-ha, Kashiwa, Chiba, Japan,
and, Mizusawa VLBI observatory, NAO,
181-8588 Osawa, Mitaka, Tokyo, Japan
Abstract: In 2009, we started Crab pulsar observation in collaboration with NICT colleagues in
Kashima. Since then we have extended the observation network for pulsars, not only for the Crabpulsar but also for a magnetar XTEJ1810-197. In
this paper I have reviewed our observational studies of pulsars in these 10 years.
1.

distribution of GRP shows power-law dependence
to the peak intensity (Fig. 3) in contrast to Gaussian or log-normal dependence to the peak intensity for ordinary pulses from the majority of otherpulsars.
In [3], we made a simultaneous observation of
GRP from the Crab pulsar with a frequency coverage from 0.3 GH to 8.4 GHz (widest-ever in 2016)
in which the Kashima observatory provided 1.51.7 GHz data. Our observation showed that about
70% or more of the GRP frequency spectra are consistent with single power law with spectral indicesfrom -4 to -1.

Introduction

The Crab pulsar, one of the most famous pulsars,
was discovered in 1968 [1], and has been playing theimportant roles in the basic study of pulsars (e.g.,[2]). In these 10 years, we have made,1. The ﬁrst simultaneous multi-frequency detection of giant radio pulses from the Crabpulsar [3],2. Search for the correlation between giant
radio pulses and X-ray pulse intensities of the
Crab pulsar [4],3. Independent identiﬁcation of the Crab pulsar’s glitch in 2017 (Hisano et al., in preparation, 2019).
For the above study of the Crab pulsar wehave organized the radio observation network
in Japan combining Kashima observatory withother laboratories, the Usuda Deep Space Center (UDSC)/JAXA, the Iitate Planetary Radio
Telescope (IPRT)/Tohoku University, the Center
for Astronomy/Ibaraki University, the ResearchInstitute for Time Studies (RITS)/Yamaguchi
University, the Solar Wind Imaging Telescope(SWIFT)/Nagoya University, and the VLBI Exploration of Radio Astrometry (VERA) project/NAO.
In addition to these studies for the Crab pulsar,
we have made,4. Campaign observation of the radio outburst of a magnetar XTEJ1810-197 in 20182019 (Eie et al., in preparation, 2019).
In the following sections we summarize the results for the above four items.

Figure 1. An example of a Crab pulsar’s GRP observedin Kashima. The upper panel shows the variation of thesquared antenna voltage (arbitrary unit) during a 20ms interval between 16:32:00.000 UTC and 16:32:00.020 UTC on
5 December 2010. The lower panel shows a corresponding
dynamic spectrum, where the abscissa and ordinate show
time and frequency. The power spectral density is shown inan arbitrary unit with pseudo-color. A GRP appeared as agreen down-tone chirp signal from 1435MHz at ∼6.5ms to1405MHz at ∼11.5ms.

Figure 2. Giant radio pulses from the Crab pulsar observedin Kashima are shown in the rotation-phase-UTC diagram.
Pulse intensities are shown in S/N with color-coded circles.
Pulses clustering around 180o are of main-pulse GRP, andthose around 326o are of interpulse GRP.

IVS NICT-TDC News No.38

24

tively. The sum of total numbers of pulses amounts
to 4.21×106, which corresponds to the eﬀective observation time of 39 hours.
Preliminary results show that there is a ﬁnite
ﬂux enhancement (several% with a signiﬁcancemore than ∼4.5σ) of X-ray pulses concurrent with
GRP. The detail of this correlation is now underinvestigation (Enoto et al., 2019, in preparation).

Figure 3. The flux distribution of the main-pulse GRP peakintensities in the 2.2-2.3GHz band.

3.

Search for the correlation between giant radio pulses and X-ray pulses of the
Crab pulsar

The Crab pulsar emits pulses in wide frequency
range, coherently in radio band, and incoherently
in optical to X/γ-ray bands. It is found that the
optical pulses from the Crab pulsar show 3% enhancement (with 7.2σ signiﬁcance) concurrent with
GRP emissions in the radio band [6, 7]. While thesedetections imply that the incoherent optical emission is physically linked to the coherent radio emission, the mechanism of this linkage is not yet understood. Despite intensive eﬀorts to search the correlation with radio emission at even higher energy
bands, only upper limits have been obtained for
X-ray and gamma-ray energy band [4, 8]. Contribution of the Kashima observatory was made in [4]where we did a correlational study of X-ray pluses
from the Hitomi satellite and Crab’s GRP from theKashima observatory. We set ∼22% upper limit of
the X-ray enhancement (with 3σ signiﬁcance) concurrent with the GRP emission in the radio band.
In 2017, the neutron star interior composition
explorer (NICER), a payload of the international
space station (ISS), started its X-ray observation(0.2-12keV) and improved the X-ray photon statistics by 1-2 or more orders of magnitude thanthe previous X-ray observations. In the periodof 2017-2019, we conducted simultaneous X-ray
pulse observation with NICER and GRP observations in Kashima and Usuda (UDSC) of the Crabpulsar (Table 1). The third column of Table 1,‘unassociated’, shows the accumulated number ofpulses other than GRP, while the fourth and ﬁfth
columns, ‘MPGRP’ and ‘IPGRP’, the accumulated
numbers of main pulse and interpulse GRP, respec-

Table 1. The accumulated numbers of pulses detected
simultaneously both in radio and X-ray observations. Total numbers of pulses are 4,186,598, 24,537, and 1,712 forunassociated, MPGRP, and IPGRP, respectively.

4.

The 2017 glitch of the Crab pulsar

The Crab pulsar has been monitored ∼50 years
since its discovery, and shown ∼0.5Hz decrease
from its value ∼30.22Hz in 1969 [9]. This decrease occurred more or less continuously except
sporadic spin-up events known as glitches, which
are regarded as physical indicators of the coupling
process between the outer region (the star crust+ the magnetosphere) and the inner region (thesuperﬂuid star core) of the pulsar. On November
2017 there was an announcement of the largestever glitch detection for the Crab pulsar [10]. At
this particular epoch, SWIFT/Nagoya University
was conducting daily 325MHz band observation of
the Crab pulsar, to prepare the solar corona observation at the time of the Crab pulsar’s solar
conjunction in June 2018 (Tokumaru et al., 2019,submitted to Solar Phys.). We then started ourDecember 201925

glitch search based on the SWIFT data, and conﬁrmed the occurrence of this glitch (Fig. 4). Whilethe daily Crab observation by SWIFT was the keyof the identiﬁcation of the glitch, the Kashima
and IPRT data were indispensable in eliminating
the 2πN uncertainty from the residual calculation.
The details of the glitch parameters based on the
SWIFT data are now being studied (Hisano et al.,2019, in preparation).
Figure 5.

Pulses from the magnetar XTEJ1810-197 ob-served in Kashima are shown in the rotation-phase-UTC
diagram. The rotation period on 14 December 2018 was
5.54143sec. Pulse intensities are shown in S/N with colorcoded circles.

Figure 4. The residual of the arrival times of Crab pulses
based on our radio observation: Till MJD 58063 (6 Nov.
2017), the pulse arrival times were consistent with the preglich model prediction [11], but showed the departure after
MJD 58064. This is the evidence of the glitch on MJD
58064, and consistent with the individual result based on
the Jodrell Bank observation [12].

5.

The XTEJ1810-197 radio outburst in
2018-2019

Magnetars are one group of pulsar population
with the characters of slow spin but fast spindown, which are generally taken as the evidence of
their very strong magnetic ﬁeld (1014−15 G). Whilethey were historically considered as radio-quiet,
recent observations reveal radio outbursts from a
few magnetars. Most recently, a magnetar called
XTEJ1810-197 showed its second outburst early
December 2018 [13].
We organized a ToO observation campaign for
XTEJ1810-197 activating the Japanese networkof pulsar observation with the frequency coverage
from 325MHz to 22GHz, and identiﬁed pulses fromthis magnetar. During this campaign, Kashimaobservatory successfully provided the 2.2-2.3GHz
band data (Fig. 5).
The ﬂux distribution of pulses from XTEJ1810197 (Fig. 6) shows a dependence to the peak intensity much softer than power law, in contrast to
the Crab pulsar case (Fig. 3). Details of our observation are being studied (Eie, Master Thesis,
Department of Astronomy, Univ. Tokyo, 2019; Eieet al., 2019, in preparation).

Figure 6. The flux distribution of the peak intensities of
XTEJ1810-197 pulses in the 2.2-2.3GHz band.

6.

Summary

As reviewed in preceding sections, the observational pulsar studies in Japan in the last 10 years
were very successful. We hope to continue/expand
this activity into the new era of the SKA/ngVLA
projects.
7.

Acknowledgments

We thank to all individuals who have contributed
to this successful pulsar studies in the past 10 years(in alphabetical order);T. Akahori, T. Aoki, K. Asano, Y. Asaoka, S.
Eie, T. Enoto, K. Fujisawa, T. Hada, N. Hiroshima,
T. Hirota, S. Hisano, M. Honma, H. Imai, W.
Ishizaki, O. Kameya, K. Kashiyama, R. Kataoka,
N. Kawai, S. Kisaka, H. Kita, E. Kobayashi, M. Kojima, T. Kondo, M. Kumamoto, S. Matsukiyo, R.
Mikami, H. Misawa, R. Morita, Y. Murata, K. Nagata, K. Niinuma, H. Odaka, M. Oeda, Y. Ohira,T. Oyama, S. Oyama, Y. Sato, M. Sekido, S. Shibata, K. Takahashi, K. Takefuji, H. Takeuchi, Y.T.

26

Tanaka, S.J. Tanaka, K. Tawara, Y. Terada, M.
Tokumaru, S. Torii, Y. Toyotome, F. Tsuchia, Y.
Yamakoshi, R. Yamazaki, Y. Yonekura, N. Yonemaru, and T. Yoshida.
References[1] Staelin, D. H., and E. C. Reifenstein III, Pulsating radio sources near the Crab nebula, Science, 162, 1483, 1968.
[2] Hankins, T. H., and J. A. Eilek, Radio emission signatures in the Crab pulsar, ApJ, 670,
693, 2007.
[3] Mikami, R., et al., Wide-band spectra of giant
radio pulses from the Crab pulsar, ApJ, 832,
212, 2016.

IVS NICT-TDC News No.38[6] Shearer , A., et al., Enhanced optical emission
during Crab giant radio pulses, Science, 301,
493, 2003.
[7] Strader, M. J. et al., Excess optical enhancement observed with ACRONS for early Crab
giant pulses, ApJL 779, L12, 2013.
[8] Lundgren, S. C., et al., ApJ, 453, 433, 1995;
Bilous, A. V., et al., ApJ, 728, 110, 2011;
Bilous, A. V., et al., ApJ, 749, 24, 2012; Aliu,
E., et al., ApJ, 760, 136, 2012.
[9] Lyne, A. G., et al., 45 years of rotation of the
Crab pulsar, MNRAS, 446, 857, 2015.
[10] Shaw, B. et al., A large glitch in the Crabpulsar (PSR B0531+21), Atel #10939, 8 Nov.
2017.

[4] Hitomi Collaboration (F.Aharonian et al.) andradio pulsar collaboration (T. Terasawa et al.), [11] Jodrell Bank Crab Pulsar Monthly Ephemeris,
http://www.jb.man.ac.uk/pulsar/crab.html.
Hitomi X-ray studies of giant radio pulses
from the Crab pulsar, PASJ, 70, 15, 2016.
[12] Shaw, B. et al., Mon. Not. Roy. Astron. Soc.,
Vol. MNRAS, 478, 3832, 2018.
[5] Argyle, E., and J. F. R. Gower, ApJL, 175,
L89, 1972; Sallmen, S., et al., ApJ, 517, 460,
[13] Lyne, A. et al., Intense radio ﬂare from the1999; Cordes, J. M., et al., ApJ, 612, 375,
magnetar XTEJ1810-197, Atel #12284, 11
2004; Bhat, N. D. R., S. J. Tingay, and H. S.
Dec. 2018.
Knight, Astrophys. J., Vol. ApJ, 676, 1200,
2008; Karuppusamy, R., et al., AA, 515, A36,
2010; Crossely, J. H., et al., ApJ, 722, 1908,
2010; Majid, W. A. et al., ApJ, 741, 53, 2011;
Mickaliger, M. B., et al., ApJ, 760, 64, 2012.

December 201927

Developmentof Wideband Antenna

Hideki Ujihara (ujihara@nict.go.jp)
Kashima Space Technology Center,
National Institute of Information
and Communications Technology,893-1 Hirai, Kashima, Ibaraki 314-8501, Japan
1.

Status of development

Figure 2.
Beam patterns of NINJA feed for MARBLE at 4.0
GHz(simulation).

MARBLE1 and 2 with 2.4 m dish and NINJA
feed were ready for intercontinental experimentof VLBI Time and Frequency Transfer, Gala-V.
MARBLE1 was transported to Medicina radio observatory and set in the summer of 2018, and surface of 34 m was well adjusted by radio holography by broad cast satellite to start the experiment(Fig.1). In 2019, lens of the NINJA feeds was
replaced to improve eﬃciency and reduce returnloss. Figure 2 and 3 shows far-ﬁeld beam patternsof the NINJA feed with new lens. They are more
ﬁtted to the subtended angle of the sub-reﬂectorof the MARBLE, 52 degrees. Also a orthogonalmode transducer(OMT) was developed to reduce
the return loss. Aperture eﬃciency of MARBELs
are estimated to reach around 50 % by this upgrade Figure 3.
Beam patterns of NINJA feed for MARBLE atand 34 m to around 40 %.
13.6.0 GHz(simulation).

References[1] Ujihara, H., (2016) Development of wideband
feed for Kashima 34m Antenna, Radio Sci., 52,
doi:10.1002/2016RS006071
Figure 1.
Transported MARBLE1 from Kashima in Japan to
Medicina radio observatory in Italy.

2.

Plans

Development of novel water vapor radiometerhas been started, which will realize rapid and accurate measurement of vapor, water and oxygen in
the atmsphere 18-64 GHz by one reciever systemwith a wideband feed. Studies to extend the band
width and improve of the eﬃciency of the wideband
OMT and feed are still keeping to improve the performance of Gala-V and for other applocations.

[2] Ujihara, H., Takefuji, K., Sekido M., andIchikawa R.,(2018) “Development of Wideband
Antennas”, International Symposium on Advancing Geodesy in a Changing World, Proceedings of the IAG Scientiﬁc Assembly, Kobe,
Japan, July 30 - August 4, 2017 (peer-reviewed),(10.1007/1345 2018 41).

IVS NICT-TDC News No.38

28

Ten years of Ibaraki Station• 2010/Jun./10, First fringe of Hitachi antenna
with Mizusawa 20-m and Iriki 20-m antennas
toward a methanol maser source G9.62 using
6.7 GHz cryogenically-cooled receiver with K4
back-end (DFC-2000, DIR-1000) (ﬁg. 4).

Yoshinori Yonekura(yoshinori.yonekura.sci@vc.ibaraki.ac.jp),• 2010/Jul./01, First light of Takahagi antenna (tracking the Sun using 6.7 GHzroom-temperature receiver and power meter)(ﬁg. 5).

Center for Astronomy, Ibaraki University,
2-1-1 Bunkyo, Mito, Ibaraki, Japan
Abstract: Ten years have passed since the beginning of the conversion from two satellite communication antennas to radio telescopes in Ibaraki Station at 2009. I brieﬂy summarize the history of the
Ibaraki Station.
1.

• 2011/Mar., Damaged by The 2011 oﬀ the Paciﬁc coast of Tohoku Earthquake (ﬁg. 6).

Introduction

The Takahagi and Hitachi 32 m antennas, which
had been used for satellite communications at 4
and 6 GHz by KDDI, were decommissioned in2007 March. These antennas were handed over to
the National Astronomical Observatory of Japan(NAOJ) on 2009 January 31, and now belong to
the Ibaraki station, which is a branch of the Mizusawa VLBI Observatory of NAOJ.
The Hitachi antenna was constructed in 1983
and the Takahagi antenna was constructed in
1992 by the Mitsubishi Electric corporation. The
mechanics and surface of the antenna were notchanged. The feed horn was newly designed andinstalled. By exchanging the bottom of the horn,
we can change the observation band from the lowfrequency band (6.7 and 8 GHz) to the highfrequency band (22 GHz) and vice versa. Antenna control software was newly developed. Twocryogenically-cooled receivers for 6–9 GHz and 20–
25 GHz were equipped for each antenna. A hydrogen maser, wide-band recording systems, singledish recording systems were also installed. For
more details, please refer to Yonekura et al. (2016).
2.

• 2010/Aug., VLBI scientiﬁc operations with
Hitachi antenna have started.

The Chronological Table of the Ibaraki
Station• 2012/Dec., Single-dish scientiﬁc operations
with Hitachi antenna have started.
3.

History of Back-end
History of back-end is summarized in ﬁg 7.

4.

Acknowledgments

We are grateful to NICT for the use of correlation
software.
References[1] Yonekura, Y., et al. 2016 PASJ, 68, 74.

Figure 1. Hitachi first light.

• 2009/Apr./23, First light of Hitachi antenna(drift-scan [no tracking] toward the Sun using
6.7 GHz room-temperature receiver and power
meter) (ﬁg. 1).
• 2009/Nov./24, Spectroscopic ﬁrst light of Hitachi antenna (tracking a methanol maser
source G9.62 using 6.7 GHz room-temperature
receiver and spectrum analyzer) (ﬁg. 2).
• 2010/Feb./25, First spectroscopic observation
with Hitachi antenna (tracking a methanol
maser source W3 using 6.7 GHz cryogenicallycooled receiver and sampler K5/VSSP32)(ﬁg. 3).

Figure 2. Hitachi spectroscopic first light.

December 201929

Figure 3. Hitachi first spectroscopic observation.

Figure 6. Damage by the earthquake.

Figure 4. Hitachi first fringe.

Figure 5. Takahagi first light.

Figure 7. History of back-end.

IVS NICT-TDC News No.38

30- News - News - News - News -Kashima 34-m Antenna–Damaged by Typhoon No.15 (Faxai), 2019–

Mamoru Sekido (sekido@nict.go.jp) , NICTSpace-Time Standards Laboratory,
Kashima Space Technology Center,893-1 Hirai, Kashima Ibaraki, Japan
1.

Kashima 34-m Antenna and its Contributions

Kashima 34-m antenna was built in 1988 for”Western Paciﬁc VLBI Network” project[1], whose
purpose is measurement of paciﬁc plate motion andregional crustal movement around Japan.
This antenna was designed as dedicated VLBI
station with multiple receiver systems. With thescope of not only geodetic VLBI but also for astronomical VLBI observations and R&D to explore the application of radio astronomical signal,
many receiver systems were implemented at initial state: 1.35-1.75GHz, 2.1-2.35GHz, 4.6-5.1GHz,
8.18-8.6GHz, 7.86-8.36GHz, 10.2-10.7GHz, 14.414.9GHz, 14.9-15.4GHz, 21.88-22.38GHz, 23.5824.08GHz, and 32.8-43.4GHz.
After accomplishment[3] of initial aim of the antenna installation, The Kashima 34-m antenna hascontributed to many projects for radio astronomy
and its applications. That include research on pulsar timing[4, 5] and astrometry[6], water maser
survey[7], interplanetary scintillation observation
to measure the solar wind speed[12, 13], Jupiter’s
synchrotron radiation[10], and bistatic Rader observation of asteroids [8, 9]. To memorize the ﬁrst
success of this Rader observation, the target asteroid was named GOLEVKA, which came fromthe participated station’s names GOLdstone, EVpatoria, and KAshima. The antenna played important role to support the ﬁrst space VLBI mission
VSOP[11], and participated many other astronomical observations, but we cannot introduce all ofthem here because of limited space of this article.
Kashima 34-m antenna has participated geodetic VLBI sessions conducted by IVS(International
VLBI Service for Geodesy & Astrometry),
AOV(Asia-Oceania VLBI Group for Geodesy andAstrometry), and domestic geodesy sessions conducted by GSI(Geospatial Information Authority of Japan). In addition, Kashima 34-m has
played important role in R&D researches for VLBI
technology developments: such as development of
PC-based VLBI data acquisition and correlation
processing system K5[14], rapid UT1 determination by combination use of the K5 and highspeed
internet[16, 17], and contribution to VSI-H and

Figure 1. Kashima 34-m station (top) and its receiver room(bottom).
VDIF standards[15]. High sensitivity Kashima 34m allowed to achieve these developments. In recent
years, Kashima 34-m was upgraded as broadband
antenna with newly developed broadband feed system. And it has been used for frequency comparison between optical time standards between INRiM(Italy) and NICT(Japan). The experiments
were successful, and we expected to continue further sessions in winter of 2019.
2.

Damage by Typhoon No.15 (Faxai)

Damages of typhoons becomes terrible year by
year. We have prepared for strong wind by repair of main reﬂector backup-structure not to makeany damages to surroundings, because it is locatedin residential area. However, the typhoon No.15(Faxai) arrived on 9th Sep., 2019 was strongest
ever experienced at our site. Maximum windspeedrecord was 56 m/s in station’s weather data. It was
fortune that there was no damage to the surrounding houses.
In the evening before typhoon’s arrival, antenna
was at zenith position as usual to prepare for strongwind. Although in the morning just after the ty-

December 201931

Figure 2. (a)Kashima 34-m antenna damaged in the morning of 9th Sep., 2019, (b)Main reﬂector panes
were removed in the same month. (c) Picture as of December, 2019, (d) Damages at elevation driving
motor, antenna structure, and working stage.
phoon has passed, we found that the main dish of [4] Hanado, Y., et al., Publ. Astron. Soc. Japan,
54, 305-313, 2002.
the antenna inclined at elevation angle about 120
degrees, which moved over elevation angle limit of
90 deg. Then it was found that elevation driving [5] Ilyasov, Yu. P., et al., Astrn. Lett., 31, 30-36,
2005.
motor, gearbox, and working stage were smashed.
Because elevation motion was free and nothing sup[6] Sekido, M. et al., Publ. Astron. Soc. Japan,
ported elevation angle motion, we had to ﬁx the51, 595-601, 1999.
elevation position by welding a thick steel plate immediately for safety. And ﬁnally, we removed the [7] Takaba, H., et al., Publ. Astron. Soc. Japan,
main reﬂector panels before arriving next coming
53, pp.517-533, 2001.
typhoon.
This is the story of how the Kashima 34-m [8] Hudson R. S., et al., Icarus, 148, 37-51, 2000.
antenna was damaged, and why its shape was[9] Koyama, Y., et al., “Rader Observations ofchanged. Although the antenna operation stopped
Near Earth Asteroids 6489 Golevka and 4197
due to the damage of typhoon, closing of operation(1982 TA)”, J. Comm. Res. Lab., Vol.48, No.1,
and dismantle of the antenna in 2020 has been de143-150., 2001.
cided before the disaster. Reason was increase ofmaintenance cost. It looks that Kashima 34-m an- [10] Miyoshi Y., et al., Geophys. Res. Lett., 26,
tenna has gave glory to the last by herself before
9-12, 1999.
dismantling.
[11] Hirabayashi, H., et al., Publ. Astron. Soc.
Acknowledgements
Japan, 52, 997-1014, 2000
We wish thank to all the people supported theKashima 34-m VLBI station. Especially, we thankto Y.Mizuno, Y.Watanabe, Y.Yamazaki, S.Suzuki,
Y.Hirose of Rikei Co. Ltd, K.Ishizu, K.Nakazono,
K.Nakazono, of Nakazono Co Ltd., H.Sakuyama,
R.Suzuki of A.E.S Co Ltd., N.Kurihara, E.Kawai,
J.Nakajima H.Takefuchi, H,Ohkubo of NICT, for
direct involvement in maintenance, for keeping running, and for improving the Kashima 34-m antenna.
References[1] Special Issue ”Western Paciﬁc Network”, J.
Commun. Res. Lab., Vol.42, No.1, pp.3-110,
1995.
[2] Imae M., et al., “The Main VLBI Station
at Kashima”, J. Commun. Res. Lab., Vol.42,
No.1, pp.5-14, 1995.
[3] Koyama Y., J. Geod. Soc., Japan, 42, 1, 4357, 1996.

[12] Tokumaru M., et al., J. Geomag. Geolelectr.,
43, 619-630, 1991.
[13] Tokumaru M., et al., J. Geomag. Geolelectr.,
46, 835-849, 1994.
[14] Kondo, T., et al., J. Geod. Soc. Japan, 54, 4,
233-248, 2008.
[15] Whitney, A., et al., Proc. IVS 2010 General
Meeting, 192-196, 2010.
[16] Sekido, M., Earth Planets and Space, 60,
pp,865-870, 2008.
[17] Matsuzaka, S., et al., Proc. IVS 2010 General
Meeting, 143-147, 2010.

“IVS NICT Technology Development Center News” (IVS NICT-TDC News) published by the
National Institute of Information and Communications Technology (NICT) (former the Communications Research Laboratory (CRL)) is the continuation of “IVS CRL Technology Development
Center News” (IVS CRL-TDC News). (On April 1, 2004, Communications Research Laboratory(CRL) and Telecommunications Advancement Organization of JAPAN (TAO) were reorganized
as “National Institute of Information and Communications Technology (NICT)”.)
VLBI Technology Development Center (TDC) at NICT is supposed
1) to develop new observation techniques and new systems for advanced Earth’s rotation observations by VLBI and other space techniques,2) to promote research in Earth rotation using VLBI,
3) to distribute new VLBI technology,4) to contribute the standardization of VLBI interface, and5) to deploy the real-time VLBI technique.
The NICT TDC newsletter (IVS NICT-TDC News) is published on a non-regular basis by NICT.

This news was edited by Mamoru SEKIDO, Kashima Space Technology Center. Inquires on thisissue should be addressed to Mamoru SEKIDO, Kashima Space Technology Center, National
Institute of Information and Communications Technology, 893-1 Hirai, Kashima, Ibaraki 3148501, Japan, e-mail : sekido@nict.go.jp.
Summaries of VLBI and related activities at the National Institute of Information and Communications Technology are on the Web. The URL to view the home page of the Radio Astronomy
Applications Section of the Space-Time Measurement Group of Space-Time Standards Laboratory is : “http://www2.nict.go.jp/sts/stmg/ivstdc/news-index.html”.

IVS NICT TECHNOLOGY DEVELOPMENT CENTER NEWS No.38, December 2019
International VLBI Service for Geodesy and Astrometry
NICT Technology Development Center News
published by
National Institute of Information and Communications Technology, 4-2-1 Nukui-kita, Koganei,
Tokyo 184-8795, Japan