ISSN 1882-3432

CONTENTS
Proceedings of the 18th NICT TDC Symposium (Kashima, October 1, 2020)
KNIFE, Kashima Nobeyama InterFErometer
Makoto Miyoshi...... 3

Space-Time Measurements Research Inspired by Kashima VLBI Group
Mizuhiko Hosokawa...... 8

ALMA High Frequency Long Baseline Phase Correction Using Band-to-band . . . . . 11
(B2B) Phase Referencing
Yoshiharu Asaki and Luke T. Maud
Development of a 6.5-22.5 GHz Very Wide Band Feed Antenna Using a New . . . . . 15
Quadruple-Ridged Antenna for the Traditional Radio Telescopes
Yutaka Hasegawa*, Yasumasa Yamasaki, Hideo Ogawa, Taiki Kawakami, Yoshinori Yonekura, Kimihiro Kimura, Takuya Akahori, Masayuki Ishino, Yuki Kawahara
Development of Wideband Antenna
Hideki Ujihara. . . . . 18

Performance Survey of Superconductor Filter Introduced in Wideband Re- . . . . . 20
ceiver for VGOS of the Ishioka VLBI Station
Tomokazu Nakakuki, Haruka Ueshiba, Saho Matsumoto, Yu Takagi, Kyonosuke
Hayashi, Toru Yutsudo, Katsuhiro Mori, Tomokazu Kobayashi, and Mamoru
Sekido
New Calibration Method for a Radiometer Without Using Liquid Nitrogen . . . . . 23
Cooled Absorber
Noriyuki Kawaguchi, Yuichi Chikahiro, Kenichi Harada, and Kensuke Ozeki
Comparison of Atmospheric Delay Models (NMF, VMF1, and VMF3) in . . . . . 27
VLBI analysis
Mamoru Sekido and Monia Negusini

HINOTORI Status Report
Hiroshi Imai

. . . . . 31

On-the-Fly Interferometer Experiment with the Yamaguchi Interferometer
Kenta Fujisawa, Kotaro Niinuma, Masanori Akimoto, and Hideyuki Kobayashi. . . . . 34

Superconducting Wide-band BRF for Geodetic VLBI Observation with . . . . . 36
VGOS Radio Telescope to Prevent Radio Frequency Interference
Ken Takahashi, Kazuhiro Tsubosaki, Hirotaka Kamide, Hiroshi Kudou, Taro
Asakura, and Masahiro Ruika
Status and Future of the Mizusawa 10m Radio Telescope
Osamu Kameya, Mareki Honma, Tomoaki Oyama, Kazuyoshi Asari, Aya Yamauchi, Yu Asakura, Gen Sato, Yuki Matsukawa, Kazuya Hachisuka, Shunsaku Suzuki, Toshio Terasawa, Kazuyoshi Yamashita, VERA development group,
VERA operations group, VERA maintenance group, JASMINE group, Balloon
VLBI group, Takeshi Fukusako, and Shotaro Haruguchi. . . . . 38

- News - News - News - News Kashima 34-m Antenna Closing Ceremony
Mamoru Sekido

2. . . . . 41

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3

KNIFE, Kashima Nobeyama the ﬁrst KNIFE experiment on June 14-15, 1989
failed, at the second KNIFE experiment on OcInterFErometer
tober 20-22, 1989, the ﬁrst 43 GHz fringes were
Makoto Miyoshi (makoto.miyoshi@nao.ac.jp)
National Astronomical Observatory, Japan,
2-21-1, Osawa, Mitaka, Tokyo, Japan, 181-8588

successfully detected from the observations of thestrong SiO masers in Orion KL and VY CMa.
2.

SiO maser observations

The SiO maser emissions originate at circumstellar envelopes of late-type stars, but they were mysAbstract: By connecting two antennas, Kashimaterious phenomena at that time. VLBI observa34 m and Nobeyama 45 m, an east-west basetions of SiO masers had been carried out, but theline of 200 km is formed. At that time, because
angular size of SiO maser was quite large, and thenNobeyama 45 m had the world’s number one sensiresolve-outs were occurred at long baselines, so thetivity in the 43 GHz band, and also Kashima 34 m
structures of SiO masers were not well understood.
was the world’s third-largest one, the KashimaThe KNIFE baseline is 200 km. At 43 GHz, theNobeyama baseline provided the highest sensitivminimum fringe spacing is about 7 mas, so it is
ity at 43 GHz VLBI (Figure 1). The construcjust suitable for observing SiO masers. Prof. Motion of the Kashima 34 m antenna began in 1988,
rimoto urged the install of the 43 GHz receiver on
also almost at the same time, a domestic project
the Kashima 34 m antenna as soon as possible beof mm-VLBI (KNIFE, Kashima Nobeyama INtercause fruitful scientiﬁc results were expected.
Frermeter) started. Nobeyama Radio Observatory
The Figure 4 shows the observational result ofprovided the ﬁrst cooled-HEMT 43 GHz receiver
SiO maser by KNIFE. The star itself cannot be dein the world to the Kashima 34 m. In October
tected, but the SiO masers are distributed around
1989, the ﬁrst fringe at 43 GHz was detected. We
the star. The distributions of SiO masers withhere review the achievements of the KNIFE at thatv = 1 and v = 2, though their excitation temtime.
peratures diﬀer about 1800 K, are almost the sameas each other. This means that the maser is in1. KNIFE project
ferred to be due to collisional excitation [1]. At
In 1987, the Radio Research Laboratory (RRL, that time, the VLBA in USA was in the process
NICT at present) decided to construct the Kashima of being constructed, but SiO maser observations
34 m antenna as the main station of the West- were performed as a test. Panel (e) in the Figure 4
ern Paciﬁc interferometer. At that time, the shows the VLBA result (SiO masers in U Her [2]).
Nobeyama Radio Observatory, a branch of the Na- If compared to the result of SiO masers by thetional Astronomical Observatory of Japan (NAOJ) VLBA four stations, It can be understood that theparticipated in the global millimeter-wave VLBI KNIFE performance was considerably good evenusing the Nobeyama 45 m and had just started though it was a single baseline.
VLBI observations. A set of Mark-3 recorder was
brought to the Usuda station from Nobeyama, and 3. Observations of high velocity watermega maser in NGC 4258
the ﬁrst space VLBI experiment with the TDRS
satellite was conducted using Usuda 64 m, and the
In 1992, Nakai et al. discovered high-velocity
fringe detection was successful. However, in Japan,
components in the water mega maser of the galaxy
independent astronomical VLBI observational reNGC 4258 by using the Nobeyama 45 m insearch has not yet been possible. In response to theJapan [3]. Relative to the known water maser comnews of the construction of the Kashima 34 m anponents, they show velocity shifts of ±1000 km/s.
tenna, Prof. Morimoto (Figure 2) noticed that the
There are two groups, one shows blue shifts of
Kashima 34 m antenna has a surface accuracy of
about 1000 km/s while the other shows red shifts
170 μas and that it is very eﬀective for millimeterof about 1000 km/s. Hearing the news, Prof. Mowave VLBI observation. Prof. Morimoto proposed
rimoto said, ”If the masers move with velocities
to RRL to conduct the millimeter-wave VLBI reabout 1000 km/s, we can easily detect the proper
searches in collaboration with NAOJ, and the joint
motions of the maser components. Observe withresearch began. Using the 43 GHz cooling receiver
Kashima 34 m immediately! Observe twice so as
dewar owned by RRL, NAOJ decided to manufacnot to fail!” What is the origin of the high-velocity
ture the world’s ﬁrst 43 GHz cooled HEMT receivermega masers of NGC 4258? In order to identify(Figure 3) and the joint research started in 1989.
the origin, observations with a high spatial resThe KNIFE experiment started at the same time as
olution are required. Prof. Morimoto suggested
the start-up and test of the 34 m antenna. Thoughthat we should investigate with KNIFE as soon as4

IVS NICT-TDC News No.39

Figure 1. KNIFE baseline: The baseline length is 197.66 km. The north-south component is only 240 m.
The fringe spacing is about 7 mas at 43 GHz observations, suitable for SiO maser observations.

Figure 2. Prof. Masaki Morimoto, who promoted Figure 3. The cooled HEMT Receiver installed to
the KNIFE project.
34 m telescope. This photo was taken at the time
when the ﬁrst 43 GHz fringe was detected in Oct.
1989.
possible. The KNIFE observations were conductedin early June 1993, less than a month after the
discovery. However, due to calculation errors of capture the structure of the water mega maser inthe observational sky frequencies, correlations by NGC 4258 even with the single 200 km baseline of
new NAOCO correlator in NAOJ [4] that has not KNIFE. The KNIFE observation of NGC 4258 also
even been bug-ﬁxed, and the use of imaging soft- shows the excellent performance of the KNIFE.
ware made tentatively, it was not possible to ﬁnd
the maser distributions exactly. However, in Au- 4. The role played by the KNIFE projectgust 1993, it is found that the high-velocity comThe results are not only those mentioned above.
ponents are also located at the center of the galaxy, Pioneering research was conducted on higherwithin 50 mas of the known main components. In
frequency geodetic VLBI using KNIFE [7]. In addi1995, VLBA observations revealed that the watertion, survey observations of water and SiO masers
mega maser in NGC 4258 was from a molecular were carried out using the single Kashima 34 m angas disk with Keplerian motion around a supertenna and KNIFE [8], [9]. The KNIFE astronommassive black hole. From the maser velocity and ical observations using the single-baseline between
structure, the mass of the black hole is estimated to Kashima 34 m and Nobeyama 45 m were carried
be 3.6 × 107 M [5]. The Figure 5 shows the megaout until 1992. In 1993 it developed into J-Net
maser distribution obtained from reanalyzing the observations with 4 stations including Mizusawa
KNIFE data and that from the result of VLBA.
10 m antenna and Kagoshima 6 m antenna; MoniAs you can see from the Figure, it is possible totoring VLBI observations of the burst phenomena

January 2021
of water masers in Orion KL were performed before its beginning by chance [10]. KNIFE paved
the way for the subsequent development of astronomical VLBI researches in Japan.
References[1] Miyoshi, M., Matumoto, K., Kameno, S., Takaba, H., & Iwata, T., Collisional Pumping of
SiO Masers in Evolved Stars, Nature, 371,
395-397, 1994.
[2] Diamond, P. J., Kemball, A. J., Junor, W.,
Zensus, A., Benson, J., & Dhawan, V., Observation of a Ring Structure in SiO Maser
Emission from Late-Type Stars, Astrophysical Journal Letters v.430, p.L61, 1994.
[3] Nakai, N., Inoue, M., & Miyoshi, M.,
Extremely-high-velocity H2 O maser emission
in the galaxy NGC 4258, Nature, 361, 45-47,
1993.
[4] Shibata, K., Sasao, T., Kawaguchi, N.,
Tamura, Y., Kameno, S., Miyoshi, M.,
Asari, K., Manabe, S., Hara, T., Kuji, S.,
Sato, K., Miyaji, M., Matumoto, K.,
Asaki, Y., Yasuda, S., & Nakamura, S., A
New Advanced One-Unit VLBI Correlator(NAOCO), Proceedings of the International
Symposium Held at Kyoto International Conference Hall on September 6-10, 1993 ’VLBITECHNOLOGY Progress and Future Observational Possibilities’ 327-331, Sasao, T., Manabe, S., Kameya, O. & Inoue, M. eds. Terra
Scientiﬁc Publishing Company, 1994.
[5] Miyoshi, M., Moran, J., Herrnstein, J., Greenhill, L., Nakai, N., Diamond, P., & Inoue, M.,
Evidence for a Black Hole from High Rotation Velocities in a Sub-Parsec Region of
NGC 4258, Nature, 373, 127-129, 1995.
[6] Herrnstein, J.R., Moran, J., Greenhill, L., Diamond, P., Inoue, M., Nakai, N., Miyoshi, M.,
Henkel, C.,& Riess,A., A geometric distanceto the galaxy NGC 4258 from orbital motions
in a nuclear gas disk, Nature, 400, 539-541,
1999.
[7] Matumoto, K., Kawaguchi, N., Inoue, M.,
Miyoshi, M., Kameno, S., Takaba, H.,
Iwata, T., & Kurihara, N., A Radio Source
Survey at 22 and 43 GHz for Geodetic VLBI,
Journal of the Geodetic Society of Japan, 40,
255-265, 1994
[8] Takaba, H., Ukita, N., Miyaji, T., &Miyoshi, M., Spectral Evolution of H2 O Maser

5in Late Type Stars, Publ. Astron. Soc. Japan,
46, 1-14, 1994.
[9] Imai, H., Miyoshi, M., Ukita, N., Morimoto, M., Iwata, T., Takaba, H., Kameno, S., Matumoto, K., Fujisawa, K., &Kawaguchi, N., Snapshot VLBI Observations of SiO v = 1 and v = 2 (J=1-0) Maser
Sources. Publ. Astron. Soc. Japan, 53, 259273, 2001.
[10] Omodaka, T., Maeda, T., Miyoshi, M.,
Okudaira, A., Nishio, M., Miyaji, M.,
Mochizuki, N., Morimoto, M., Kobayashi, H.,& Sasao, T., The Enormous Outburst of the
7.9km/sec Water-Maser Feature in Orion KL,
Publ. Astron. Soc. Japan, 51, 333-336, 1999.

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IVS NICT-TDC News No.39

Figure 4. SiO masers observed with KNIFE. (a) VY CMa J=1-0, v=1, (b) VY CMa J=1-0, v=2, (c)W Hya J=1-0, v=1, (d) W Hya J=1-0, v=2, (e) U Her observed with VLBA [2], (f ) u-v coverage of
KNIFE baseline for them.

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Figure 5. The oldest map of high velocity water mases in NGC 4258 observed with KNIFE in 1992
June (top panel). The image was produced in 1997. This map demonstrates that the KNIFE had theperformance to clarify the structure of water mega masers in NGC 4258. The bottom panel shows theimage with VLBA observations [6].

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IVS NICT-TDC News No.39

Space-Time Measurements
Research Inspired by Kashima
VLBI Group
Mizuhiko Hosokawa(hosokawa@nict.go.jp )Principal Reseacher,
NICT, 4-2-1 Nukui-Kita-Machi, Koganei, Tokyo
194-8795, Japan
Abstract: I have done some works on the relativistic
eﬀects and precise space-time measurements, suchas Time Geostationary Orbis, stellar mass measurements by using astrometric gravitational lensing, and positional ﬂuctuation of quasars causedby gravitational lensing stars in our galaxy. Mostof them are inspired by the members and legendsof Kashima VLBI group. Here I introduce thesememories.
1.

Introduction

Kashima VLBI group is famous on the development of K-series VLBI systems and on thelots of applied research using large radio antennae. Kawajiri-san, Kawano-san, TakahashiFujinobu-san, and Kawaguchi-san are some of thevery early member of this group, followed by many
researchers, such as Yoshino-san, Takahashi Yukiosan, whom I mentioned here, and many excellent
researchers. Encounters with them changed my research life. In this article, I would like to introduce how they inspired me on some of my researchworks.

Figure 1. Time geostationary orbits around Earth,
Sun and Jupiter.

an idea on one orbit seems not so much for an article. So I considered a bit more and got an idea toseek orbits in the solar system whose proper times
are the same as that on the Earth geoid surface insimple circular orbit approximation. After a few
2. Takahashi Fujinobu-sancalculations, I found two more such orbits, one is
I joined Communications Research Laboratories that around Sun, a little bit inside the Earth or(CRL) in 1990. A few months after I joined CRL, bit, and the other is that around Jupiter. They are
I had a chance to listen to a talk on the proper shown in Fig.1. I brought these results to Fujinobutime of an atomic clock onboard satellite by Taka- san, had some discussion with him and presented
hashi Fujinobu-san. In the talk, he said his group a paper [1].
was learning General Relativity to calculate theclock’s proper time. I said to him that I can help 3. Kawaguchi-sansuch calculation, as in the graduate school I had
This research started my join to the space-timelearned General Relativity. Soon after hearing his measurements project in CRL, that include thetalk, I tried to calculate a proper time of a space VLBI group. In 1991 and 1992, I gave some preclock in the simplest circular orbit case by using sentation to the Japan astronomical society annual
Schwarzschild metric solution. Then, I found an in- meeting on the proper time of space crafts, includteresting case. In a special orbit radius, the proper ing Time Geostationary Orbits. At that time, sometime of the space clock become just same as that feasibility studies on space-VLBI using Japaneseof the clock on the Earth geoid surface. I told Engineering Test Satellite VI. I considered the relFujinobu-san that result. He said he had never ativistic eﬀect on ETS-VI, including frequency shift
heard about such orbit. Then I named this orbit with ground station. Kawaguchi-san heard my preas “Time Geo-stationary orbit”. I also thought we sentation and interested in that. He said that heshould write a paper on this result. However, just wished to visit me to discuss on the relativistic

January 2021

9

Figure 3. Apparent motion of the lensed image S
caused by the proper motion of a lensing star P.
Figure 2. Parallactic variation of gravitational
lensing and measurement of stellar mass.

In 1997, he visited CRL with Hanada-san in NAO.
He had an idea which he called inverse VLBI. He
eﬀect for space VLBI project. One or two years wanted Imae-san in CRL to join the system delater, when he engaged in the HALCA space VLBI sign investigation of this idea. Imae-san asked me
project, he visited my oﬃce. I thought that his to join this research. So I became a member of
purpose should be to hear from me the relativistic this investigation. His idea was that if we had two
eﬀects for satellites. However, what he said was “In transmitters with atomic clocks on Mars stations,order to consider the frequency transmission sys- and had synchronized the clocks by two way sateltem between ground station and HALCA, I studied lite time transfer via Mars GEO satellite, we canGeneral Relativity and did some calculations. Are measure many geodetic phenomena precisely by rethese results coincide with your result in ETS-VI?” ceiving the transmitters’ signals with one antenna
Yes, he already learned General Relativity and ﬁn- on an Earth station. After long discussions, thisished the calculations. His results agreed with my work was published in 1999 [3]. It is good memory
calculations. I had heard Kawaguchi-san was pow- for me to collaborate with one of Japanese VLBI
erful and great. At that time, I really understood legends.
that reputation.
5. Takahashi Yukio-san and ATNF
4. Kawano-san
In 1993, Takahashi Yukio-san had stayed at AusIn late 1990 and early 1991, I had some chances tralian National Telescope Facility (ATNF) for one
to listen to the plan of VERA project, which will year. I envied him and wished that someday I could
realize the 10 micro-arc second position measure- stay at ATNF, too. In 1996, I got a chance toment precision by using phase referenced VLBI apply to Japan-Australia researcher exchange protechnique to measure the scale of our galaxy. The gram. On December 1996, an international workkey persons to this project were Sasao-san and shop called TWAA was held in Kashima. There
Kawano-san. At that time, Kawano-san was al- I gave a presentation on the ﬂuctuation of celesready moved to National Astronomical Observa- tial reference frame due to the gravitational lenstory. But I have heard Kawano-san was once a ing caused by stars in our galaxy [3]. It was one ofmember, and a leader, of Kashima VLBI group. the works under the collaboration with FukushimaThat made me easy to talk with Kawano-san. To san and Ohnishi-san. In this workshop, Dr D.
hear this very ambitious plan, I was very much as- Jauncey at ATNF attended. After my talk, I asked
tonished, and thought that such high precision can Dr Jauncey that I was planning to apply Japanbe used to many other observations. An idea I Australia researcher exchange program and if I segot was, to measure the stellar mass using gravita- lected as the program member, I wanted ATNF totional lensing eﬀect. Fukushima-san and Ohnishi- accept my stay. Dr Jauncey assured me that ATNF
san were interested in this idea and we started would accept me when the program committee seour collaboration. After many discussions, we pre- lected me as the member. His words encouraged me
sented a paper [2] as our ﬁrst collaboration work. and I was selected as the program member 1998.
An illustration of the paper is shown in Fig.2. This My stay at ATNF is ﬁrst for me, but for ATNF, I
work was thus inspired by the precision of the phase was not the ﬁrst visitor from CRL. So many procereferenced VLBI technique proposed by Sasao-san, dures went smoothly. At the begging of my stay, I
Kawano-san and many radio astronomers. I had had a chance to talk with Prof. R. Ekers, the Direca chance to present some work with Kawano-san. tor of ATNF at that time. Hearing a summary of

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IVS NICT-TDC News No.39

what I did in CRL, he gave me an excellent suggestion. “There are three Quasars in the direction of
Galactic Center, within one degree. It may be interesting to investigate the gravitational lens eﬀect
for these quasars.” I appreciated this good suggestion very much. During my stay, with the help of
ATNF VLBI group, I considered on this suggestion.
I also discussed with Japanese colleagues, since it
was already easy to communicate with Japanesecolleagues with email at that time. Applying theidea of apparent motion of the quasars induced by
the stars in our galaxy, used in [4] and shown in
Fig.3, we wrote a paper whose title is “Possible
ﬂuctuation of the position of Sagitarius A* relative
to extragalactic radio sources” [5]. Publication ofthis paper was a token that my stay at ATNF wasvery fruitful.
6.

Yoshino-san

When I joined CRL in 1990, Yoshino-san wasa senior researcher of Fujinobu-san’s laboratory.
In those days, when I had some question about
VLBI and space-time measurements, I often asked
him. And also, every time when I got some idea
on space-time measurement, I told him and asked
his opinion about that. His comment was, in mostof the time, in a soft tone, but severe. With myidea. Many time he told me “It may be interesting. But it seems not yet so realistic.” After I
came back from ATNF, I told Yoshino-san about
the idea of ﬂuctuation of Sagitarius A* relative toquasars beyond galactic center. At that time, hesaid that it seems much realistic idea than ever. I
was very glad to hear his words, as he at last ﬁndsome value in my new idea. His warm and severe
comments has been giving encouragement to me.
7.

Summary

Meeting with the legends and members of
Kashima VLBI group has been given me lots ofidea and has changed my research life very fruitfully. I deeply appreciate the group. Though 34
m antenna operation has stopped, I wish the successors of the group will open up the new research
frontier.
References[1] Hosokawa, M. and Takahashi, F., “TimeGeostationary Orbits in the Solar System”,
Publ. Astron. Soc. Japan, 44, 1992, pp.159162.
[2] Hosokawa, M., Ohnishi, K., Fukushima, T.,& Takeuti, M. “Parallactic variation of gravitational lensing and measurement of stellarmass”, A&A, 278, L27, 1993.

[3] Kawano, N., Hosokawa, M., Hanada, H. &Imae, M., “Inverse VLBI method for Planetodesy”, J. of the Geodetic Society of Japan.
45, p181, 1999.
[4] Hosokawa, M., Ohnishi, K., & Fukushima,
T. “Fluctuation of Extragalactic Reference
Frame Due to Gravitational Lensing in Our
Galaxy”, Astronomical Journal, v. 114, p.
1508, 1997.
[5] Hosokawa, M., Jauncey, D., Reynolds, J.,
Tzioumis, A., Ohnishi, K., & Fukushima,
T. “Possible Fluctuation of the Position of
Sagittarius A* Relative to Extragalactic Radio Sources”, APJL, 580, L43, 2002.

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11

ALMA High Frequency Long To reduce the atmospheric phase ﬂuctuations,referencing is a frequently adopted techBaseline Phase Correction Us- phasenique for interferometer phase correction. The
ing Band-to-band (B2B) Phase general implementation of phase referencing is tofrequently visit a phase calibrator, typically a
Referencing
bright point-like quasar (QSO), close to the tarYoshiharu Asaki,
Chile Observatory, National Astronomical
Observatory of Japan,
Alonso de Córdova 3788, Oﬃce 61B, Vitacura,Santiago, Chile
Luke T. MaudESO Headquarters, Karl-Schwarzchild-Str. 2,
D-85748 Garching, Germany
Abstract: In 2017, an Atacama Large Millimeter/submillimeter Array (ALMA) high-frequency
long baseline campaign was organized to test image capabilities with baselines up to 16 km at submillimeter (submm) wavelengths. We investigatedimage qualities using ALMA receiver Bands 7, 8, 9,
and 10 (285–950 GHz) by adopting band-to-band(B2B) phase referencing in which a phase calibrator is tracked at a lower frequency. In this paper wereport on QSO observation experiments at Band 7(289 GHz) and Band 8 (405 GHz) using B2B phasecorrection with a phase calibrator 0◦ .7 away observed in Band 3 (96 GHz) and Band 4 (135 GHz),
respectively. The novel B2B phase referencing successfully correct the high frequency interferometer
phases, achieving the angular resolutions of 15 and11 mas at Bands 7 and 8, respectively.
1.

Introduction

The Atacama Large Millimeter/submillimeter
Array (ALMA) is a very powerful instrument toinvestigate emissions in millimeter/submillimeterwaves with very high angular resolutions. With the
16 km ALMA conﬁguration, the spatial resolutions
will be 12, 7, and 5 milliarcseconds (mas) at the observing frequencies of 400, 650, and 950 GHz, respectively. The long-baseline capabilities of ALMA
have been tested since ALMA early science operation started in 2011 [1–5]. ALMA has opened
16 km baseline observations to the user community up to Band 7 frequencies in Cycle 7 (≥ 0.8 mm
wavelength, or ≤ 373 GHz), with the ability to regularly achieve resolutions higher than ∼ 20 mas.
To achieve the highest angular resolutions with
ALMA, interferometer phase stability for the
16 km baselines in high-frequency (HF) bands(Bands 8–10, or the frequency range of 385–
950 GHz) is crucial. The longer the baseline, theworse the phase stability limited by atmosphericphase ﬂuctuations, and thus the synthesized image
becomes blurred [6].

get source at the same frequency. However, theﬂux density of a typical QSO decreases in the millimeter/submillimeter waves and the system noise
increases, meaning that the integration time toachieve suﬃcient signal-to-noise ratio (S/N) becomes too large. This combination makes it unlikely to ﬁnd a bright enough nearby phase calibrator for a randomly located target.
In order to mitigate the diﬃculty in ﬁndinga phase calibrator at HF, the so-called band-toband (B2B) phase referencing provides an alternative technique, employing observations of a sufﬁciently close phase calibrator at a low-frequency(LF) Band [7–9].
B2B phase referencing is currently being implemented at ALMA. The motivation of the researchis to verify the technical feasibility of B2B phasereferencing and evaluate its performance in synthesized images with the high angular resolutions
in ALMA. We report on HF long-baseline (14–16 km) image capability tests for a point source
QSO J2228−0753 in Band 7 (289 GHz) and Band 8(405 GHz) observed together with a closely locatedphase calibrator, J2229−0832, in Band 3 (96 GHz)
and Band 4 (135 GHz), respectively, to demonstrate B2B phase referencing. In Section 2, thebasic strategy of B2B phase referencing for HF
long-baseline observations is mentioned. In Sections 3, we outline our ALMA B2B phase referencing observation experiments. Section 4 provides
the results of the experiments and discussions. The
ALMA high frequency long baseline campaign arepresented in detail by other papers [10–12].
2.

Basic strategy of ALMA high frequency
phase correction

The main goal of the interferometer phase correction is to remove systematic phase errors due
to instrumental electrical path length errors and apriori antenna position errors, atmospheric delay
models over the array, and short-term atmosphericdelay variations. An important delay correction
produced by water vapor over each 12 m antenna
is obtained by water vapor radiometers (WVR) on
each antenna [13]. The WVR phase correction isalways applied to the ALMA data and typically removes the majority of the water vapor delay [14],although it is somewhat dependent on the quantity of the water vapor content in the atmosphereas precipitable water vapor (PWV), the depth of

12

the water in a column of the atmosphere to thezenith.
In 2017, HF-long baseline campaign was organized to test the submillimeter-wave (wavelength ofshorter than 1.1 mm) imaging capability of ALMA
using the longest baselines up to 16 km. One focus of the campaign was to extensively test theB2B phase referencing method and begin providing
appropriate observing scripts and data reduction
scripts. In B2B phase referencing, a close phasecalibrator is observed at an LF. The LF phase of
the phase calibrator is multiplied by the HF/LF
ratio.
An important observation parameter in phasereferencing is the switching cycle time tswt that ismeasured by a time diﬀerence between two calibrator scans. Our experiments adopted a short switching cycle time of 20 s because HF observations may
require more frequent visits to a phase calibrator
to compensate for the atmospheric phase ﬂuctuations. The ALMA antennas can quickly change
their position by several degrees in a few seconds
to accommodate such a quick source change.
Use of B2B phase referencing requires an additional phase correction to remove an instrumental
phase oﬀset diﬀerence between the two respective
Bands. This can be achieved through a cross-band
calibration, which is referred to as diﬀerential gain
calibration (DGC). In the DGC, a bright QSO is
observed at both the HF and LF, using frequency
switching, such that the phase oﬀset diﬀerence canbe solved.

IVS NICT-TDC News No.39

by internal electronic phase diﬀerences between the
HF and LF [10]. We switch between the HF and
LF with only a ∼ 2 s delay. We used DGC scans
for a bandpass calibration at both the LF and HF,
as well as the ﬂux calibration.
A single observation consists of repetitions of theB2B phase referencing block for the target andphase calibrator and regular visits to the DGC
block where we only switch the frequency while
pointing at the DGC source. The B2B phase referencing block between the target and phase calibrator lasts 12 minutes, after which a 2.5 minute
DGC block is inserted. Such a frequent insertionof the DGC block was made not only to checkthe DGC solution repeatability and stability but
also to avoid a heavy load to the antenna control
computers by a long sequence of antenna pointingchanges.
Each Band ampliﬁes two linear polarizations (Xand Y ) separately at each 12 m antenna. The
ampliﬁed signal is split into four intermediatefrequency signals, which are referred to as basebands (BBs) with the bandwidth of 2 GHz. In our
experiments, the digitized BBs were ﬁltered out tohave a 1.875 GHz bandwidth each with 128 frequency channels for two polarization pairs of XX
and Y Y to form a spectral window. The integration time used for recording each of the four spectral windows at both the HF and the LF respectively was 1.01 s.
For the B2B phase referencing blocks, a typicalscience target scan length was 8 s at the HF, while
the phase calibrator scan length at the LF was 6 s.
3. Experiments
For the DGC blocks, the above frequency switchingsequence was adopted, and the scan length for the
We conducted the QSO observing experiments
HF and LF DGC scan was typically 8 and 6 s, rein 2017 October and November with the maximum
spectively. The antenna slew and/or the frequency
baseline lengths of 16 and 14 km, respectively, withswitch are done simultaneously and take 2–3 s. The
40–50 12 m antennas. Overall four experiments
total length for the one switching cycle is ∼ 20 s inwere conducted, two in October (9th and 10th) andboth the B2B phase referencing sequence and DGC
two in November (2nd and 3rd). We arranged B2B
blocks.
phase referencing experiments at two speciﬁc frequency pairs. One is at 289 GHz in Band 7 for
4. Results and Discussions
the target with the phase calibrator observed at
96 GHz in Band 3, while the other is at 405 GHz 4.1
Image coherencein Band 8 for the target with the phase calibraFigure 1 shows synthesized images of the target
tor at 135 GHz in Band 4. There are totally two
J2228−0753 at 289 GHz (Band 7) and at 405 GHz
frequency sets of two observations at Bands 7 and(Band 8) after B2B phase referencing. The highest
8.
angular resolution of 14 × 11 mas was achieved at
We observed the QSO J2228−0753 as a target at
Band 8 with a maximum projected baseline length
the HF, while the QSO J2229−0832, with a separaof ∼ 13 km. To verify the image quality using
tion angle of 0◦ .7, was observed as a phase calibraB2B phase referencing, phase self-calibration was
tor at the LF. For DGC, we selected a very brightadditionally conducted after B2B phase referencQSO J2253+1608 located 25◦ away from the taring. The peak ﬂux densities of the self-calibrated
get. Note that the phase oﬀset diﬀerence measured
images are consistent within 5% between the two
with a DGC source is independent of the QSO locaepochs in both Bands 7 and 8. Since a good amtion in the sky, since it is predominately determinedJanuary 2021
plitude repeatability is achieved in the synthesis
images by reducing the phase errors as much as
possible, the phase self-calibrated results can be
regarded as a reference.

13

so that in our tests B2B phase referencing with aclose phase calibrator and 20 s switching cycle time
proves to be an eﬀective phase correction schemeat the high frequencies.
4.2

Astrometry

Baseline vector errors cause a systematic phaseerror after B2B phase referencing, and thus cancause blurring of the target image. The positional
shift of the target Δθ (rad) due to a baseline error
vector Δρ (in meters) is expressed by ρ · Δθ =Δρ · (st − sc ), where ρ is a baseline vector andst and sc are unit vectors from an observer to thetarget and phase calibrator in the sky, respectively.
In ALMA, Δρ is 2.5 mm for the longest baseline[15].
We investigated the target peak position in
the synthesized images before self-calibration
using a two-dimensional Gaussian ﬁtting in
the image plane.
The averaged positionof J2228−0753 for the two epochs and two
Bands is (αJ2000 , δJ2000 )=(22h 28m52s .607590,
with
a
1σ
error
of
−7◦ 53 46 .64238)(0.24, 0.40) mas.
The measured position of
J2228−0753 is consistent with that in the
ALMA calibrator catalogue within the quoted
uncertainties. In centimeter-wave observations,
J2228−0753’s position is reported in the VeryLong Baseline Array (VLBA) calibrator list1
with a 1σ error of (0.10, 0.21) mas. Our B2B
image position is also consistent with that in the
VLBA calibrator list within 0.4 mas, less than a20th of the Band 8 synthesized beam. For thetarget–phase calibrator separation of 0◦ .7, Δθ
is equivalent to ∼ 0.4 mas at most for a 16 km
baseline, so that the obtained target position has areasonable accuracy compared to the ALMA calibrator source catalogue and the VLBA calibrator
Figure 1.
CLEAN images of the target list.
J2228−0753 at Band 7 (top) and Band 8 (bottom)
with B2B phase referencing. The abscissa and or- We used the ALMA data ofdinate are R.A. and decl., respectively. The con- uid://A002/Xc5802b/X5bb3,
tours are drawn for a factor of −3 (dash line), 3, uid://A002/Xc59134/Xd47,
6, 12, 24, 48, 32, 64, and 128 (solid lines) multi- uid://A002/Xc65717/X56f,
andplied by the rms noise of the images. The synthe- uid://A002/Xc660ef/X8e0 for the research.
sized beam is shown in the lower left corner of each ALMA is a partnership of ESO (representing its
panel: 18 × 15 mas and 14 × 11 mas at Bands 7 member states), NSF (USA) and NINS (Japan),and 8, respectively.
together with NRC (Canada), NSC and ASIAA(Taiwan), and KASI (Republic of Korea), in
One of the important indicators to evaluate the cooperation with the Republic of Chile. The
image quality is an image coherence, which is the Joint ALMA Observatory is operated by ESO,
ratio of the image peak ﬂux density compared to AUI/NRAO, and NAOJ. The authors thank to
the true value. We assume that the phase self- all the Joint ALMA Observatory staﬀs in Chilecalibrated image represents the true value for each for performing the challenging HF-LBC-2017
experiment. The image coherence is larger than successfully. L. T. M. was adopted as a JAO
90% and 80% at 289 and 405 GHz, respectively,
1 http://www.vlba.nrao.edu/astro/calib/

14

ALMA expert visitor during his stay. This work
was supported by JSPS KAKENHI grant No.
JP16K05306.
References[1] Matsushita, S., Asaki, Y., Fomalont, E. B. etal., ALMA Long Baseline Campaigns: Phase
Characteristics of Atmosphere at Long Baselines in the Millimeter and Submillimeter
Wavelength, PASP, 129, 973, 035004, 2017.
[2] ALMA Partnership, Fomalont, E. B., Vlahakis, C. et al., The 2014 ALMA Long Baseline Campaign: An Overview, ApJL, 808, 1,
L1, 2015.
[3] ALMA Partnership, Hunter, T. R., Kneissl, R.
et al., The 2014 ALMA Long Baseline Campaign: Observations of Asteroid 3 Juno at 60
Kilometer Resolution, ApJL, 808, L2, 2015.

IVS NICT-TDC News No.39
ing in Submillimeter Waves, ApJS, 247, 23,
2020.
[11] Asaki, Y., Maud, L. T., Fomalont, E. B. etal., ALMA Band-to-band Phase Referencing:
Imaging Capabilities on Long Baselines andHigh Frequencies, AJ, 160, 59, 2020.
[12] Maud, L. T., Asaki, Y., Fomalont, E. B. et al.,
ALMA High-frequency Long-baseline Campaign in 2017: A Comparison of the Bandto-band and In-band Phase Calibration Techniques and Phase-calibrator Separation Angles, ApJS, 247, 23, 2020.
[13] Nikolic, B., Bolton, R. C., Graves, S. F. etal., Phase correction for ALMA with 183 GHz
water vapour radiometers, A&A, 552, A104,
2013.

[14] Maud, L. T., Tilanus, R. P. J., van Kempen, T. A. et al., Phase correction for ALMA.
[4] ALMA Partnership, Brogan, C. L., Pérez,
Investigating water vapour radiometer scalL. M. et al., The 2014 ALMA Long Baseline
ing: The long-baseline science veriﬁcation
Campaign: First Results from High Angular
data case study, A&A, 605, A121, 2017.
Resolution Observations toward the HL Tau[15] Hunter, T. R., Lucas, R., Broguiére, D. et
Region, ApJL, 808, 1, L3, 2015.
al., Analysis of antenna position measure[5] ALMA Partnership, Vlahakis, C., Hunter,ments and weather station network data durT. R. et al., The 2014 ALMA Longing the ALMA long baseline campaign of 2015,
Baseline Campaign: Observations of the
2016SPIE 9914, 99142L, 2016.
Strongly Lensed Submillimeter Galaxy HATLAS J090311.6+003906 at z = 3.042, ApJL,
808, L4, 2015.
[6] Carilli, C. L. & Holdaway, M. A., Troposphericphase calibration in millimeter interferometry,
RaSc, 34, pp.817-840, 1999.
[7] Asaki, Y., Shibata, K. M., Kawabe, R. etal., Phase compensation experiments with thepaired antennas method 2. Millimeter-wave
fringe correction using centimeter-wave reference, RaSc, 33, 5, pp.1297-1318, 1998.
[8] Rioja, M. J., Dodson, R., Jung, T.-H. et al.,Veriﬁcation of the Astrometric Performance ofthe Korean VLBI Network, Using Comparative SFPR Studies with the VLBA at 14/7mm, AJ, 148, 84, 2014.
[9] Pérez, L. M., Lamb, J. W., Woody, D. P.
et al., Atmospheric Phase Correction Using
CARMA-PACS: High Angular Resolution Observations of the FU Orionis Star PP 13S*,
ApJL, 724, pp.493-501, 2010.
[10] Asaki, Y., Maud, L. T., Fomalont, E. B. et al.,
ALMA High-frequency Long Baseline Campaign in 2017: Band-to-band Phase Referenc-

January 2021

15

Development of a 6.5-22.5 GHz of 66% on the fundamental transmission mode. In
practice, approximately 58% of the ALMA bandvery wide band feed antenna
7+8 ( 275–500 GHz ) is one of the maximum, andusing a new Quadruple-Ridged trying to broaden above 60 % is so diﬃcult on the
waveguide devices such as the Ortho–Mode TranAntenna for the traditional
scucer and Branch–Line Coupler. To solve thisproblem, antennas such as Eleven Feed[1],Quasi
radio telescopes
Yutaka Hasegawa∗1(s y.hasegawa@p.s.osakafu-u.ac.jp),
Yasumasa Yamasaki∗1 , Hideo Ogawa∗1 ,Taiki Kawakami∗2 , Yoshinori Yonekura∗2 ,Kimihiro Kimura∗3 , Takuya Akahori∗4,
Masayuki Ishino∗5 , Yuki Kawahara∗5*1 : Department of Physical Science, Graduate
School of Science, Osaka Prefecture University,
1-1, Gakuen-cho, Nakaku, Sakai, Osaka 599-8531,
Japan*2 : Ibaraki University*3 : ISAS/JAXA*4 : SKAO/NAOJ*5 : Kawashima Manufacturing Co.
Abstract: We have developed a new Quadruple–
Ridged Antenna to realize a very wideband feed
antenna, which is highly required in astronomicaland geodetic observations using radio telescopes at
microwave band. This new QRA has a very low
return loss of −19 dB or less over the entire 6.5–
22.5 GHz, and can provide a good radiation beam
pattern and sharp directivity, with small frequency
dependence on its beam size, so that it can be installed into the conventional radio telescopes.
1.

Introduction

One of the recent development issues in the radio telescope receiver systems is to broaden theirfractional bandwidth (deﬁned as its coverage bandwidth / its center frequency) to above 100%. Recent observation devices using the current mainstream waveguide front end, such as in ALMA,
which revolutionized radio astronomy, usually havea fractional bandwidth of approximately 30%. On
the other hand, in future large radio astronomy
projects such as SKA (MID) and ngVLA, which
will consist of hundreds of parabolic antennas, they
aim at a fractional bandwidth over 100%, and will
explore new science ﬁelds such as the epoch ofreionization, multi-messenger astronomy including
gravitational waves, and explorations of extraterrestrial life.

Self–Complementary Antenna[2], and Quadruple–
Ridged Antenna (QRA)[3, 4, 5] have been developed, which can be cover more than 100% of fractional bandwidth. However, these have a verystrong frequency dependence on its radiation beam
pattern, and matching with the optical transmission system of the conventional (existing) telescopecannot be ensured in the entire band. In addition,
its reﬂection amount (return loss) of the coaxial
feeding unit is so large. These problems have prevented the spread of high–performance and very
wideband receivers, and thus it becomes one ofthe development issues in recent years. In orderto solve these problems, we focused on the QRA,
which is considered to be compatible with the radio telescope feed, and we developed a completely
new QRA model that is improved its antenna characteristics (Return Loss, Beam Pattern, Frequency
Independency).
2.

Features of general QRAs

Unlike other well known self–complementary
wideband antennas, QRA is a super wideband
waveguide antenna extending the fundamental cutoﬀ frequency toward a lower frequency by the Ridge
Waveguide. The structure of QRA is an open
waveguide antenna of Quadruple-Ridged Waveguide (QR–WG) which has four ribs on four sides
of a square or circular waveguide (Fig.1). In addition, a back short is provided at one end of the
QR–WG, and two coaxial core wires are embedded
at the resonance points as two feeding units. This
feeding part is an important part that determines
the reﬂection characteristics of QRA, and is hereafter referred as QR–Coax. The general electricalfeatures of QRA are as follows.
1. Very broad bandwidth over 100% of the fractional bandwidth based on the wideband characteristic of QR–WG and QR–Coax.
2. Separation and transmission of V/H orthogonally polarized waves by two fundamental transmission modes that occur in QR–WG.
3. The adjustability of the radiation beam size
and its pattern by optimizing the transmission
mode at the apperture, such like a waveguide hornantenna.
4. The simplicity of QR–Coax with only 2 coaxial outputs.

The problem is that such a wide bandwidth mustbe achieved simultaneously with high performancethroughout the band. Although the conventional
On the other hand, most of the QRAs on the
waveguide type feed has very high performance, itmarket are diﬃcult to match with the optical transhas the physical limit with a fractional bandwidth
mission system of the conventional radio telescopes

16

IVS NICT-TDC News No.39

due to its so large beam size, and also they have better in very wide band, hence the matching reso large input return loss of approximately −10 dB duces the return loss of the QR–Coax.
not adequate for a radio telescope feeder.
Thus, in this research, we designed a special
QRA which solves the problem described above.
The developing QRA will meet our requirements
as a wideband radio telescope feed antenna.

Figure 2. A schematic of the 6.5–22.5 GHz QRA
HFSS simulation model.
Fig.3 shows the HFSS analysis result of the input return loss at QR–Coax of this QRA model.
The Front–Port return loss which is near the apperture is approximately −19 dB in the entire bandof 6.5–23 GHz, and the Rear–Port closed to the
Back-Short is better than −20 dB in 6.5–22.5 GHz.
These are the greatly smaller return losses than
general QRAs and also enough to install into theradio observation system front–end.
Figure 1. A schematic and photograph of the open
type QRA.

3.

Designing of the 6.5–22.5 GHz QRA

Fig.2 shows a schematics of the developed 6.5–
22.5 GHz QRA using a 3D–FEM simulated CAD
model. In this QRA designing, we adopted awaveguide horn type QRA instead of the general
open type, in order to obtain the sharp beam sizeand perform mode matching to adjust radiation
beam pattern. The circular shaped aperture dimeter is 38mm, and the distance to the back–shortend (Length) is 138 mm. This is very small asa horn antenna that receives a radio wave at 6.5
GHz. The four tapered ridges, which are the characteristic structure of QRA, were shaped as an exponential curve which has less design parameters.
The two pairs of taper ridge curves corresponding to each of V and H–pol were independently
optimized by 3D–FEM electromagnetic simulator
Ansys HFSS. In addition, the back–short end isprovided with a hemispherical protrusion on thecentral portion of the circular waveguide as a characteristic structure of this QRA, which is called as”Plusphere back–short”. This hemisphere is considered to make matching the resonance between
the feeding point on QR–Coax and the back–short

Figure 3. Calculated input return loss of the 6.5–
22.5 GHz QRA.
Fig.4 shows another schematic of QRA, where
in which a conical horn antenna with a wider ﬂare
angle was directly connected to the apperture ofthis QRA, and Fig.5 shows the simulated radiation
beam pattern in the case that a signal was input tothe rear port of this model. A mode matching partwith an exponential curve was provided near the
QRA apperture, and it was optimized to matchthe beam pattern and its width designed by theconnected conical horn.
Accordance from Fig.5, it can be seen that theradiated beam pattern of this model has a relatively well–ordered and close to circular shaped.
The beam size, i.e. the FWHM of the beam, at
6.5 GHz was approximately 28 deg in both the Az

January 2021

17

and El direction, so that the pattern was close tocircular. On the other hand, though 20 deg in theAz at 22 GHz, 14 deg in the EL was relatively
narrow so that the pattern was eliptical. This ﬂattening of the beam pattern at the high frequency
band is considered to be inﬂuenced by the radiation pattern of the conical horn itself, and it can be
expected to be improved by optimizing the mode
matching in the future. Furthermore, the beam
size ratio of 3 : 1.5∼2 between 6.5 GHz and 22
GHz does not match the wavelength ratio of approximately 3.5 : 1. This is considered to be owing to the QR–WG property, which provides thewide bandwidth of QRA, and this feature can be astrong advantage in optimizing broadband opticaltransmission systems.

Figure 5. Calculated radiation main beam patternof the 6.5–22.5 GHz QRA with conical horn.
Figure 4. A schematic of the 6.5–22.5 GHz QRA
with wide–ﬂare conical WG horn.
4.

Summary

We develop a new type Quadruple-Ridged Antenna (QRA) as a very broad band feed antenna ofradio telescope. Unlike commercial products, thisspecial QRA has a very low input return loss of under −19 dB in a very wide band of 6.5–22.5 GHz,and its radiation beam pattern is circular and relatively sharp, also these are able to modify by optimizing mode matching conical horn part. Thesefeatures are very adequate as a radio telescope feed
antenna including the conventional (existing) telescopes. The current status of this development isongoing to manufacture QRA feed part called as
QR–Coax, and near future we will manufacture allpart of this 6.5–22.5 GHz QRA, also use it in practice to a VLBI observation.
References[1] J. Yin, J. Yang, M. Pantaleev and L. Helldner, ”The Circular Eleven Antenna: A New
Decade-Bandwidth Feed for Reﬂector Antennas With High Aperture Eﬃciency,” in IEEE
Transactions on Antennas and Propagation,
vol. 61, no. 8, pp. 3976-3984, Aug. 2013, doi:
10.1109/TAP.2013.2263275.

[2] G. Cortes-Medellin, ”Novel non planar ultra wide band quasi self-complementary antenna,” 2007 IEEE Antennas and Propagation Society International Symposium,
Honolulu, HI, 2007, pp. 5733-5736, doi:
10.1109/APS.2007.4396853.
[3] A. Akgiray, S. Weinreb, W. A. Imbriale and C. Beaudoin, ”Circular QuadrupleRidged Flared Horn Achieving Near-Constant
Beamwidth Over Multioctave Bandwidth:
Design and Measurements,” in IEEE Transactions on Antennas and Propagation, vol.
61, no. 3, pp. 1099-1108, March 2013, doi:
10.1109/TAP.2012.2229953.
[4] Y. Ma et al., ”A 10:1 Bandwidth Cryogenic
Quadruple-Ridged Flared Horn Design for Reﬂector Antennas in Radio Astronomy,” inIEEE Access, vol. 8, pp. 81101-81115, 2020,
doi: 10.1109/ACCESS.2020.2991350.
[5] A. Martellosio, M. Pasian, R. Rayet, S.
Rawson and T. Bonhoure, ”Wideband Cryogenic Receiver for Very Long Baseline Interferometry Applications,” in IEEE Antennas and Wireless Propagation Letters, vol.
17, no. 2, pp. 275-278, Feb. 2018, doi:
10.1109/LAWP.2017.2785958.

18

IVS NICT-TDC News No.39

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, Japan1.

FarewellThis is the ﬁnal report from Kashima about thedevelopment history for wideband antenna system.
So, I would like to look back the development
history of the wideband antennas brieﬂy and say
farewell to the Kashima VLBI group.
2.

SIRIUSAt ﬁrst, From 2009 to 2011, weak radiation measurement system was developed by the radio usage
fund. The system was named SIRIUS and aimed
to investigate the very weak noise from Ultra Wide
Band(UWB) communication equipments or PLC,
millimeter wave rader, or other communication devices. However they are weak, they can be serious
interference source for radio astronomy, geodetic
VLBI or some other scientiﬁc observation for weaksignal detection, if their population grows. And
also, their performance may be down by their own
mutual interference without suitable regulation for
their own noise levels. Thus our systems are aimed
for so weak level of noise as -90dBm/MHz withtime variation, out of the licensed bands, by detecting statistical abnormal radiation with two-element
interferometer. Detection technique is based on radio astronomy, which are switching radio meter andlong term averaging. The total bandwdith is from
0.8GHz to 26 GHz over, but that was divided into
3 bands because of limited bandwidth of receiverhorns and IQ mixers. The low band system is designed for 0.8 GHz - 3 GHz, mid if for 3 GHz 18 GHz, high-band is for 2GHz-26GHz. The received signal by the double-ridged horns were converted to IQ signals in video frequencies and sampled by VSSP32 VLBI sampler with USB 2.0 to
be processed in the personal computers. The image is rejected by the software before several hours
averaging. However narrow the bandwidth of 32
MHz which is maximal sampling rate for 8-bit in
VSSP32, this wideband receiver system was theetude for the Gala-V.

3.

Gala-V

Gala-V(Galapagos VLBI) project was started in2011, which was aimed to compare the atomic clock
frequencies over an inter-continental baseline with
2 small dish antennas. SNR of the total system isenhanced by the large dish of Kashima 34 m. Various wideband antennas were compared such as the
Eleven feed, Open Boundary Quad-ridged Horn,Taper slot antenna, and other conventional feedhorn before developing wideband antennas withnarrow beam width for conventional Cassegrain antennas such as Kashima 34 m. Finally, coaxial
multimode horn and axial corrugated multimodehorn were selected for the candidates of the wideband feed, because they are simple and can be arranged various beam width. The Eleven feed has
no room for changing its beam width, because thefacing elements should keep speciﬁc angles to eachother and the ground shield. Also, it needs lossy
balance-unbalance transition circuit and and many
LNAs. Open Boundary Quad-ridged Horn has difﬁculty for beam shaping. Coaxial multimode horn,
named IGUANA[1], is an extension of the multimode horns for VSOP-2/ASTRO-G VLBI satellite
and 6.7GHz methanol maser horns for VERA 20
m and Shanghai 25 m antenna. Axial corrugatedmultimode horn, named NINJA[2], is compact andhas a potential for arranging the beam width. They
are quite simple and easy to be manufactured by
NICT. Small portable VLBI station MARBLE1
and 2 were 1.6 m and 1.5 m parabola antennas
with a Open Boundary Quad-ridged Horn. They
were modiﬁed into 2.4 m dish and NINJA feed toobtain 3 times better sensitivity. MARBLE1 wastransported to Medicina radio observatory in thesummer of 2018, and the main reﬂector of 34 m
was well adjusted by radio holography with broad
cast satellite to start the experiment. Finally theexperiment shows 10−16 accuracy of comparison[3].
4.

Next generation microwave radiometers(no name yet)

Development of a novel water vapor radiometerhas been started in 2018. The will realize rapid andaccurate measurement of vapor, water and oxygenin the atmosphere 15 - 60 GHz by one receiver system with a wideband feed. OMT is used to divide the signal to lower LNA(15 - 33 GHz) andhigher LNA(26 - 60 GHz) because full band LNA
was not found. Calibration scheme by RF switch
before LNA, like SIRIUS, is not practical, because
it will be increase input noise by several dB. Thus apair of cold and hot absorbers well measured theirtemperatures will be set and used for reliable cali-

January 2021
bration. The 22 GHz multimode horn for VSOP-2
Bread Board Model was set on NICT Okinawa 3.7
m antenna and used in 16 - 33 GHz for an experiment to measure the water and vapor absorption
in atmosphere in 2019. Wideband OMTs and feeds
are developed in 2020 for 3.7 m antenna and otheroptics1, but experiments were postponed becauseof COVID-19.
5.

BRAND

BRAND is the project to expand the bandwidthof European radio telescope network. Successors of
IGUANA and NINJA feeds have been in evolution
to 1.5 - 15.5 GHz bandwidth for secondary focus
solutions2.

19

Acknowledgement. Next generation radiometerhas been supported by JSPS KAKENHI Grant
Number JP18H03828. Development of wideband
feeds was supported by NAOJ Joint Development
Research Grant in FY2009-FY2014 and NICT incentive fund in FY2013.
References[1] Ujihara, H., (2016) Development of wideband
feed for Kashima 34m Antenna, Radio Sci., 52,
doi:10.1002/2016RS006071[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,
6. Lessons Learned
Japan, July 30 - August 4, 2017 (peer-reviewed),“Keep it simple, stupid” is a familiar phrase(10.1007/1345 2018 41).
for software developers, especially Linux users.
Kashima 34 m has complex receivers systems on [3] Pizzocaro, M. et al., (4th:Ujihara H.), Intercontinental comparison of optical atomic clocks
4 facing trollies, which push up maintenance timethrough very long baseline interferometry, Naand cost but never guarantees position accuracy
ture. Physics. , doi:10.1038/s41567-020-01038-6
under mm. Also some feed can not be placed on(2020).
the center of focus and reﬂectors were lost the origin of alignments. However radio holography wasdone with CW from a broad cast satellite, it isstill unsure alignments on the reﬂectors and feed.
Thus it is hard to guarantee the best surface by theholography over the whole band of 3 - 15 GHz, justsay it is enough to the Gala-V experiments. Simply, if the symmetry of the main beam in E-cut andH-cut have been measured, alignment origin might
be recovered and more eﬃciency was obtained. Instead of trolly system, revolving systems seemed tobe reliable, or more simply tilting sub-reﬂector to
select a horn seems to be better to make maintenance easier and better position accuracy. Also, if
Figure 1. Wideband radiometer feed system.
some feeds were combined and replaced by wideband feed, including the huge L-band horn, 34 m
could have more research areas. Metabolism is important to extend the life time and applications ofhuge equipment like 34 m antenna, but it was notconcerned to the end of 34 m.
Gala-V was named after the unique evolutions
in Galapagos islands, it has several meanings andirony based on biology. What is need to survive is the variation of the food, which is the research themes and applications for the researchers.
Kashima 34 m has been stick to VLBI and not so
active to other studies and destroyed by a typhoon
blast with slightly less speed than designated windlimit speciﬁcation. I am still wondering Kashima Figure 2. Simulated beam patterns of a wideband
feed for BRAND.
34 m was designed appropriately.

20

IVS NICT-TDC News No.39

Performance survey of super- suppresses only the RFI in S-band. This report
the results of survey and test observations
conductor filter introduced in presents
to conﬁrm the performance of the ﬁlters.
wideband receiver for VGOS of
2. Characteristic survey of superconductor
the Ishioka VLBI station
filter introduced in the Ishioka VLBI
station1

Tomokazu Nakakuki(nakakuki-t96nm@mlit.go.jp),
Haruka Ueshiba1 , Saho Matsumoto1 ,Yu Takagi1 , Kyonosuke Hayashi1 ,Toru Yutsudo1 , Katsuhiro Mori1 ,Tomokazu Kobayashi1 , Mamoru Sekido2

We introduced a superconductor ﬁlter to attenuate only the RFI in S-band. Because a superconductor ﬁlter has steep cut-oﬀ characteristics. In
addition, cryogenic dewar of current receiver system was suitable for installation of the superconductor ﬁlter. Figure 1 shows the reﬂective andtransmission characteristics of the ﬁlter from 1 to
1 Geospatial Information Authority of Japan,
15 GHz for both polarizations (H-pol: Horizon1 Kitasato, Tsukuba Ibaraki 305-0811, Japantal linear, and V-pol: Vertical linear) of the in2 NICT Space-Time Standards Laboratory,
troduced ﬁlter. It is designed to attenuate speciﬁc
Kashima Space Technology Center,
frequencies around 2 to 3 GHz. Details from 2 to
893-1 Hirai, Kashima Ibaraki 314-8501, Japan 3 GHz are shown in Figure 2. It can be seen thatwhile tiny transmission loss of 2.2-2.4 GHz used
1. Introduction
for legacy S/X session, upper and lower band to be
The Ishioka VLBI Station operated by the suppressed are attenuated with sharp cutoﬀ charGeospatial Information Authority of Japan (GSI) acteristic.
participates in legacy S/X sessions (2 GHz and 8
GHz band) conducted by the International VLBI
Service for Geodesy and Astrometry (IVS). The
Ishioka station also participates in VGOS (3-14GHz) sessions, which is VLBI observation systempromoted by IVS. Because the Ishioka station employs diﬀerent receiving systems for legacy S/X sessions and VGOS broadband sessions, it is necessary to exchange receivers in order to participate
in each observation. Since it takes about a week toreplace and adjust the receiver, it is practically difﬁcult to replace frequently to avoid interruption ofobservations. Therefore, the Ishioka station is currently regularly implementing legacy S/X sessions
to maintain geodetic reference system of Japan,
and is participating in VGOS sessions for several
months in a year. GSI has a plan to participate
in legacy S/X sessions with broadband receiver of
VGOS speciﬁcation in order to participate in bothsessions constantly and more ﬂexibly. Although Figure 1. Green and orange curves represent thethe broadband receiver is capable to receive the 2- transmission and reﬂection characteristics of the
14 GHz frequency, the signal of S-band is seriously superconductor ﬁlter, respectively. The upper panel
contaminated by strong radio frequency interfer- is for the H-pol and lower panel for V-pol.
ence (RFI) caused by artiﬁcial radio waves aroundthe Ishioka Station. To avoid saturation of the receiver stem by the RFI, we are currently attenuat- 3. Detailed investigation using VLBI antenna of Ishioka
ing signal below 3 GHz with applying a high-passﬁlter after Low-Noise-Ampliﬁer (LNA), but there
We investigated the performance of RFI attenustill remains a problem that the signals around 2 ation after installation of superconductor ﬁlter in
GHz are not available. Thus, substantial counter- the cryogenic receiver system. Similary to the lastmeasures against the RFI have been required to year’s RFI survey (Nakakuki et al., 2019), the oututilize the full-band data; 2-14 GHz. To solve the put of LNA of broadband receiver was measured
problem of the RFI, we introduced a ﬁlter that using spectrum analyzer of Agilent Technologies,

January 2021

21

Figure 4. Averaged (100 scans) power spectrumof received signal for 2-14 GHz frequency range at
Az=90 deg, El=5 deg.
time duration, and that does not appear in time
averaged spectrum. (Figure 4)
Figure 2. The same as Figure 1, but the range of
frequency is 2-3 GHz.

Figure 3. Receiving intensity from 2 to 14 GHz
for the azimuth of 90 degrees and the elevation of
5 degrees. Blue and Orange curves represent theintensity of received signals in the surveys wtih (in2020) and without (in 2019) superconductor ﬁlter,
respectively.

4.

Confirmation of fringe

Since the attenuation of the RFI was conﬁrmed,
we conducted VLBI experiment to check fringe detection. The test observations were carried out
with the 11 meter antenna of the Koganei station operated by National Institute of Information
and Communications Technology (NICT). Table 1
shows the conditions for the test observations. Correlation processing was carried out by NICT and
GSI using K5-correlation system respectively.
Table 1. The setting of the test observations.
Obs. dateJun. 29, 2020 5:00-6:00 (UTC)Obs. band
S-band 6 Chs. (16MHz/1ch)
X-band 10 Chs. (16MHz/1ch)
Obs.Celestial 29 sources (Obs. time of 60 tobody
180 s/sources)

Inc. We observed the spectrum while changing theelevation angle with the azimuth angle of 90 degrees (The azimuthal angle is measured in clockwise from the North. The elevation is from horizon). Figure 3 shows the received power for the frequency range from 2 to 14 GHz with an elevation
angle of 5 degrees, where RFI was the strongest as
in the RFI survey before the introduction of superconductor ﬁlter. As shown in this ﬁgure, both radio waves, emitted from mobile phone base station(2.1 GHz) with almost constant power and intermittently emitted radar radio waves (2.8 GHz), areattenuated. The radar power changes from time
to time and it cannot always be fully attenuated. Figure 5. Correlation result of the Ishioka station
However, we suppose that is acceptable due to low (H-pol) and the Koganei station(R-pol).

22

IVS NICT-TDC News No.39
The investigation of the received signal revealed
that the signal which caused the RFI was attenuated. In addition, conducting the test observation with the Koganei station, we conﬁrmed thatfringes were successfully detected with received signal which passed through the superconductor ﬁlter. From these results, we can say that one of the
major obstacles for the observations of legacy S/X
sessions with broadband receiver was resolved. For
the next step to realize participation in legacy S/X
sessions with broadband receiver, polarization conversion is necessary from linear to circular. We willcontinue eﬀort to solve this issue.
References

Figure 6. Correlation result of the Ishioka station(V-pol) and the Koganei station(R-pol).
5.

Summary

The Ishioka station introduced superconductor
ﬁlters on the broadband receiver which successfully
attenuate the RFI in S-band in order to participate
in legacy observations with the broadband receiver.

[1] Nakakuki T., K. Hayashi, N. Ishikura, M.
Umei, S. Matsumoto, T. Kikkawa, T. Yutsudo, H. Munekane, M. Sekido, Radio Frequency Interference research a round Ishioka VLBI Station using VGOS broadband
receiver, Japan VLBI Consortium, pp.6669,2019.

January 2021

23

New calibration method
depending strongly to a loss of the liquid nitrogenand slightly to barometric pressure.
for a radiometer without using container
The receiver noise temperature is determined byliquid nitrogen cooled absorber equation (2), but requires an absorber cooled in liqNoriyuki Kawaguchi (kawagu.nori@nao.ac.jp),
National Astronomical Observatory,
2-21-1 Osawa, Mitaka, Tokyo, Japan
Yuichi Chikahiro, Kenichi Harada andKensuke Ozeki
Elecs Industry Co. Ltd.,
1-22-23 Shinsaku, Takatsu, Kawasaki, Japanuid nitrogen. Care must be taken when handlingliquid nitrogen, and it is diﬃcult to perform calibration work frequently. In addition, calibration
at various elevation angles cannot be performed.
Figure 1 shows how liquid nitrogen is ﬁlled in thecontainer containing the absorber.

Abstract: In radio astronomy receivers and microwave radiometers, it is important to know thereceiver noise temperature accurately in order toknow the radiance temperature of the measurementtarget. Liquid nitrogen-cooled absorbers are often
used to calibrate the receiver noise temperature.
The liquid nitrogen-cooled absorber is a standard
noise source of the most reliable, but cannot be
used frequently during measurements or for calibration at various elevation angles.Therefore, we Figure 1. Liquid Niterogen ﬁlling into the absorber
have developed a simple method for calibrating the container at Tianma 65m site in Shanghai Astroreceiver noise temperature that does not require nomical Observatory
liquid nitrogen, and here the reliability experimenIn recent years, the use cases of array antennas
tally conﬁrmed is reported.
are increasing in the ﬁeld of radio astronomy. Especially in SKA-related ﬁelds, we are promoting the
1. Introduction
introduction of ‘Aperture Array’ technology withIn a microwave radiometer, noise temperature
the aim of improving survey speed and wide-ﬁeld
of a target source is measured by so-called as aobservation.However, it is very diﬃcult to measure
Y factor method. The noise temperature of thethe noise temperature of the Aperture Array retarget, Ttg can be obtained from the received power
ceiver with a built-in low noise ampliﬁer (See Figratio in the equation (1).
ure 2).
Tamb + Tr
Y1 =
,
(1)
Ttg + Tr
where Tamb is ambient temperature of a non-cooled
absorber or a non-reﬂective terminator, Tr is receiver noise temperature. The noise power ampliﬁed by the receiver is expressed as P = kT BG,
in which k is the Boltzmann constant, B is bandwidth, G is the receiver gain, and T is noise temperature. By taking the received power ratio, thecommon terms, k, B, and G are removed, and only
the noise temperature terms are left as shown inEquation 1. Tamb is easy to measure but Tr is not.
In order to know the receiver noise temperature, Figure 2. Liquid Niterogen Calibration for a large
Tr , it is necessary to measure the Y factor by the array anntena [1]next equation, Equation 2 in advance.
Y2 =Tamb + Tr,
Tc + Tr(2)

where Tc is noise temperature of a liquid nitrogencooled absorber, exactly 77.364K and around 80K

We devised a new calibration method that canmeasure the noise temperature of the receiver andtarget by a simple method without using liquid nitrogen (patent pending). In the next section theory

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IVS NICT-TDC News No.39

of calibration is presented and the Section 3 experimental results are shown.
2.

Theory of calibration

We introduce a microwave switch in front of aﬁrst stage ampliﬁer, a Low Noise Ampliﬁer, LNA,
which switches between an antenna, ‘Horn Antenna’, a non-reﬂective terminator, ’Hot Terminator’ , and a fully reﬂective terminator, ‘Short Terminator’ as indicated in Figure 3.

measured in advance.
The parameter k is measured by placing an absorber cooled with liquid nitrogen on the horn antenna, the noise temperature of which is taken as areference of known temperature, around 80K. The
ﬁrst of all, Tr is measured from a power ratio of
Y4 expressed by Equation 5 by switching ‘Hot Terminator’ to ‘Antenna’ fully covered with a cooled
absorber. The power ratio can be written as
Y4 =Tterm + Tr.
Tc + Tr(5)

where Tterm is terminator temperature measured
with a thermometer attached on the terminator
and Tc is liquid nitrogen temperature. Both Tterm
and Tc are known so that Tr is derived from Equation 5.
In the next. switching between ‘Hot’ and ‘Short’Figure 3. A schematic diagram of a new calibration and the power ratio, Y5 expressed by Equation 6 issystemmeasured. By introducing the relation, Equation 4
and the Tr measured, the k is derived..
When the switch is connected to the ‘Antenna’,
Tterm + Tr
Tterm + Tr
a sky noise through the horn antenna is received,
Y5 =
=
,
(6)
when connected to the ‘Hot Terminator’, ambient
Tshort + Tr
(k + 1)Tr
temperature noise is received, when connected toOnce k is ﬁxed, then we can know the receiver‘Short Terminator’, a part of LNA noise is reﬂected
noise temperature as
and received.
Now we deﬁne a new Y-factor, a power ratio ofthe ‘Hot’ to ‘Short’. expressed by an equation beTterm
,
(7)
Tr =
low.
Y5 (k + 1) − 1
Tamb + Tr
Y3 =,(3) without using an absorber cooled with liquid niTshort + Tr
trogen, only switching ‘Hot’ and ‘Short’, and wewhere Tamb is ambient temperature to be measured can measure the target noise temperature, Ttg by
by a thermometer attached on the terminator and a power ratio of Y1 in equation (1).
Tshort is a reﬂected LNA noise. The reﬂected noise
temperature Tshort is a part of the receiver noise. 3. Experiment Results.
If we assume the factor of of the reﬂection is k, we
3.1 Receiver noise by using k
can know the Tshort as,Spectra of receiver noise temperature measured
Tshort = k · Tr .
(4) with a conventional way (Hot/Cold) using a liquid
nitrogen cooled absorber are compared with thosemeasured with a new method, switching between‘Hot’ and ‘Short’ using a k-factor in Figure 5.
Both are in good agreement and the correlationis shown in the Figure 6.
Figure 4. A simpliﬁed noise equivalent circuit of a
LNA
A simpliﬁed noise equivalent circuit is shown in2
Figure 4. The noise generated at the ‘Gate’, vRgappeared not only on the load resistance, RL as areceiver noise, but also on the input resistance, Rs .
The coeﬃcient, k depends on the internal parameters of the semiconductor device and needs to be

3.2

Reliability of a new calibration method

In order to test the reliability, the horn antenna
is covered with an absorber cooled with liquid nitrogen, the noise temperature of which is supposed
to be around 80K. Figure7 shows the test results.
As a result of 10 trial tests, the mean value was
80.26K and the standard deviation was 0.25K.
In the next time, the horn antenna is covered
with a non-cooled absorber, the noise temperature

January 2021

25

Figure 5. Receiver noise spectra measured by‘Hot’/’Cold’ and ‘Hot’/’Short’ using ‘k’

Figure 7. Cooled noise temperature of 10-times
trial at a day

Figure 6. Receiver noise temperature measured by‘Hot’/’Cold’ using liquid nitrogen cooled absorberand ‘Hot’/’Short’ using ‘k’

Figure 8. Terminator temperature of 10-times trials at a dayof which is supposed to be room temperature, 302K
measured by a thermometer. Figure 8 shows thetest results.As a result of 10 tests at room temperature of 28,9◦C (302K), the average value of thenoise temperature was 303.4K and the standard
deviation was 0.8K. From the above test measurements for absorbers of noise temperature known,
the certainty of the new calibration method was
conﬁrmed.

it also can be seen that the value of k tends to decrease as the temperature rises by about 0.003/deg.

In Figure 11, noise temperature of liquid nitrogen cooled absorber is shown. Blue dots indicate
noise temperature when we take the k-factor calibrated on each day, the mean over the 10 days andthe standard deviation are 80.30 and 0.21. Red
dots indicate noise temperature when the k-factor
is ﬁxed to that calibrated in the ﬁrst day, the mean
and the standard deviation are 80.39 and 0.82.
3.3 Long-term stability
In Figure 12, noise temperature of la non-cooled
In a previous section, results of 10-times trials on absorber is shown. Blue dots indicate noise tema day are reviewed. Here we present a 10-days ex- perature when we take the k-factor calibrated on
periment repeated same switching operation, ‘Hot’, each day, the mean over the 10 days and the stan‘Short’, ‘Antenna’. At the ‘Antenna’ position we dard deviation are 301.31 and 1.56. Red dots incovered the horn antenna with a cooled absorber dicate noise temperature when the k-factor is ﬁxed
and a non-cooled one, respectively.
to that calibrated in the ﬁrst day, the mean and
Figure 9 shows values of k obtained on each day. the standard deviation are 302,55 and 1.44. Blackthe averaged value is 0.446 with the standard de- dots are room temperature.
viation of 0.0054. At the 6th day, it was clearly
Figure 13 is receiver noise spectrum measured 10
smaller than the other days. We found at the the days after using the k-factor obtained in the ﬁrst
6-th day, the room temperature and the termina- day. As compared to Figure 5, the receiver noise
tor temperature were extraordinary high by 3 to 4 spectrum measured in the ﬁrst day and the k-factor
degree as is shown in Figure 10. Only the day the on the day, Figure 13 shows small errors at some
air conditioning system was not working. Similarly,

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IVS NICT-TDC News No.39

Figure 9. Stability of a parameter k over 10 days

Figure 11. Stability of noise temperature of a cooled
absorber over 10 days

Figure 10. Room temperature variation over 10
days
Figure 12. Stability of noise temperature of a noncooled absorber over 10 days
frequencies, but not so large, a few kelvin in the
average. We concluded the k-factor and the noise
calibration without using a liquid nitrogen cooled
absorber work eﬀectively for at least 10 days, The
longer test for one year is now going on at the
VERA VLBI station in Iriki.
4.

Summary

We successfully performed the noise calibration
without using liquid nitrogen. The new calibration system is quite useful and will be introduced
to a water vapor radio meter mounted on a radio
telescope. This system is also eﬀective for noise
temperature calibration of array antennas. In particular, this method is indispensable for sensitivity
calibration of individual array elements. Optimal
beam forming will be performed in consideration ofvariations in the sensitivity of each array element.
It is also possible to make calibration of a large
array antenna without using a large container for
liquid nitrogen,5.

Acknowledgments

We thanks to Dr. Akiji Nakagawa and his staﬀ inKagoshima University for their helps on introducing the water vapor radiometer on Iriki 20-m radio

Figure 13. Receiver noise spectrum by using a kfactor calibrated in the past 10 days
telescope. We would like to thank Dr. Takashi
Maeda for his valuable advice regarding array antennas.
References[1] A. P. Chippendale, et al., Measured ApertureArray Noise Temperature of the Mark II
Phased Array Feed for ASKAP, ISAP 2015,
IEICE

January 2021

27

Comparison of atmospheric delay models (NMF, VMF1, andVMF3) in VLBI analysis
Mamoru Sekido1 (sekido@nict.go.jp ) andMonia Negusini2
1 NICT Space-Time Standards Laboratory,
Kashima Space Technology Center,893-1 Hirai, Kashima Ibaraki, Japan
2 Istituto Nazionale di Astroﬁsica/
Istituto di Radioastronomia,
Via Gobetti, 101, 40129 Bologna, Italy
Abstract: In case of VLBI observation of a longsingle baseline, observing antenna directions distributes in a small part of sky. Hence, accurate atmospheric delay parameter, especially anisotropic
component, is diﬃcult to estimate from VLBI data
itself. Therefore, atmospheric delay model for apriori delay calibration is essentially important toachieve better accuracy in VLBI delay analysis of along single baseline. We compared atmospheric delay models of NMF, VMF1, and VMF3 with broadband VLBI observation data of the Medicina - Koganei baseline. As a result of the comparison, we
found that VMF3 provides the most accurate a priori calibration among them.
1.

Introduction

Figure 1. Medicina (Italy) and Koganei (Japan)
baseline (8,785 km) was used for precise frequency
link between Yb and Sr optical lattice clocks operated at INRIM and NICT, respectively.
cle describes comparison of these atmospheric delay
models applied to the Medicina – Koganei baseline.
2.

Data and analysis

The VLBI data used for the comparison is thesame dataset used for the frequency link[1].
Table 1 shows the list of VLBI sessions conducted between Medicina (MARBLE1) and Koganei (MARBLE2) antennas. The Kashima 34m antenna had been participating all the sessions
to form virtual delay observable via NHS scheme[2, 3], although this is not the subject here.
Fig. 2 shows the occurrence histograms of azimuthal and elevation angles at each station during these observation sessions. Corresponding histograms of standard geodetic VLBI for the period 2018-2020 are plotted for reference. Atmospheric delay contributions are parameterized in
VLBI analysis by its elevation and azimuthal angles dependencies. Indeed, standard geodetic VLBI
sessions are scheduled with care for sky coverage ateach station. That is possible by network observation with globally distributed stations. Instead, incase of a single long baseline, observing directions
are limited in a small part of sky due to mutual visibility. Atmospheric delay parameters, especiallyazimuth angle dependent anisotropic components,
are diﬃcult to be estimated by observation data
itself as it is performed in standard geodetic VLBI
analysis. Thus, we need to calibrate the atmospheric delay by a priori data as accurate as possible before least square analysis.
We examined NMF, VMF1, and VMF3 atmospheric delay models in data analysis with batchSolve [7]. The NMF model is available in the Solveas bult-in model, and the atmospheric delay prediction with VMF1 and VMF3 were externally computed and stored in MK3 database as ’cable delay’data1 . Atmospheric delay is separately modeled by
the contribution of hydrostatic component referred

We have conducted international experiments toapply VLBI for precise frequency link under thecollaboration among National Institute of Information and Communications technology (NICT), Istituto Nazionale di Ricerca Metrologica (INRIM), Istituto Nazionale di Astroﬁsica/Istituto di Radioastronomia (INAF/IRA), and Bureau Internationaldes Poids et Mesures (BIPM). By using novel techniques in observation, signal processing, and data
analysis, we successfully demonstrated the precise frequency link over intercontinental distance[1]. Moreover, geodetic performance of our system with transportable VLBI stations turned out
to be comparable to standard VLBI observations
made within IVS [2]. One of the key elementsin the success of the experiments was calibrationof atmospheric delay. Because of a long single
baseline (Fig. 1), sky coverages at each antenna
were limited. Hence, atmospheric delay parameters, especially anisotropic contribution, cannot
be estimated accurately. In this project, we haveexamined atmospheric delay models provided by
Niell Mapping Function (NMF) [4], Vienna Map1 It was because ’cable delay’ data was convenient to inping Function (VMF1) [5], and its extension Vienna Mapping Function 3 (VMF3) [6]. This arti- clude external calibration data in Solve analysis.

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IVS NICT-TDC News No.39

Table 1. VLBI sessions between two 2.4 m diameter antenna located at Medicina and Koganei. WRMS
represents weighted root-mean-square of post-ﬁt residuals in the baseline analysis for the NHS virtual
delay.
Date

MJD

5 Oct. 201814 Oct. 2018
4 Nov. 2018
14 Nov. 2018
24 Nov. 2018
4 Dec. 201815 Dec. 201825 Dec. 201815 Jan. 201925 Jan. 20194 Feb. 201914 Feb. 2019

58396
58406
58426
58436
58447
58457
58467
58477
58498
58508
58518
58528

Session
[h]
31.4
28.9
30.6
29.0
28.8
29.0
29.5
28.9
29.0
30.6
31.0
35.8

#Scans(Used/Total)
1366 / 1470
1155 / 1415
1452 / 1645
1419 / 1539
1291 / 1435
1344 / 1511
1379 / 1470
1439 / 1501
1363 / 1437
1336 / 1591
1342 / 1500
1341 / 1585

#Scans
/hour
46.9
49.0
53.7
53.1
49.8
52.2
49.9
51.8
49.6
56.4
48.4
44.3

WRMS
[ps]
30
32
39
24
29
33
26
22
24
26
30
29

Figure 2. Occurrence of azimuthal (left) and elevation (right) angles of observation directions from eachstation in VLBI observation of Medicina - Koganei baseline. All the scans for twelve sessions in Table
1 are used. As a reference of standard geodetic VLBI observation, corresponding occurrence histogram
data are plotted for the same stations based on the schedule ﬁles of 59 IVS sessions of R4, RDV, T2,
CRF, and EURO for Medicina (•) and 91 sessions of T2, CRF, AOV, and APGS for Koganei () in
the period of 2018–2020.
as ’dry’ and others dominated by water vapor referred as ’wet’. Each of them is parameterized byzenith delay and a factor called mapping function(mf ()), which is a ratio of slant delay in the line ofsight and that in zenith direction. In addition, 2-D
gradient vector (Gn , Ge ) are used in modern VLBI
data analysis to represent anisotropic contribution.
The Atmospheric delay model is expressed asτatm = τhz · mf h () + τwz · mf w ()+ mf g () · [Gn · cos(α) + Ge · sin(α)],(8)

where  and α represent elevation and azimuthal
angles.
Both VMF1 and VMF3 data are based on

ray tracing simulation with numerical weather
data of the European Centre for Medium-Range
Weather Forecasts (ECMWF). An advanced feature of VMF3 model is that provides not only mapping functions, but also gradient coeﬃcients computed by ray tracing, and they can be used for apriori delay calibration. Table 2 shows the atmospheric delay models used for a priori delay calibration. The dry components of zenith delay werecomputed with ground pressure data using Saastamoinen model [9] for all the cases. The wet zenith
delay parameters of both VMF1 and VMF3 andgradient coeﬃcients of VMF3 are given by 6 hours
interval, and the data are provided by the VMF
data server [8]. Estimated parameters in the least

January 2021

29

Table 2. Atmospheric delay calibration performance is compared for NMF, VMF1, and VMF3 via
estimated atmospheric zenith delay parameters. Average and standard deviation computed over all thesessions listed in Table 1 are presented.
Model

Delay components

NMF
VMF1
VMF3

Dry
Dry+WetDry+Wet+Gradient

700
600
500
400
300
200

VMF1

MARBLE1

VMF3

MARBLE2

100
0
-100
-200
-300
14OCT.2018:MBL1
04NOV.2018:MBL1
14NOV.2018:MBL1
24NOV.2018:MBL1
04DEC.2018:MBL1
15DEC.2018:MBL1
25DEC.2018:MBL1
15JAN.2019:MBL1
25JAN.2019:MBL1
04FEB.2019:MBL1
14FEB.2019:MBL1
14OCT.2018:MBL2
04NOV.2018:MBL2
14NOV.2018:MBL2
24NOV.2018:MBL2
04DEC.2018:MBL2
15DEC.2018:MBL2
25DEC.2018:MBL2
15JAN.2019:MBL2
25JAN.2019:MBL2
04FEB.2019:MBL2
14FEB.2019:MBL2
MBL1-All-data
MBL2-All-data

Average of atmospheric zenith delay [ps]NMFz
Avg. δτatm± σatm [ps]MARBLE1 MARBLE2
231 ± 157
237 ± 183
-40 ± 67
0 ± 64-11 ± 45
7 ± 60Figure 3. Estimated atmospheric zenith delay residual averaged over each session per station. Threecases of atmospheric delay calibration with NMF, VMF1 and VMF3 are compared.
square analysis are baseline vector, clock parameter (rate and oﬀset), and zenith atmospheric delay
at one-hour interval for each station.
3.

Results and conclusion

Fig. 3 shows estimated atmospheric zenith delay parameters averaged over each session per station. While delay calibration using NMF results in
atmospheric delay residual of hundreds of picoseconds, the remaining atmospheric delays when using VMF1 or VMF3 are of the order of tens of ps.
This indicates that numerical weather model-based
ray tracing technique provides very realistic atmospheric delay prediction. The smaller magnitude
and lower scatter of estimated atmospheric delay(Table 2) of VMF3 compared to VMF1 can be at-tribute to the contribution of the gradient component of the VMF3, which represents the anisotropic
property of atmosphere. This advantage of VMF3
is essentially important for the single long baseline,
where estimating the atmospheric delay parameteris diﬃcult from VLBI data itself.
Acknowledgements
We thank to all the members of collaboration team from INRIM (M. Pizzocaro, C. Clivati,
D. Calonico, et al.), INAF (F. Perini, G. Maccaferri, R. Ricci, et al.), BIPM (G. Petit and J.Leute),
and NICT (K. Takefuji, H. Ujihara, T. Kondo,
N. Nemitz, T. Ido, et al.) for making possible
this international VLBI project for precise frequency transfer. We thank J. Böhm and D. Land-30

skron of the Vienna University of Technology (TU
Wien) for providing the VMF1 and VMF3 data.
We thank VLBI group of the NASA/GSFC/NVI(E. Himwich, J. Gipson, D. Gordon, et al.) for
the FS9, Calc/Solve software resources that were
instrumental in this project.
References[1] Marco Pizzocaro, et al., “Intercontinental
comparison of optical atomic clocks via very
long baseline interferometry”, Nat. Phys.,
2020, doi: 10.1038/s41567-020-01038-6.
[2] Mamoru Sekido, et al., “A broadband
VLBI system using transportable stations for
geodesy and metrology – An alternative approach to the VGOS concept–”, J. Geod., inprinting, 2021.
[3] Mamoru Sekido, “’Node - HUB’ StyleVLBI with Broadband System”, IVS NICTTDC News. 37, pp. 22-25, 2017, URL
https://www2.nict.go.jp/sts/stmg/
ivstdc/news_37/pctdc_news37.pdf[4] Arthur E. Niell, “Global mapping functions
for the atmosphere delay at radio wavelengths”, J. Geophys. Res. , 101, pp. 32273246, 1996.

IVS NICT-TDC News No.39[5] Johannes Böhm, Birgit Werl, and Harald
Schuh, “Troposphere mapping functions for
GPS and VLBI from ECMWF operational
analysis data”, J. Geophys. Res. 111, B02406,
2006, doi:10.1029/2005JB003629.
[6] Daniel Landskron & Johannes Böhm,“VMF3/GPT3: reﬁned discrete and empirical
troposphere mapping functions”, J. Geod.,
92:349-360, 2018, doi: 10.1007/s00190-0171066-2.
[7] David Gordon, et al., “GSFC VLBI analysis center annual report”, In: Armstrong
KL, Baver KD, Behrend D (eds) IVS
2015+2016 Biennial Report, pp. 224-228,
2017, URL https://ivscc.gsfc.nasa.gov/
publications/br2015+2016/acgsfc.pdf.
[8] VMF Data Server; editing status 2020-1214; re3data.org - Registry of Research Data
Repositories, 2020. DOI: 10.17616/R3RD2H[9] J. Saastamoinen, “Atmospheric correction fortroposphere and stratosphere in radio ranging of satellites”, in The Use of Artiﬁcial
Satellites for Geodesy, 15, 247, 1972 doi:
10.1029/GM015p0247.

January 2021

HINOTORI Status Report
Hiroshi Imai (hiroimai@km.kagoshima-u.ac.jp),
and HINOTORI Team
Amanogawa Galaxy Astronomy Research Center,
Graduate School of Science and Engineering,Kagoshima University, 1-21-35 Korimoto,Kagoshima 890-0065, Japan
Abstract: Here we introduce our Hybrid Integration Project in Nobeyama, Triple-band Oriented(HINOTORI), which is composed of the development of a quasi-optics system capable of simultaneous signal reception in the triple bands (22,43, 86 GHz) and the upgrade of the VLBI backend system in the Nobeyama 45 m telescope. The
triple-band receiving system has been in commission since 2020 September with the expected performance. The new VLBI backend system is still
being tested so as to achieve a recording rate of
32 Gbps. This report describes the speciﬁcations
and the present statuses of these major HINOTORI components.
1.

Introduction

Single-dish and VLBI observations that enable
us to simultaneously receive and record radio signals in multiple frequency bands have demonstrated their high scientiﬁc potentials in terms of
not only a higher observation eﬃciency but alsoreliable registration of radio emission maps of themultiple bands onto a common coordinate system.
One of the scientiﬁc demonstrations of the simultaneous multiple band VLBI observations is presented by Yoon et al. [5], who observed circumstellar H2 O and four transitions of SiO masers infour (22, 43, 86, and 129 GHz) bands. They revealed the whole morphology of the SiO maser ringvisually clariﬁed by multiple transition maser lines
and biased H2 O maser distributions with respect
to the mass-losing star located at the ring center.
This will shed light on the property of inhomogenous mass loss from a red supergiant.
Millimeter (especially 86 GHz) VLBI observations have made use of great beneﬁt of such a simultaneous multiple frequency band receiving technique, which enables us to apply the band-to-band
phase transfer technique and the source-frequency
phase-referencing technique [1]. They shall provide
solutions of fringe-phase calibration from the data
in a lower frequency band (e.g., 22 GHz) to the datain higher frequency bands (e.g., 43, 86 GHz) for realizing longer coherent integration of the higher frequency band data. The special system that enables
us to simultaneously observe in multiple frequency

31

bands have been in commission at ﬁrst in the Korean VLBI Network (KVN), then 40 m telescopein Yebes, Spain. Here we introduce such a systemdeveloped for the 45 m telescope of Nobeyama Radio Observatory (NRO) by HINOTORI, which has
been promoted since 2016.
2.

HINOTORI specification

Figure 1 shows the schematic diagram of thewhole HINOTORI system.
The quasi-optics of the Nobeyama 45 m telescope
enables us to select one of the receivers by inserting mirrors to and releasing them from the optics.
Instead of these mirrors, the dichroic ﬁlter plate
we have developed shall guide the transferred radio signals into two passes in the optics. By inserting two such ﬁlter plates, we can simultaneously receive signals in three frequency bands. The
gain losses due to the ﬁlters have been measured tobe ∼0.22 dB, causing slight increase in the systemtemperatures by ∼12 K in the 43 GHz band [3].
This gain loss is higher (1.0–1.5 dB, dependent onradio frequency) in the 86 GHz band [4]; it will be
mitigated by revising the design of the ﬁlter in the
future.
In order to ﬂexibly select the multiple frequency
band signals, HINOTORI shall utilize the intermediate frequency (IF) switch box operated by NRO.
HINOTORI also has installed an IF selector box
in which one to six out of eight IF passes are connected to the inputs of the two digital samplers
OCTAD (Oyama et al., in preparation). Thus theIF signals shall be digitized in the samplers in thelower cabin of the telescope before being transferred to the VLBI backend system far away fromthe telescope, with keeping the instrumental excess path length stable during a VLBI observation.
Moreover, HINOTORI has installed the new highspeed signal recording systems, VSREC and OCTADISK2. They shall record the signals at ratesof 2×n Gbps, where n = 2 − 24, with a base-bandchannel width of 512 MHz .
3.

Status of HINOTORI development

In the ﬁrst stage of HINOTORI, we have developed a system to simultaneously receive 22 and43 GHz band signals [3]. The ﬁrst light of this system was yielded in 2018 April and has been in scientiﬁc operation since 2019 December. The ﬁrst
user project of this system, FLASHING (Finest
Legacy Acquisitions of SiO-/H2 O-maser Ignitions
by Nobeyama Generation) [2] is intensively monitoring circumstellar H2 O and SiO masers for ﬁnding periodic behaviors and secular evolutions of thespectra of maser sources associated with high velocity jets so called “water fountains”.

32

IVS NICT-TDC News No.39

Figure 1. Overview of the HINOTORI system, which describes the signal ﬂows in it. Color shaded
components are provided from HINOTORI. Two perforated dichroic ﬁlter plates can be remotely inserted
to and released from the quasi-optics so as to transfer sets of the wanted frequency band signals. Four
receives named H22, H40/Z45, and TZ receive radio signals in K-band (20–25 GHz), Q-band (42.5–
44.5 GHz/40–45 GHz), and W-band (80–116 GHz), respectively. The W-band signals shall be transferredto the backend in the frequency range of 85–93 GHz (lower side band) when the dichroic ﬁlter plate isinserted. In VLBI observations in the East Asia VLBI Network (EAVN) open-use program, analogue
intermediate frequency (IF) signals are transferred to the digital sampler ADS3000+ (provided by NICT)
for digital sampling at a rate of 4 Gbps in total in the operation building. In new VLBI experimentswith the HINOTORI backend, the IF signals are digitized with two OCTAD samplers in the lower cabinof the telescope, yielding a higher recoding rate (48 Gbps in maximum) and higher fringe phase stability.
The triple-band receiving system has completed
by 2020 September through re-commissioning the
80–116 GHz receiver, TZ, with its major repair in
the SIS mixer and new development for its remotecontrol system[4]. TZ itself can receive the upper and lower side band signals, while HINOTORI
accepts only the latter signals due to the limited
passes to the VLBI backend system and the characteristics of the dichroic ﬁlter plate. Note that TZ
shall receive radio signals in two orthogonal linear
polarizations. These signals will be converted tocircular polarization signals by either using 1/4-λplate or post-processing in the software that deals
with cross-polarization correlation data. The latterapproach is still in development.
The IF signals for VLBI were injected at proper
power levels in 2020 September after installing thesignal ampliﬁers[4]. Although the two OCTAD
samplers can transfer the digitized IF signals at
a rate of 48 Gbps in total, recording at the maximum data sampling rate is in a future work because
recording with the new VLBI recorders is still being tested.

HINOTORI has also developed the script thatenables us to operate VLBI observations remotelyin the new HINOTORI system. This script shallconvert a vex ﬁle provided by the principal investigator of the VLBI observation to a set of an execution ﬁle for the telescope (.start) and parameterﬁles the frequency setup (.ndevice and .dat). It
may need to be updated so as to handle a varietyof the frequency setups.
4.

Future perspectives

HINOTORI will reach its completion by the endof 2021 March as planed. Actualy, there still remain further development items. They are in remote IF selection, recording at the maximum sampling rate, polarization handling, and real-time
spectroscopy using a PC (PolariH) for quick lookand measurement of system noise temperatures in
the whole baseband channels. For improving thesensitivity in the 43 GHz band by using the Z45
receiver (receiving radio signals in two orthogonallinear polarizations) in the 22/43 GHz simultane-

January 2021

33

supported by NICT Kashima Space Research Center for installing its ADS3000+ sampler in NRO.
References[1] Dodson, R., et al., Astrometrically registeredsimultaneous observations of the 22 GHz H2 O
and 43 GHz SiO masers toward R Leonis Minoris using KVN and source/frequency phasereferencing, AJ, 148, 97 (9 pages), 2014[2] Imai, H., et al., FLASHING: New highvelocity H2 O masers in IRAS 18286−0959,
PASJ, 72, 58 (11 pages), 2020.
Figure 2. Perforated dichroic ﬁlter plate installedin the Nobeyama 45 m telescope. It transmits radio signals in 85–93 GHz, which shall be received
with the TZ receiver, and reﬂects those in otherfrequency ranges, received with H22 and H40 receivers.
ous observation system, another dichroic ﬁlter isneeded to be installed (in Mirror #7 in Figure 1).
For development of the band-to-band phasetransfer technique for the 45 m telescope, a calibration scheme for diﬀerential instrumental delays
among the diﬀerent IF signals is an open question.
Eventually we need to often observe delay calibrators. We will further need to consider to install asystem to inject reference signals for such calibration.
Even with the challenging issues in the HINOTORI system, it will become a key component in
future VLBI observations in the 86 GHz band in
East Asia, with KVN, and the Paciﬁc Ocean region with JCMT, Greenland Telescope (GLT), andATCA. This style of future collaboration is underdiscussion over 10 years, but coming true with the
beginning of scientiﬁc operation of the HINOTORI
system.
Acknowledgments
HINOTORI is one of the Development Proposal
projects technically supported by Nobeyama Radio Observatory (NRO) and its proposed VLBI experiments have been conducted as commissioning
works approved by Mizusawa VLBI Observatory
and KVN/KASI. HINOTORI has been promotedby the team members from Kagoshima University,
Yamaguchi University, Osaka Prefecture University, and the National Astronomical Observatory ofJapan, with the supports by KAKENHI 16H02167
and 18H03721 and that partially by Amanogwa
Galaxy Astronomy Research Center, Kagoshima
University. HINOTORI and EAVN have also been

[3] Okada, N., et al., Development of a 22/43GHz-band quasi-optical perforated plateand dual-band observation system of the
Nobeyama 45-m Telescope, PASJ, 71, 7 (23pages), 2020.
[4] Tsutsumi T., in master thesis of Yamaguchi
University, 2021 (in Japanese).
[5] Yoon, D.-H., et al., Astrometrically registeredmaps of H2 O and SiO masers toward VX
Sagittarii, Nat. Comm., 9, 2534, 2018.

34

IVS NICT-TDC News No.39

On-the-Fly Interferometer Ex- beam/phase center, the amplitude shows the beamand the phase should change linearly withperiment with the Yamaguchi pattern
time.
Interferometer
Since the integration time of the raw data is
Kenta Fujisawa1 (kenta@yamaguchi-u.ac.jp),
Kotaro Niinuma2 , Masanori Akimoto2 ,Hideyuki Kobayashi3(1) The Research Institute for Time Studies,
Yamaguchi University, Yoshida 1677-1,
Yamaguchi-city, Yamaguchi, Japan(2) Graduate School of Science and Technology,
Yamaguchi University, Yoshida 1677-1,
Yamaguchi-city, Yamaguchi, Japan(3) National Astronomical Observatory of Japan,
2-21-1, Osawa, Mitaka, Tokyo, Japan
Abstract: We report on an interferometric observation technique for scanning the wide ﬁeld of thesky. When an interferometric observation of radio
astronomy is carried out for a ﬁeld wider than the
beam of a single telescope, it is usual that imageof each small ﬁeld is synthesized in an ordinarymanner and combine the images after the observation. We are studying an On-the-Fly interferometer (OtFI) observation method that continuously
changes the beam position of the telescope and thephase center of the correlation process. So far, wehave succeeded in scanning the beam and the center of phase, consequently detecting the fringes ofmultiple sources in 1 square degree ﬁeld.
1.

short, the signal-to-noise ratio of the data is low,
and the amplitude information alone can provide
positional accuracy equivalent to the angular resolution of the main beam. A time window is set
to this data and a Fourier transform is performed,
then the temporally integrated and the S / N is
improved, and consequently the positions of radiosources can be estimated with the accuracy of theinterferometer.

Figure 1. The scan pattern of this experiment.

Introduction

The goal of the On-the-Fly interferometer mapping method is to detect radio sources whose positions are unknown in a wide ﬁled of the sky with aradio interferometer. While continuously scan the
beam of the element antenna in the sky (On-theFly), it operates as an interferometer to continuously detect radio sources in the ﬁeld. Then, it
will be possible to detect and list radio sources inthe ﬁled with positional accuracy of the interferometer. In this research, OtFI is tested with the
Yamaguchi interferometer (YI) to demonstrate thefunction and the eﬀectiveness of this new method.
2.

Goals of this Experiment

The ﬁrst step of this experiment is to successto scan the beam of the telescope and the phasecenter of the correlation processing continuously ina ﬁeld of the sky. The target point of the beamand the phase center is changed at every second,
and the correlation data were output every second
too. When passing through a radio source with the

Figure 2. The observed ﬁeld of 1 square degree.
This image is obtained by NVSS (Condon & Kaplan 1998). The brightest source at the south partis J1055+201. The second brightest source at thenorth part is also detected in this experiment.

January 2021
3.

35

Experiment

The OtFI experiment with the YI is carried out
on September 27, 2020, 03:00:00–05:05:00 UT. The
Yamaguchi Interferometer is a single baseline interferometer with Yamaguchi 32m and Yamaguchi
34m radio telescopes. The baseline length is 109 m.
The observed ﬁeld is centered at RA of 10h 58m
17.9008s and Dec of +19d 51’ 50.869” (J2000.0),
and ranged 1 degree both in RA and Dec. There
are some radio source in this ﬁeld. J1055+201 is
the brightest one with ﬂux density of 2 Jy at 1.4
GHz. The second brightest source is 400 mJy at
1.4 GHz. The scan direction is in RA with a speed Figure 4. Enlarged plot of 5500 to 6000 sec of ﬁgureof 15 arcsec/sec. The scan was repeated 30 times 3. Two weak peaks are seen at 5660 and 5900 sec
changing the declination of 2 arcmin at each scan. which are scans of the second brightest source.
The scan pattern is shown in ﬁgure 1, and the radio image of the ﬁeld by NVSS is shown in ﬁgure
2. The observing frequency is 8192-8704 MHz.
4.

Results

The fringe amplitude with time is shown in ﬁgure
3. Three prominent peaks around 1800 to 2300
sec are from three scans passed over J1055+201.
Two other peaks are also seen at 5660 and 5900
sec which are from the second source. The enlarged
plot of this part is shown in ﬁgure 4.

Figure 5. Time averaged plot of 5500 to 6000 sec.
Two peaks are clearly seen.
sources, and (ii) to calibrate the ﬂux density of thedetected sources.
References[1] Condon, J. J. & Kaplan, D. L. 1998, Astronomical Journal, 115, 1693

Figure 3. Fringe amplitude with time. Three
prominent peaks are scans passed over J1055+201.
Figure 5 shows the time averaged plot of the data
shown in ﬁgure 4. The averaging time of 16 seconds
are set to the beam (FWHM) passing time. Twopeaks of the second brightest source are clearly detected. At this moment, the peak height is not correct because the averaging was performed ignoring
the phase rotation. By correcting the phase rotation and performing the averaging operation, theS / N is further increased, and the third brightest
source can be detected.
We have succeeded in the scan and detection of
the radio sources. The next step to be carried out is(i) to derive the accurate positions of the detected

36

IVS NICT-TDC News No.39

Superconducting
wide-band material in our work, and the HTS ﬁlter minimizes
loss at an upper boundary temperature of 77
BRF for Geodetic VLBI Ob- its
Kelvin or lower. Under 77 Kelvin, the HTS ﬁlter
servation with VGOS Radio can be designed according to the operating temTelescope to prevent radio perature.
frequency interference
Ken Takahashi (ken25.takahashi@toshiba.co.jp),
Kazuhiro Tsubosaki, Hirotaka Kamide,
Hiroshi Kudou, Taro Asakura, andMasahiro Ruika
Figure 1. Schematic block diagram of receiverToshiba Hokuto Electronics Corp. Business
Development Dept., New Business PlanningGroup,
1-1-1 Shibaura, Minato Tokyo, Japan

Spectrum mask requirements for the BRF are
indicated in ﬁgure 2. This ﬁlter has a hex-band
transmit response for rejecting RFI. Required inAbstract: The Toshiba Group has developed wide- sertion loss is less than 0.5dB. The attenuation ofband (1-14GHz) multi-band (hex-band) BRF us- RFI is more than 15dB, 22dB, 32dB at each RFI
ing high-temperature superconductor (HTS) . The frequencies. In 2.1-2.4GHz, some of RFI is very
VLBI Global Observing System (VGOS) speciﬁca- close to passband. So narrow bandwidth and sharp
tions require wide-band operation of the receiver cut oﬀ ﬁlter characteristics are needed.
from 2 to 14 GHz. And in the L-S band, there is
Deviations in group delay of desgined ﬁlter wereradiofrequency interference (RFI) from other wire- veriﬁed. The simulation result was conﬁrmed to beless systems. We conﬁrmed that new developed less than 5nsec.
BRF achieves required attenuation at each RFI frequencies and transmit response. This ﬁlter was
adopted at the wide-band receiver in the Ishioka
VGOS station of the Geospatial Information Authority of Japan.
1.

Introduction

Superconducting ﬁlter can realize both sharp
cutoﬀ characteristics and low insertion loss. Sharp
cutoﬀ characteristics realize narrow band ﬁlter.
Low insertion loss characteristics realize ﬁlters withhigh ﬂexibility in design, such as BPF, BRF andthose mixed Multi-band ﬁlter.
We have developed HTS ﬁlter[1] and small highsensitibity receiver using HTS ﬁlter[2][3] for radio astronomy. In this article, we introduce the
new HTS BRF for wide-band VGOS receiver(2-14
GHz).

Figure 2. Spectrum mask requirements for multiband BRF

3.

Experimental Result

Figure 3 shows measured response of the BRF.
Spurious response frequencies which are generated
in the high frequency side (7-8GHz and 13-14GHz)
are tuned to avoid observation band. Insertion loss
2. Design Specification of BRF
is less than 0.5dB in the 1-10.5GHz, which satisﬁes required speciﬁcs. Insertion loss at high freFigure 1 outlines a block diagram of the receiver
quency edge (10.5-10.7GHz) exceeding 0.5dB andincluding BRFs. The cryostat with the feed horn,
within 0.61dB, which is not the best but acceptBRFs and cryogenic low noise ampliﬁers (LNAs) toable. At frequencies above 10GHz, unavoidable inbe cooled down to 20 Kelvin in a vacuum chamber.
sertion loss of SMA connector, which is used input
The receiver is equipped with two channels for sepand output port of BRF module, is added to meaarated vertical and horizonal polarization signals.
sured response.
To prevent RFI eﬀects, BRFs are placed at the inFigure 4 shows measured detail response of rejecput of the LNA at each channel. Yttrium bariumtion band(countermeasure for RFI). The attenuacopper oxide (YBCO) is used as a superconductortion of RFI is more than 15dB, 22dB, 32dB at each

January 2021

37

Figure 5. BRF ﬁlter module
Figure 3. Measured response of BRF(1-14GHz)
RFI frequency. We conﬁrmed that this response of
rejection band satisﬁes required speciﬁcs.

References[1] T. Kawaguchi, N. Shiokawa, K. Nakayama,
M. Yamazaki, H. Ikeuchi, H. Kayano, DualBand High Temperature Superconductor Filter for Radio Astronomy Station, IEICE technical report, Vol. 114, No. 11, pp. 7–10, 2014
(in Japanese).
[2] T. Kawaguchi, H. Kayano, K. Tsubosaki, H.
Kamide, T. Asakura, HTS Quad-band Highsensitivity Receiver for a 4.5-m Radio Telescope, 2019 49th European Microwave Conference (EuMC), pp. 248–251, 2019.

Figure 4. Detail of rejection band(countermeasure
for RFI)
Figure 5 shows a photograph of BRF module.
The BRF module measures approximately 79 mm
long, 88 mm wide and 20 mm high. A small module size will allow it to adapt even to the existingreceiver.
4.

Summary

We have developed the new HTS BRF for wideband VGOS receiver(2-14 GHz). The results ofthe experiments conﬁrmed that it attenuates morethan 15 dB, 22 dB, 32 dB at each RFI frequency
and insertion loss of the passband is less than 0.5
dB in the 1-10.5GHz, less than 0.61 dB in the 10.510.7GHz. This ﬁlter is expected to realize highsensitivity wide-band VGOS receiver.
5.

Acknowledgments

We would like to thank the Geospatial Information Authority of Japan for their RFI data andrequired ﬁlter speciﬁcations.

[3] T. Kawaguchi, K. Tsubosaki, H. Kamide, T.
Asakura, HTS L-S band High-sensitivity Receiver for a 20-m Cassegrain Radio Telescope,
2019 IEEE Asia-Paciﬁc Microwave Conference(APMC), pp. 1360–1362, 2019.

38

IVS NICT-TDC News No.39

Status and Future of the Mizusawa 10m Radio Telescope
Osamu Kameya (osamu.kameya@nao.ac.jp),
Mareki Honma, Tomoaki Oyama,
Kazuyoshi Asari, Aya Yamauchi,
Yu Asakura, Gen Sato, Yuki Matsukawa,Kazuya Hachisuka, Shunsaku Suzuki,
Toshio Terasawa, Kazuyoshi Yamashita,
VERA development group, VERA
operations group, VERA maintenancegroup, JASMINE group, and Balloon
VLBI group
Mizusawa VLBI Observatory, National
Astronomical Observatory of Japan, 2-12
Hoshigaoka-cho, Mizusawa-ku, Oshu, Iwate
023-0861, Japan
Takeshi Fukusako, and Shotaro Haruguchi
Faculty of Advanced Science and Technology,
Kumamoto University, Kurokami South W2,
2-39-1, Kurokami, Kumamoto 860-8555, Japan
Abstract: The Mizusawa 10m radio telescope has
been used as a VLBI station and a test bench of
new VLBI and radio experiment since 1992. Now
we have a plan to use it for as the downlink stationof the Nano-JASMINE (Japan Astrometry Satellite Mission for INfrared Exploration） near future.
We also think a new plan to use it for the low frequency observation which is below 2 GHz.
1.

Introduction

Figure 1. The Mizusawa 10m radio telescope.
2.

Recent status of the 10m radio telescopeWe
haveinstalled
S-bandreceiver,22GHz/43GHz receiver on the telescope, andwe can observe both S-band and the other one
band simultaneously.
While the budget for
Mizusawa VLBI Observatory has been reduced
from last year, there is no prospect of operating
costs from this year onward. Therefore, we areconsidering operation and development withcompensation from individual research expenses,
scientiﬁc research expenses, and others. We arealso responding with the help of everyone involved.
Over the past year, the 10m radio telescope has
had some problems due to its aging. That is, theair conditioning in the receiver room was stopped,the LAN in the receiver room was malfunctioning, the K-band compressor was stopped, theperformance of the optical transmission device
was deteriorated, the OS of the antenna drive
computer (Windows7) should be updated, and the
S band reception performance was deteriorated.
These have been resolved individually.

The Mizusawa 10m radio telescope is 28 years
old now, and it has been used for Japanese VLBI
Network (J-Net) which is the ﬁrst VLBI network
for common use observations [1]. After the construction of the VERA (VLBI Experiment of Radio
Astrometry) it has been used for many new VLBI
experiments: daily observations of Sgr A* at 22
GHz [2], phase diﬀerential VLBI test observations
for VERA, and RISE (Research in selenodesy; [3]).
It will be used for test observations for the Balloon
VLBI experiment. It has been used for single dishdaily monitor observations for Orion-KL [4], andeducation of radio astronomy for students. It isplanned to be used as the downlink station of the
Nano-JASMINE (Japan Astrometry Satellite Mission for INfrared Exploration）satellite near futurewhen the launch system becomes available [5]. Ourgroup has a plan to use the 10m radio telescope forthe low frequency observation, which is below 2 3.
GHz, and want to use it for pulsar and fast radio
burst observations.

Current status andfrequency receivers

futureoflow-

The purpose of the low-frequency observations issecuring radio telescopes capable of low-frequency

January 2021

39

Figure 3. Schematics of the L-band system

Figure 2. The L-band system.

Figure 4. Schematics of the lower frequency system.

observation of pulsars and transient sources withan eye on SKA, etc. We consider the following
three future items: (1) Improvement of the existing
system in the S-band, (2) Development of L-band
system, (3) Aiming for the lower frequencies.

signal comes to the observation building by using
an E/O and O/E and optical ﬁber line and directly
sampled by backends. The schematics of the Lband system is shown in Fig.3. At present, interference is strong, and the performance of the receiver
3.1 Improvement of the existing system insystem is not suﬃcient. It is necessary to optimize
the S-band
the level and improve the horn position. In the fuThe purpose is to enable observation of the gi- ture, when a good-performance L-band system isant pulse of the crab pulsar with high temporal constructed on this telescope, simultaneous obserresolution. Therefore, RF is sent to the observa- vation with a large-diameter radio telescope suchtion building using E/O and O/E without using as FAST would be expected.
the existing down converter, and RF is recorded
directly in the observation building. Currently, we 3.3 Aiming for the lower frequencies
are re-inspecting the system that has already been
Observations of the frequency lower thaninstalled. Since the performance of the horn deteri- 650MHz will be challenging and very interestingorated due to rust, the rust was removed to improve because many pulsars have strong emission at suchit.
low frequency. In Japan, Iidate station of To3.2

Development of L-band system

This system was developed by Prof. Fukusako
and his student in Kumamoto University. The purpose is to study an L band receiving system for thistelescope. Initially, we considered 1.4 GHz to 1.6
GHz, but changed it to 1.05 GHz to 1.45 GHz in
consideration of the FIRST frequency. For use inthis radio telescope, we considered improving gain
and reducing cross-polarization [6]. We will show
the schematics of the L-band system. L-band RF

hoku University can observe such low frequency.
If we make the low frequency system on the 10m
radio telescope, we can conduct some simultaneous observations with the Iidate and the other low
frequency stations. We have planned to installsuch low frequency using Yagi-Uda antennas. The
planned schematics is shown in Fig. 4.
4.

Summary

The Mizusawa 10m radio telescope has been used
as a VLBI station and a test bench of new VLBI

40

and radio experiment for more than 28 years. This
telescope is planned to be the down link stationof the Nano-JASMINE satellite near future. We
also have plans to use this telescope at the lower
frequency which is lower than 2 GHz including Lband and lower than 650MHz.
References[1] Omodaka, T., O. Kameya, M. Miyoshi, T.
Miyaji, T. Sasao, T. Kondo, Y. Koyama, J.
Nakajima, M. Sekido, E. Kawai, H. Osaki, H.
Okubo, M. Imae, Y. Takahashi, N. Kurihara,
H. Takaba, T. Iwata, N. Kawaguchi, K. Iwadate, H. Kobayashi, K. Shibata, S. Kaneko, Y.
Murata, and H. Imai, Operations and Resultsof J-Net (Japanese Domestic VLBI Network),
Journal of the Communications Research Laboratory, Vol.48, No.1, 23-29, 2001.
[2] Kameya, O., M. Tsuboi, Y. Asaki, Y.
Yonekura, Y. Miyamoto, H. Kaneko, M.
Seta, N. Nakai, M. Miyoshi, H. Takaba, K.
Wakamatsu, Y. Fukuzaki, T. Morimitsu, M.
Sekido, 22 GHz MONITORING OF SGR A*
DURING THE G2 PERI-CENTER PASSINGWITH THE JAPANESE VLBI NETWORK,
Publications of The Korean Astronomical Society, vol. 30, issue 2, pp. 649-650, 2015.

IVS NICT-TDC News No.39[3] Hanada, H., T. Iwata, N. Namiki, N. Kawano,
S. Sasaki, K. Matsumoto, H. Noda, S. Tsuruta, K. Asari, T. Ishikawa, F. Kikuchi, Q.
Liu, S. Goosens, Y. Ishihara, N. Petrova, Y.
Harada, K. Shibata, K. Iwadate, O. Kameya,
Y. Tamura, X. Hong, J. Ping, Y. Ali, S.
Ellingsen, and W. Schluter, The Explorationof Lunar Gravity by VLBI Observations of SELENE (Kaguya), Trans. JSASS Space Tech.
Japan, Vol. 7, pp.Tk 1-Tk 5, 2009.
[4] Horiuchi, S. and O. Kameya, Highly polarized emission from the bursting water maserin Orion-KL, Publ. Astron. Soc. Japan, 52,
545-550, 2000.
[5] Gouda, N., T. Yano, Y. Yamada,Y.
Kobayashi, T. Tsujimoto, and the JASMINE Working Group, Japan Astrometry
Satellite Mission for Infrared Exploration
ESA RSSD (Research and(JASMINE),
Scientiﬁc Support Department),
2004,
URL: http://www.rssd.esa.int/SA/GAIA/
docs/Gaia_2004_Proceedings/Gaia_2004_
Proceedings_77.pdf[6] Haruguchi, S., Dual-polarized Horn Antenna
with Low-proﬁle Structure, Master s thesis,
Kumamoto University, 2019.

January 2021
- News - News - News - News -Kashima 34-m Antenna
Closing Ceremony
Mamoru Sekido (sekido@nict.go.jp )NICT Space-Time Standards Laboratory,
Kashima Space Technology Center,893-1 Hirai, Kashima Ibaraki, Japan1.

A brief hisotory of Kashima 34-m antenna and VLBI researh

Kashima 34-m antenna closing ceremony washeld on 3rd Oct. 2020, which was about one year
after the stop of operation. Elevation drive mechanics and part of reﬂector structure was damaged by strong typhoon Faxai attacked to Kashima
on 9th Sept. 2019. That was decided to be dismantled in 2020, and that work is underway now(as of January 2021). History of Kashima SpaceTechnology Center of NICT goes back to 1964.
The Kashima branch of Radio Research Laboratory (RRL; Name has been changed to CRL in
1988 and to NICT in 2001) was established in 1964
for technology development of satellite communication and broadcast. Steerable parabola antennasof 30 m diameter (1964) and 26 m antenna (1968)
were built for technology development in that ﬁeld.
Television picture of the Olympic games held inTokyo 1964 was transferred by satellite link from
Kashima to Ponit Mugu station of California in
USA, then broadcasted in the USA and Europe.
The 1960s was era of historical big discoveries inradio astronomy, such as quasar (1963), Cosmic microwave background (1965), and pulsar (1967). Radio astronomers in Japan focused on these antennas and had visited Kashima to use these high gain
parabola antennas for radio astronomical observations, and then collaboration has started between
RRL and National Astronomical Observatory. In
1971, VLBI research at Kashima has started by anoﬀer of collaboration about VLBI proposed fromNASA Goddard Space Flight Center (GSFC). A
VLBI group was formed in Kashima in 1975, andthe ﬁrst domestic VLBI experiment was successfully performed in 1977. Since then, Kashima 26m antenna (Fig. 1) joined international geodetic
VLBI experiments conducted by NASA. These experiments resulted the detection of contraction of
baseline length between Kashima and Kauai, andit provided a proof of plate tectonics. Based on thegreat success of geodetic VLBI with 26-m antenna,
Kashima VLBI group has made big progress. The
Kashima 34-m antenna was built in 1988 as one of41

radio telescopes of Western Paciﬁc VLBI Network(WPVN) project, which is targeting measurementof plate motions around the Japan Islands.
Not limited to geodetic measurement of paciﬁcplate motion, Kashima 34-m antenna has madecritical role in many technology developments.
Three generations (K3, K4, and K5) of VLBI observation systems were developed in Kashima, andthey have been used as standard geodetic VLBI
system in Japan. Our activity as the IVS technology development center has been published bythis IVS NICT-TDC news (https://www2.nict.
go.jp/sts/stmg/ivstdc/news-index.html ).
2.

Kashima 34-m antenna closing ceremony

Closing ceremony of the Kashima 34-m VLBI
station was held as the memorial event on 3rd
Oct. 2020. Originally it was planned in the endof April 2020 having about 200 attendees including invited individuals and citizens. However, due
to COVID-19, that was postponed to October andit was changed to hybrid form with small number of attendee and online broadcast to the public
via YouTube (https://youtu.be/kXaWl3zDP8g).
The ceremony started from the opening addressby the president of NICT. Then, video messages of invited participants were showed. Although, the invited guests (Prof. Axel Nothnagel
as the chair of IVS directing board, Prof. Zhiqiang
Shen as director of Shanghai Astronomical Observatory, Ed. W. Himwich from NASA/GSFC,
Prof. Hideyuki Kobayashi of NAOJ, and Satoshi
Fujiwara as director of GSI/Space Geodesy Division) could not attend due to the disaster, they
contributed the ceremony by kindly sending video
or letter messages. Following to the ceremony,
public lectures were given by Mizuhiko Hosokawa(a former member of directors’ board of NICT),
Prof. Kosuke Heki (Hokkaido Univ.) and Takashi
Takahashi (director of Kashima Space Technology
Center (KSTC)). Hosokawa talked about achievement of Kashima 34-m VLBI station in the ﬁeldof geodesy and astronomy. Prof. Heki used to bea member of Kashima VLBI group, and he gave
lecture on recent results of space geodesy. Takahashi talked on past and future of satellite communication technology in relation with KSTC. Morethan 13 old VLBIers, who used to be worked at
Kashima, got together for the ceremony and we
said goodbye to the Kashima 34-m antenna. The
video of the whole event is still available from
YouTube (https://youtu.be/kXaWl3zDP8g).

42

IVS NICT-TDC News No.39

Figure 1. Top left: Kashima 26-m antenna (1968–2002) played an important role in the history ofgeodetic VLBI in Japan. Kashima 34-m antenna was built on the success of this antenna. Top right:
Raibow above the antenna on 12th Nov. 2001. Bottom: Kashima 34-m antenna on 5th Jun. 2006.

Figure 2. Group photo of the attendees to the ceremony on 3rd Oct. 2020.

January 2021

43

“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.39, January 2021
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