ASTRO-F Observer’s Manual
Version 3.2— for Open Time Observation Planning —ASTRO-F User Support Teamin
Institute of Space and Astronautical Science / JAXA
contact: iris help@ir.isas.jaxa.jp
European Space Astronomy Centre / ESA
contact: http://astro-f.esac.esa.int/esupport/

November 29, 2005

Version 3.2 (November 29, 2005)

i

Revision Record
2005 Nov. 29:

2005 Nov. 8:

2005 Sep.20:

Version 3.2 released.
Updated description of IRC FoV and slit (Section 5.1.4 and 5.1.5).
Updated IRC04 detection and saturation limits.
Also improve the description (Section 5.5.9).
Section 5.4.2 revised to clarify the point.
Units for Ho given value in p.113.
Version 3.1 released.
Updated Visibility Map for Open Time Users (Figure 3.4.3)
Information for the handling of Solar System Object observations (Section 3.4)
Updated saturation limits for FTS (FIS03) mode (Table 4.4.16)
IRC04 detection limits for NG+Np added (Table 5.5.25,5.5.26)
Updated worked examples using the latest versions of the Tools (Section B)
ESAC web pages URL and Helpdesk contact address updated
Numerous errors and typos corrected
Version 3.0 released.

Contents
1 Introduction
1.1 Purpose of this document . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
1.2 Relevant Information . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

1
1
3

2 Mission Overview
2.1 The ASTRO-F Mission . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.2 Satellite . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.2.1 The Bus Module . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.2.2 Attitude Determination and Control System . . . . . . . . . . . . . . . . .
2.2.3 Cryogenics . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3 Telescope . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.1 Speciﬁcation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.2 Pre-ﬂight performance . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.4 Focal-Plane Instruments . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.4.1 Speciﬁcation Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.4.2 Focal-Plane Layout . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.4.3 The Focal-Plane Star Sensor (FSTS) . . . . . . . . . . . . . . . . . . . . .
2.5 The ASTRO-F Observation Programms . . . . . . . . . . . . . . . . . . . . . . .

5
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3 Satellite Operations
17
3.1 Orbit and Observing Attitude . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 17
3.2 Attitude Operation Modes for Observations . . . . . . . . . . . . . . . . . . . . . 19
3.2.1 Survey Mode . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 19
3.2.2 Pointing Mode . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 19
3.3 Mission Phases . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 22
3.4 Sky Visibility . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 23
4 FIS: Far-Infrared Surveyor
4.1 Hardware Speciﬁcation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.1 Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.2 Optics and Filters . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.3 Detector System . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.4 The Fourier Transform Spectrometer (FTS) . . . . . . . . . . . . . . . . .
4.1.5 Instrument Operation . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2 Pre-ﬂight performance . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2.1 Optics . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2.2 Detector system . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2.3 FTS mode . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.3 The FIS All-Sky Survey . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.3.1 Operation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
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4.4

4.3.2 Expected Performance . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
The FIS Pointed Observations: AOTs . . . . . . . . . . . . . . . . . . . . . . . .
4.4.1 FIS01: Photometry / Mapping of Small Area . . . . . . . . . . . . . . . .
4.4.2 Expected Performance of FIS01 . . . . . . . . . . . . . . . . . . . . . . . .
4.4.3 FIS02: Mapping of large area . . . . . . . . . . . . . . . . . . . . . . . . .
4.4.4 Expected Performance of FIS02 . . . . . . . . . . . . . . . . . . . . . . . .
4.4.5 FIS03: FTS Spectroscopy . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.4.6 Remarks and Constraints of the FIS Observations . . . . . . . . . . . . .

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57
58
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63

5 IRC: Infrared Camera
65
5.1 Hardware Speciﬁcation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 66
5.1.1 Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 66
5.1.2 Optics . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 67
5.1.3 Filters and Dispersion Elements . . . . . . . . . . . . . . . . . . . . . . . 68
5.1.4 Field of View (FoV) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 69
5.1.5 Slits for Spectroscopy . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 69
5.1.6 Detectors . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 70
5.2 Pre-ﬂight performance . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 71
5.2.1 Optics . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 71
5.2.2 Detector system . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 78
5.3 The IRC Instrument Operation . . . . . . . . . . . . . . . . . . . . . . . . . . . . 80
5.3.1 Pointed Observations . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 80
5.3.2 The IRC Scan Observation . . . . . . . . . . . . . . . . . . . . . . . . . . 82
5.4 The IRC All-Sky Survey . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 83
5.4.1 Operation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 83
5.4.2 Expected Performance . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 84
5.5 The IRC AOTs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 85
5.5.1 Detection Limits and Saturation Limits: General Remarks . . . . . . . . . 86
5.5.2 IRC00: Deep Imaging Mode . . . . . . . . . . . . . . . . . . . . . . . . . . 88
5.5.3 Expected Performance . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 89
5.5.4 IRC02: Imaging with Two Filters . . . . . . . . . . . . . . . . . . . . . . . 90
5.5.5 Expected Performance . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 91
5.5.6 IRC03: Imaging with Three Filters . . . . . . . . . . . . . . . . . . . . . . 92
5.5.7 Expected Performance . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 93
5.5.8 IRC04: Spectroscopic mode . . . . . . . . . . . . . . . . . . . . . . . . . . 94
5.5.9 Expected Performance . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 96
5.5.10 IRC11: Slow-Scan mode . . . . . . . . . . . . . . . . . . . . . . . . . . . . 98
5.5.11 Expected Performance . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 99
5.6 Notes and Restrictions for the IRC Observations . . . . . . . . . . . . . . . . . . 100
6 Data Reduction and Products
101
6.1 Basic Policy . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 102
6.2 Large Area Surveys . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 103
6.2.1 All-Sky Survey . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 103
6.2.2 ASTRO-F Point Source Catalogues . . . . . . . . . . . . . . . . . . . . . . 103
6.2.3 Image Maps . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 104
6.3 Pointed Observation Data . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 105
6.3.1 FIS Data Reduction: AOT FIS01, 02 . . . . . . . . . . . . . . . . . . . . . 105
6.3.2 FIS Data Reduction: AOT FIS03 . . . . . . . . . . . . . . . . . . . . . . . 105
6.3.3 IRC Data Reduction: AOT IRC00, 02, 03 . . . . . . . . . . . . . . . . . . 106

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iv
6.3.4
6.3.5

IRC Data Reduction: AOT IRC04 . . . . . . . . . . . . . . . . . . . . . . 106
IRC Data Reduction: AOT IRC11 . . . . . . . . . . . . . . . . . . . . . . 107

A Confusion Limits and ASTRO-F
109
A.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 109
A.2 Confusion due to Infrared Cirrus . . . . . . . . . . . . . . . . . . . . . . . . . . . 109
A.3 The Extragalactic Source Confusion Limit . . . . . . . . . . . . . . . . . . . . . . 110
B ASTRO-F Cookbook for Open Time Observations
115
B.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 115
B.2 Example 1: Imaging of extragalactic ﬁelds with the IRC and FIS . . . . . . . . . 116
B.2.1 Scientiﬁc Background . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 116
B.2.2 Selection of Target Field and Target Visibility . . . . . . . . . . . . . . . . 117
B.2.3 Preparation of Target List . . . . . . . . . . . . . . . . . . . . . . . . . . . 119
B.2.4 Target List Validation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 123
B.2.5 Duplication Check Tool . . . . . . . . . . . . . . . . . . . . . . . . . . . . 125
B.2.6 Estimation of Sensitivity for IRC . . . . . . . . . . . . . . . . . . . . . . . 127
B.2.7 Estimation of Sensitivity for FIS . . . . . . . . . . . . . . . . . . . . . . . 127
B.2.8 Submission of Proposal . . . . . . . . . . . . . . . . . . . . . . . . . . . . 131
B.3 Example 2: Spectroscopy of Planetary Nebulae with the IRC . . . . . . . . . . . 132
B.3.1 Scientiﬁc Background . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 133
B.3.2 Target Visibility . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 134
B.3.3 Preparation of Target List . . . . . . . . . . . . . . . . . . . . . . . . . . . 135
B.3.4 Target List Validation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 137
B.3.5 Duplication Checker . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 138
B.3.6 Estimation of Sensitivity for IRC . . . . . . . . . . . . . . . . . . . . . . . 139
B.3.7 Submission of Proposal . . . . . . . . . . . . . . . . . . . . . . . . . . . . 141

Chapter 1

Introduction
1.1

Purpose of this document

This ASTRO-F Observer’s Manual, is issued as the reference for ASTRO-F observers, for observation planning and the preparation for the data analysis. The current version (3.0) is mainly
used for the preparation of Open Time Proposals.
The document covers the speciﬁcations of the satellite and the instruments, their performance, ﬂight operation plan, including deﬁnitions of the Astronomical Observation Templates(AOTs), and the data reduction system. It contains the most up-to-date estimates of the detection limits of the instruments (based on the performance evaluation tests carried out over thelast few months). This version (version 3) contains more a detailed introduction to the mission
and instruments, guidance to the users on how to prepare their observations, and the scope forthe data reduction.
Information about the scientiﬁc programmes and proposal handling are not included in thisdocument. Users are suggested to refer to other information resources (see the next page).
Note that some of the instrumental performance information was not yet available by thetime of the publication of this edition version. Frequent updates will be given on the Observer’s
support Web page.

1

2

ASTRO-F Observer’s Manual

Figure 1.1.1: ASTRO-F on the vibration test bench (May, 2005).

Version 3.2 (November 29, 2005)

1.2

3

Relevant Information

ASTRO-F Observer’s Web
The ISAS Web page contains some additional information and data useful for the observers. Also links to the user support tools are given:
URL:http://www.ir.isas.jaxa.jp/ASTRO-F/Observation/

The ESAC page includes mission overview, fact sheet, and user support policy for theEuropean Open Time proposals. Web tools are also available.
URL:http://astro-f.esac.esa.int/
Helpdesk
Any questions and comments on ASTRO-F observations and user support are addressedto the ASTRO-F Helpdesks.
iris help@ir.isas.jaxa.jp: for Japanese & Korean Open Time Users and general inquiry.
http://astro-f.esac.esa.int/esupport/ : for ESA Open Time Users

4

ASTRO-F Observer’s ManualChapter 2

Mission Overview
2.1

The ASTRO-F Mission

ASTRO-F (previously known as IRIS) is the second Japanese space mission for infrared astronomy. The satellite will be launched by JAXA’s M-V rocket. The launch is expected to takeplace in January–February 2006 1 . The orbit is a Sun-synchronous polar orbit with an altitudeof 745 km and a period of 100 minutes. The cryogen lifetime is expected to be 550 days. A
pair of Stirling cycle mechanical coolers enables this long lifetime with only 170 liters of liquid
Helium. Even after the liquid Helium has boiled-oﬀ, observations at near-infrared wavelengths
can be continued.
ASTRO-F is equipped with a 68.5 cm cooled telescope and two scientiﬁc instruments, namely
the Far-Infrared Surveyor (FIS) and the Infrared Camera (IRC). The FIS has two 2-dimensional
detector arrays and observes in four far-infrared bands between 50 and 180 µm. The IRC consistsof three cameras covering 1.7–26 µm in 9 bands with ﬁelds of view of approximately 10 × 10 .
Both instruments have low- to moderate-resolution spectroscopic capability.
A major goal of the mission is to carry out an All-Sky Survey with the FIS and additionallywith the IRC at 9 and 18 µm. The ASTRO-F All-Sky Survey will signiﬁcantly surpass theprevious all-sky survey in the infrared by IRAS in both spatial resolution and sensitivity. The
results will be published as ASTRO-F infrared source catalogues. In addition to the survey
observations, ASTRO-F will carry out thousands of dedicated pointed observations with both
the FIS and IRC. Deep imaging as well as spectroscopic observations are available. Two highvisibility regions, the North Ecliptic Polar (NEP) region and the Large Magellanic Cloud (LMC)
will be observed intensively by the pointed observations. These surveys are operated by the
ASTRO-F Project in a similar vein to the All-Sky Survey. Many pointed observations are used
for the legacy observations by the ASTRO-F project team members. Such programmes arereferred to as “Mission Programmes”. In addition, a signiﬁcant fraction of pointed observation
opportunities will be open for general users in Japan/Korea and ESA related countries.
Many international collaborations are ongoing with the ASTRO-F project. The European
Space Agency (ESA) will supply a ground station operated by ESOC, and ESAC will carryout the pointing reconstruction of the All-Sky Survey observations. ESAC will also handle theuser support for the European Open Time observing programmes. A consortium of Imperial
College University of London, University of Kent, Sussex University, and SRON-Groningen with
the University of Groningen (IKSG consortium) participates on the data reduction of the FIS
All-Sky Survey. Seoul National University representing the Korean community also joins thedata reduction activity. Several collaborations at the personal level are also ongoing especiallyon the celestial calibrators.
1

Precise launch schedule will be given later this year.

5ASTRO-F Observer’s Manual

6

2.2

Satellite

Figure 2.2.1 shows the overall structure of the ASTRO-F spacecraft. The hight is about 3.7 m(excluding the aperture lid) and the launch weight will be 960 kg. The satellite consists of twoparts, the bus module and the cryostat. The two parts are connected by a truss structure.

20

00

2170

Telescope

Sun Shield

Stirling-Cycle Cooler
Star Trackers
Focal-Plane Instrument Electronics

Sol

ar P

add

le

S-band Antenna

60
Thrusters

S-band Antenna
Earth Sensor

1550

18

930

Figure 2.2.1: An overall view of the ASTRO-F spacecraft (scales are in mm).

2.2.1

The Bus Module

The bus module takes care of the power supply, house keeping, spacecraft attitude control, dataacquisition and the telecommunication link to the ground station. ASTRO-F uses the S-band(mainly for command and house-keeping telemetry data) and the X-band (mainly for scientiﬁcdata) for telecommunication. Due to its’ near-Earth orbit, communication between the satellite
and the ground stations is limited. Therefore, ASTRO-F has a two-Gigabyte data recorder.

Version 3.2 (November 29, 2005)

2.2.2

7

Attitude Determination and Control System

The attitude of the ASTRO-F satellite is determined and controlled by on-board sensors and acomputer unit. Figure 2.2.2 shows the block diagram of the attitude and orbit control system(AOCS). The framework of the system is as follows: The IRU (Gyro) measures the motion of
the satellite. The TFSS (Two-dimensional Fine Sun Sensor) monitors the position of the Sun
in two-dimensions perpendicular to the Sun-satellite direction, and corrects any long-term driftof the IRU signal. A pair of Star Sensors (STT) mounted on the wall of the cryostat are used
to observe the attitude along the third axis and the alignment change between the bus module,
where IRU and TFSS are located, and the cryostat in which the telescope is installed. Other
instruments provide additional information for attitude determination. A dedicated computer(AOCU) will determine the satellite attitude. A set of reaction wheels (RW) will control the
spacecraft attitude, with the help of Magnetic Torque Rods (MTQ). In the case of any verylarge required motion, a set of thrusters may be used, but this is not expected during the survey(until liquid Helium boil oﬀ).

STT-A

RW-A

STT-B

RW-B

TFSS
IRU-A
IRU-B
NSAS
CES-A

RW-C

AOCU

RW-D

AOCP-AMTQ-X

AOCP-B

MTQ-Y

AOCP-C

MTQ-Z

CES-B
RCS20Nx4

GAS

RCS 3Nx4

ACM
Sensors

Data Handling Unit

Drivers

Figure 2.2.2: Block diagram of the ASTRO-F attitude determination and control system.

Table 2.2.1: ASTRO-F AOCS Performance
Absolute Pointing Control Accuracy
Scan Speed Stability (in-scan)
Scan Speed Stability (cross-scan)
Pointing Stability

30 arcsec< 0.1 per cent3 × 10−5 arcsec/min< 1 arcsec/min (peak-to-peak)ASTRO-F Observer’s Manual

8

2.2.3

Cryogenics

Scientiﬁc instruments are all stored in the cryostat and maintained at cryogenic temperaturesby 170 liters of super-ﬂuid liquid Helium with the help of two sets of Stirling-cycle coolers. The
expected lifetime of the cryogen is 550 days. The near-infrared instrument can be operated withthe mechanical coolers even after the liquid Helium boil-oﬀ. ASTRO-F will be operated with
the Sun-shield always directed to the Sun. The cryostat is covered by silver coated ﬁlm thatreﬂects visible light but radiates in the infrared.
The Scientiﬁc Instrument Assembly (SIA), namely the telescope and the Focal-Plane Instruments (FPI) are thermally shielded by two layers of Vapour Cooled Shield (VCS) in which
Helium vapour evaporated from the tank is running through. The inner VSC is also cooled by
the mechanical coolers. The SIA is mainly cooled by the Helium vapour down to about 6 K.
The FIS body and two detectors are directly connected to the Helium tank via thermal straps
and are operated at 2.0–2.2 K.
The two mechanical coolers are designed to work symmetrically to cancel the momentum
that may interfere with the spacecraft attitude control system. High-frequency noise remains,
and we recognize that it interferes with the FTS mode observation by the FIS. See Section 4.1.4
for details.

Figure 2.2.3: The ASTRO-F Cryogenic system in the clean room for maintenance. (Left) The
Cryostat. (Right) Close-up of the Two-Stage Stirling cooler. (Photo: July, 2005)

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2.3

Telescope

2.3.1

Specification

9

The telescope is a Rithcy-Chrétien system with an eﬀective aperture of 68.5 cm and a focal ratioof 6.1 (Figure 2.3.4). The mirrors are made of SiC porous+CVD coated material. This rigidmaterial allows us to reduce the weight of the mirrors signiﬁcantly. The primary mirror, witha physical diameter of 71 cm, weighs only 11 kg. The total weight of the telescope system is
42 kg. The focus may be adjusted in orbit during the PV phase by moving the position of thesecondary mirror. The parameters for the telescope are summarized in Table 2.3.2.
The primary mirror was replaced in 2004 as a part of a refurbishment of the telescope system.
The eﬀective aperture has been increased slightly (previously 67 cm). The telescope has beentested in the cryogenic vibration tests, and the mirror support system which previously had aproblem has been conﬁrmed to be secure.
The telescope of ASTRO-F is cooled mainly by Helium vapour evaporated from the tank.
The nominal operation temperature is 5.8 K.

Figure 2.3.4: The ASTRO-F telescope system.

Table 2.3.2: Telescope Speciﬁcation
Eﬀective diameter
Focal length
Focal ratio
Optical design
Temperature

68.5 cm (primary mirror)
420 cm
6.1
Ritchey-Chrétien type
∼ 5.8 K

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2.3.2

Pre-flight performance

Mirror surface accuracy has been measured at cryogenic temperatures (as low as 9 K). Due tothermal stress at the supporting points, the wave front error becomes larger at low temperaturesand is at its’ worse at around 50 K, but improves slightly beyond that.
The analysis of the cryogenic tests of the telescope predict that the system reaches thediﬀraction limit at 6.2 µm, slightly longer than the original design goal (at 5 µm).
The focus position of the telescope has been measured accurately in cryogenic optical tests,
and the telescope will be adjusted to the best position at the launch. Further focus adjustments
may be carried out in orbit during the performance veriﬁcation (PV) phase.
No indication of stray light has been observed, though the tests were not intended to simulate
the complete ﬂight environment. This will also be measured in the PV phase.

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2.4

Focal-Plane Instruments

2.4.1

Specification Overview

11

Two scientiﬁc instruments are placed in the Focal-Plane of the telescope; the Far-Infrared Surveyor (FIS) and the Infrared Camera (IRC). In addition, two sets of Focal-Plane Star Sensors(FSTS) are installed for observing guide stars in the Near Infrared band (∼ J-band) for pointing
reconstruction during the survey observations.
The FIS is equipped with two sets of two-dimensional detector arrays. The Short-Wavelength
channel (SW; 50–110 µm) uses two monolithic Ge:Ga arrays2 of 20 × 3 and 20 × 2 pixels, while
the Long-Wavelength channel (LW; 110–180 µm) has a stressed Ge:Ga array of 15 × 3 + 15 × 2
pixels. Each of the four arrays correspond to a diﬀerent wavelength band. The pixel sizes are
27 and 44 arcsec2 , respectively.
The IRC consists of three independent cameras; the NIR (2–5 µm), the MIR-S (5–13 µm),
and the MIR-L (13–26 µm). The NIR camera has a 512 × 412 pixel InSb array, of which
100 × 412 pixels are dedicated for slit spectroscopy. The two MIR cameras use 256 × 256 pixel
Si:As detectors.
Table 2.4.3: Observation Capabilities of ASTRO-F.
Channel

IRC

FIS

2

Band

N2
N3
NIR
N4
NP
NG
S7
S9W
MIR-S
S11
SG1
SG2
L15
L18W
MIR-L
L24
LG1
LG2
SW
N60
WIDE-S
LW
WIDE-L
N160
FTS

Wavelength(µm)
1.7–2.7
2.7–3.7
3.7–5.5
1.7–5.5
2.5–5.0
5.8–8.4
6.5–11.6
8.6–14.1
5.5–8.3
7.4–13.0
12.4–19.4
13.9–25.3
20.4–26.5
17.7–25.0
50–80
60–110
110–180
140–180
60–180The SW detector was manufactured by NICT.

Resolution(λ/∆λ)
3.6
2.8
3.4
22@3.5 µm
135@3.6 µm
2.8
2.1
2.3
47@6.6 µm
34@10.6 µm
2.5
1.8
4.3
Not Available
27@20.2 µm
2.2
1.8
2.0
4.0
150–450

Detector Array(pix2 )

FoV
(arcmin2 )InSb
512 × 412

9.5 × 10.0

Si:As
256 × 256

9.1 × 10.0

Si:As
256 × 25610.3 × 10.2

Ge:Ga 20 × 2
Ge:Ga 20 × 3
Stressed Ge:Ga 15 × 3
Stressed Ge:Ga 15 × 2
WIDE-S + WIDE-L

10 × 1.5
10 × 1.0
12 × 2.5
12 × 1.6

ASTRO-F Observer’s Manual

12

The wavelength range and resolving power provided by the scientiﬁc instruments are shownin Figure 2.4.5. The spatial resolution and the Field of View (FoV) of the instruments are
summarized in Figure 2.4.6. The detection limits are illustrated in Figure 2.4.7.

IRC

1000

FTS
Prism & Grism

100
λ/∆λ

FIS

10

Wide-band PhotometryWide-band (Imaging)

1
1

10
Wavelength (µm)

100Ver. 05/08/15 I.Yamamura

Figure 2.4.5: The wavelength coverage and resolving power of the ASTRO-F instruments.

1000

IRC

FIS

Field of View

Field of View
(cross-scan)arcsec100
10Telescope Limit

Pixel Size

Limit
ction
DiffraPixel Size

1
1

10
Wavelength (µm)

100Ver. 04/07/27 I.Yamamura

Figure 2.4.6: Pixel size and ﬁeld of view (FoV) of the ASTRO-F on board instruments.

Version 3.2 (November 29, 2005)

13

Figure 2.4.7: Detection limits of the ASTRO-F onboard instruments in various observing modes.

ASTRO-F Observer’s Manual

14

2.4.2

Focal-Plane Layout

Figure 2.4.8 shows the Focal-Plane layout projected onto the sky. There are three light entrances
to the instruments; one for the FIS and two for the IRC. The two instruments essentially canobserve simultaneously, but they see diﬀerent areas of the sky as shown in the ﬁgure. Therefore,
observations of a sky position with diﬀerent aperture have to be made on diﬀerent orbits.
The incident infrared radiation coming into the FIS aperture is divided into two spectral
domains by a dichroic beam splitter, as a result of which the FIS-SW and LW detector arrays
observe almost the same sky position. A small diﬀerence remains due to diﬀerence in pixel sizeof each detector array (see Figure 4.1.3). Similarly, the NIR and MIR-S cameras of the IRC
share the same FoV so as to observe the same sky position. Only the MIR-L channel observesa diﬀerent sky position.
The scan path is approximately parallel to the Ecliptic meridian line. The scan direction
depends on the launch conﬁguration, namely the launch time (morning or evening). The current
plan is a morning launch, with the FoV moving downward, i.e., stars go through the detector
from bottom to top on the ﬁgure.
Note that the real position and the shape of each Focal-Plane Instrument on the sky will beslightly diﬀerent from those shown in Figure 2.4.8, due to optical aberration. See Section 4.1.3(FIS) and 5.1.4 (IRC) for details.

FSTS-SSCAN
DIRECTION.3'

1226.5deg
2.5' 6.5'

38.7 deg

FIS

IRC

9.4' 0.6'

7.0' 2.5'

NIR
&
MIR-S

8.2'

44.4deg

7.0'

MIR-L
10.7'

8'

Telescope
Axis
10.0'

10.2'

25.0'

0.8'
for all the FSTS detectors

19.2'

18.

20.5'

8.2'

7.0'

12.4'

44.4deg
'

.3

12

23.3'

20

04

Oc

t. 5

FSTS-L

Figure 2.4.8: Focal-Plane layout projected onto the sky. The distortion due to the telescope andinstrumental optics have to be taken into account for the data reduction. See the instrument
pages for details.

Version 3.2 (November 29, 2005)

2.4.3

15

The Focal-Plane Star Sensor (FSTS)

A pair of star sensors are installed in the Focal-Plane. The FSTS-L consists of four Ge detector
elements3 , while the FSTS-S has three. They observe stars with their intrinsic wavelength band
proﬁle close to the J-band (1.25 µm). The readout is a TIA (TransImpedance Ampliﬁer) circuitwith a sampling rate of about 50 Hz. The seven detectors are sampled almost simultaneously.
As the satellite scans the sky, stars cross the detectors in the in-scan direction one after another. The detection interval between a pair of two horizontal detectors placed in parallel deﬁnes
the scan speed, and the interval between the tilted and horizontal detectors can be converted
to the cross-scan position. The two FSTS’s located in opposite corners provide information on
the roll angle of the Focal-Plane with respect to the scan direction, as well as a redundancy
for the star detection. The absolute positional information is given by comparing the detection
information with the position reference catalogue. The pointing reconstruction using the FSTS
assumes (all-sky) survey operation mode. It is not decided whether the FSTS data can be used
for the Slow-Scan operation for pointed observations.
Although the FSTS is controlled via the FIS electronics and software, and the data is transfered within the FIS data telemetry packet, the data may not be provided to the users for analysis. The pointing reconstruction processing is carried out by a dedicated team at ESAC/ESA
and the results are provided to the All-Sky Survey data reduction team.

3
One of the four elements is known to be dead. The eﬀect of this loss to the pointing reconstruction is expected
to be small.

16

2.5

ASTRO-F Observer’s Manual

The ASTRO-F Observation Programms

Observations with the ASTRO-F can be divided into three categories.
Large Area Survey (LS): Survey programmes of wide areas of the sky which use the ASTROF capabilities most eﬃciently. Three observing programmes are included in this category;
the All-Sky Survey (LSASS), the North Ecliptic Pole Survey (LSNEP) and the Large
Magellanic Cloud Survey (LSLMC). The latter two surveys will use the pointed observation
mode. Planning and data processing of these programmes are led by dedicated teams in
the ASTRO-F Project, and the products will be released to the communities.
Mission Programmes (MP): Seventy per cent of the pointed observation opportunities areallocated to the ASTRO-F team members as a kind of guaranteed observation time. Speciﬁc working groups in every scientiﬁc ﬁeld have worked to establish the observing programmes. In total 15 Mission Programmes are assigned. The data will be used by theteam members for a one year prioritized period and then will be opened to the public. The
MP observations, as well as the LS are regarded as the legacy of the ASTRO-F mission.
Open Time Programmes (OT): The remaining 30 per cent of the total pointing opportunities are available to the general research community. Two thirds of the opportunities(= 20 per cent of the total) are reserved for Japanese/Korean astronomers, and 10 per
cent are for the ESA member countries. Details of the Call for Open Time Proposals are
explained in a separate document.

Chapter 3

Satellite Operations
3.1

Orbit and Observing Attitude

The orbit of ASTRO-F is a Sun-synchronous polar orbit along the twilight zone. The nominal
altitude is 745 km, corresponding to an orbital period of 100 min. Once ASTRO-F is brought
into this orbit in the initial maneuver operation, no adjustments or changes of the orbit during
the mission will be performed in principle until the boil-oﬀ of the liquid Helium, in order toavoid contamination by the thruster gas.
In principle, the spacecraft always points in the direction perpendicular to the Sun–Earthline, keeping the Sun-shield towards the Sun, and rotates once per orbit to look in the opposite
direction to the Earth. The spacecraft attitude in the cross-scan direction1 has a ﬂexibility of±1 degree about the canonical scan path.
89 - 91 degfrom the Sun
Return to
Survey( 7.5 min )

Sun
Light

Pointing(10 min )Stabilize( 5 min )Attitude Change
Maneuvour( 7.5 min )EARTHContinuous Survey

Figure 3.1.1: The in-orbit attitude of ASTRO-F (left) and the pointed observation maneuver(right).
Observations by ASTRO-F will be carried out either in survey mode or pointing mode. In
the survey mode, the satellite is operated such that the telescope scans the sky with a constant
speed (3.6 arcmin / sec) usually along the great-circle at a solar-elongation of 90 degrees. In the
pointing mode, the telescope can point at an area of the sky along the scan path for up to 10
1

the direction perpendicular to the scanning (in-scan) direction.

17

ASTRO-F Observer’s Manual

18

minutes, with the total cost of 30 minutes including maneuver for the operation. Fine controlof the observing position in a pointing operation is possible. Observations in the pointing mode
will be deﬁned by a set of Astronomical Observation Templates (AOT; see Sections 4.4 and 5.5).
Table 3.1.1: ASTRO-F Orbit.
AltitudeOrbital Period

745 km (−0 km +145 km)
100 min

The satellite attitude and orbit control system (AOCS) has the capability of determining/ controlling the absolute attitude to an accuracy of 30 arcsec. The dominant source of thepointing uncertainty is the alignment of the focal-plane axis with respect to the AOCS reference
frame on the bus module and its time variation due to thermal distortion of the spacecraft.
Table 3.1.1 summarizes the nominal operation parameters and expected performance of the
ASTRO-F AOCS for astronomical observations.

Version 3.2 (November 29, 2005)

3.2

Attitude Operation Modes for Observations

3.2.1

Survey Mode19

In the survey mode, ASTRO-F performs a continuous scan of the sky. The spacecraft spinsaround the Sun-pointed axis once every orbit, keeping the telescope pointed toward the opposite
direction to the Earth (see left-hand of Figure 3.1.1); as a result, the telescope traces a great
circle with a solar elongation of 90 deg. The orbital period of 100 minutes corresponds to ascan-speed of 3.6 arcmin s−1 . In Table 3.2.2, the speciﬁcation and performance of the survey
mode operation are summarized.
Table 3.2.2: Summary of Survey mode operation (nominal).
Speciﬁcation of Survey mode operation
Survey Scan Speed
3.6 arcmin s−1
Survey In-scan Stability1≤ 0.1 per cent
Survey Cross-scan Stability1∼ 3 × 10−5 deg/sec (3σ)
Pointing determination error (goal)
On-board< 30 arcsec (in-scan and cross-scan)
G-ADS
15–30 arcsec (in-scan and cross-scan)
Pointing Reconstruction (During Mission) < 5 arcsec (in-scan), < 7 arcsec (cross-scan)
Pointing Reconstruction (Post-Mission)< 3 arcsec (in-scan), < 5 arcsec (cross-scan)
1 Based on the performance evaluation tests.

The accurate position of the survey scan path on the sky will be determined during thepointing reconstruction processing. There are three levels of the processing: on-board determination by the attitude and orbit control system (AOCS), on-ground processing using the samedataset (G-ADS), and detailed analysis together with the Focal-Plane Star Sensor (FSTS) and
IRC survey data.
The outline of the attitude determination process is summarized in Figure 3.2.2. The onboard attitude analysis has an absolute error of ∼ 30 arcsec. The main source of this error
is the alignment uncertainty between the satellite AOCS and the telescope’s focal-plane. The
results of the on-board analysis are downloaded to the ground with the sensor signal data. The
data is re-analyzed by the G-ADS for cross-check and to prepare the initial pointing data forreconstruction. At this stage, the pointing data (in the survey mode) still has an error of 15–30
arcsec.
Then, the data from the FSTS and the IRC survey will be analyzed together with the initialpointing data for reﬁnement (pointing reconstruction). The ESA European Space Astronomy
Centre (ESAC) is in charge of the ﬁnal pointing reconstruction. The goal of the pointing accuracy
after the reconstruction is 5 arcsec (in-scan) and 7 arcsec (cross-scan) during the survey. Thisaccuracy will be improved to 3 arcsec (in-scan) and 5 arcsec (cross-scan) when the survey iscompleted and all data are analyzed and reprocessed.

3.2.2

Pointing ModeIn the pointing mode, ASTRO-F will observe a speciﬁc sky position. Both imaging and spectroscopic observations with longer exposure times than the survey are possible. A single pointed
observation consists of four continuous operations (Figure 3.1.1, right);
1. Slewing maneuver to point the telescope at the target position,

ASTRO-F Observer’s Manual

20

Control

On-Board
ADS

FSTS
Data

AOCS
Sensors

IRC ScanData

Realtime~ one day

ISAS
Ground ADS
Pointing
Obs. data

1~2 weeks

All-Sky Survey
Data Reduction
On-flyCatalogue" (in-scan)7" (cross-scan)
After Phase 2

All-Sky Survey
Data Reduction
ASTRO-F
Catalogues

ESA
Pointing
Reconstruction~

Improvements

Pointing
Reconstruction3" (in-scan)5" (cross-scan)Figure 3.2.2: The outline of the ASTRO-F attitude determination for the All-Sky Survey, including time scales and envisioned accuracy improvement. Note that the pointing reconstruction
by ESA will not be carried out for the pointed observations.
2. Stabilization of the telescope,
3. Observation of the target,
4. Returning to the survey mode.
The total operation is completed in 30 minutes. This means that the maximum number ofpointed observations per orbit is three, if and when we can ﬁnd a suitable set of targets separatedby ∼ 120 deg. The actual net observation time per pointing is however limited to 10 minutes.
To increase the ﬂexibility of the observation opportunity for any target, the telescope can pointat any position on the sky in a stripe of 2 degree (±1 deg) perpendicular to the nominal scanpath. However, this allowance is much smaller than previous missions such as IRAS, ISO andSpitzer. Observers have to understand that the visibility of ASTRO-F to observe a speciﬁctarget is extremely limited.
In order to avoid the Earth-light, the satellite starts slewing back to the survey mode automatically if the avoidance angle exceeds the limit. This happens without notifying the instruments to complete the observation. As a result, the very last part of some observation data
may be lost in some unfortunate cases. In Table 3.2.3, we summarize the speciﬁcations of thepointing mode.
The FoV is allowed to shift by a small amount in a pointed observation in the way summarizedin Table 3.2.4. The Step-Scan and the Micro-Scan are essentially the same operation, being
diﬀerent only in the amount of shift for one step. For the Slow-Scan operation, the spacecraft
scans the sky with constant speeds of 1–15 arcsec/sec. An additional speed of 30 arcsec/sec may
be available if it is validated during the PV Phase.
These ﬁne control operations are all included in the AOTs, and users do not have to concern
themselves with them. An IRC observation uses Micro-Scan for dithering. The FIS imaging
observations are carried out in Slow-Scan mode in the in-scan direction with the Step-Scan used

Version 3.2 (November 29, 2005)

21

Table 3.2.3: Summary of Pointing mode operation (nominal)
Speciﬁcation of Pointing mode operation
Total operation time
30 minMaximum observation time
10 minAbsolute accuracy of pointing 30 arcsec
Cross-scan Oﬀset±1 deg
Performance
Pointing Stability1≤ 1 arcsec (peak-to-peak) in 1 minutes
1 Based on the performance evaluation tests.

to shift the FoV in the cross-scan direction. See the AOT description page of each instrument
for more details.
Table 3.2.4: Pointing adjustment during a Pointed observation.
Amount of shift
Direction Purpose
Micro-Scan
15–30 arcsec/step
arbitrary Dithering of IRC imaging
Step-Scan
3–7 arcmin/step
arbitrary Cross-scan shift of FIS imaging
Slow-Scan
1–15(–30) arcsec/sec arbitrary FIS imaging
Use of these operations are deﬁned in the AOT and little room for the observers to change.

The Step-Scan and Micro-Scan will cause a disturbance of the satellite attitude at the levelof a few arcsec. Therefore, each exposure has to be started after waiting for the stabilizationof the telescope attitude. The IRC waits for notiﬁcation from the AOCS and then starts an
exposure. This waiting time will be a dead time for observers. The dead time for Micro-Scan
and Step-Scan operations are typically about 30 sec and 40–50 sec, respectively. It is noted thatthe actual stabilization time varies depending on the given condition. The Slow-Scan operationneeds 30–40 seconds until the scan speed becomes stable after the change.
These dead times are taken into account in the AOTs of the IRC and the corresponding
sensitivity estimates. The amplitude of the disturbance is much smaller than the spatial resolution of the FIS. Hence the FIS may simply wait for a certain small amount of time after anoperation.
No pointing reconstruction using the Focal-Plane instruments is planned for the pointed
observations. The users will receive position information from the G-ADS. Since the source ofthe pointing uncertainty is the alignment of the telescope and the AOCS, it is expected that theaccuracy will be improved as the analysis progresses. The users can reﬁne the position using
the data themselves for the IRC imaging observations. It is under discussion if it is feasible tocarry out the pointing reconstruction for the FIS scanning observation.

ASTRO-F Observer’s Manual

22

3.3

Mission Phases

The mission lifetime of ASTRO-F, expected to be 550 days, is divided into three observation
phases after a performance veriﬁcation (PV) phase.
Phase 0 (PV Phase) The ﬁrst two months after the launch are assigned to the PV phase.
In this period, a check out of the spacecraft system and the Focal-Plane Instruments will
be made, and then the in-ﬂight performance will be evaluated and the initial calibrationof the Focal-Plane Instruments will be carried out. Test observations will be performed
accordingly.
Phase 1 Approximately the ﬁrst half year, by which time ASTRO-F will have scanned theentire ecliptic longitude is referred to as Phase 1. In this period, the Large Area Survey
Programmes, namely the All-Sky Survey and the NEP / LMC pointing surveys will be
carried out with the highest priority. Some Mission Programme observations are also
carried out.
Phase 2 This phase will continue until all liquid Helium evaporates (∼300 days for the nominalcase). Many pointed observations of Mission Programmes and Open Time proposals will
be performed as well as supplemental scan observations to complete the All-Sky Survey.
We expect that supplemental survey observations will be especially required and performednear the ecliptic plane, where the scan paths are relatively sparse. Since Phase 2 is expectedto last for 10 months, about 2/3 of the sky area will be observed twice in this phase. The
regions of second observation opportunity depend on the launch date and can be checkedwith available visibility maps and visibility tools (see section 3.4).
Phase 3 After the boil-oﬀ of the liquid Helium, the mechanical coolers will keep the temperature
low enough to observe with the IRC/NIR camera. Phase 3, where only the IRC/NIR
camera is operated, will continue until any on-board instruments (the NIR camera itself,mechanical coolers etc.) cease to function.

Version 3.2 (November 29, 2005)

3.4

23

Sky Visibility

As noted earlier, the direction of the telescope pointing is always perpendicular to the direction
toward the Sun with a ﬂexibility of ±1 deg in the cross-scan direction. Since ASTRO-F is in apolar-orbit, the visibility, deﬁned as the number of scan paths covering a particular sky position,
is a strong function of ecliptic latitude. Targets in the ecliptic pole regions are observable ona number of orbits, while targets near the ecliptic plane are visible by ASTRO-F for only twodays (29 orbits) in a half year at the most. In addition, visibility is also aﬀected by the SAA(South Atlantic Anomaly; where a large ﬂux of charged particles disturbs the observation when
the satellite pass through the region) and the Moon (the moonlight can enter the telescope
aperture. The avoidance angle is 33 deg).
Figure 3.4.3 shows the ASTRO-F visibility for Open Time observations from October 24,
2006 to August 21, 2007 (current estimates of Phase 2). The visibility map updated from
Observer’s Manual Version 3.0 indicates more practical numbers than previously given. The
visibility is calculated based on the simulation of observation scheduling. In the calculation,
the nominal orbit of the spacecraft, the ±1 deg cross-scan oﬀset operation, Moon avoidance
angle, and South Atlantic Anomaly are taken into account. The available resources for OpenTime observations are the reserved orbits (6 per 29 orbits) and the orbits that are not used byeither the Large Area Surveys (including complementary scans for the All-Sky Survey to satisfy
the 2 orbital passes over every point on the sky, the NEP and LMC surveys) or the Mission
Programme observations. At every point on a 20 arcmin resolution grid, 50 observations arerequested and the number of successfully allocated observations is regarded as the visibility.
The data are then smoothed with a 40 arcmin Gaussian to remove over detailed informationthat may change easily in the real scheduling. The number indicated by the tool is the best
estimate of the visibility at present, but may not necessarily be true in the actual mission
operation.
Observers should understand that opportunities to observe particular objects or ﬁelds with
ASTRO-F are strictly limited and are a strong function of the visibility.
It is strongly recommended to the observers to check the visibility at each target position. A
Web interface visibility check tool is provided, with which the observers can check the visibility
either interactively or in a batch ﬁle mode.
(http://www.ir.isas.jaxa.jp/ASTRO-F/Observation/vis/otvis/)

24

ASTRO-F Observer’s Manual

Figure 3.4.3: ASTRO-F Visibility for Open Time Observations during the Phase 2 period (current assumption: from October 24, 2006 to August 21, 2007) in Ecliptic coordinates. The
notations are the maximum possible number of observations that can be allocated at any position. In this map, the number given is limited to up to 50 by computational constraints.
The calculation was carried out on a 20 arcmin resolution grid which was then smoothed witha 40 arcmin Gaussian to remove over detailed information that may change easily in the realscheduling. The number indicated by the tool is the best estimate of the visibility at present,
but may not necessarily be true in the actual mission operation.

Version 3.2 (November 29, 2005)

25[Notes for the Solar System Observation]
Due to the severe visibility constraint, observation planning for Solar System Targets requires
additional procedures. The following steps are suggested.
1. Calculate the target position for corresponding observation dates. This can be done with
NASA/JPL’s HORIZON system (http://ssd.jpl.nasa.gov/cgi-bin/eph) by giving the condition of solar elongation angle as 90 ± 1 deg.
2. Use the original visibility tool (http://www.ir.isas.jaxa.jp/ASTRO-F/Observation/vis/, orfollow the link from the Open Time visibility tool page) to check if the target is reallyvisible at that position at that date.
3. Create the Target List with the coordinates. The corresponding observing date has to begiven in the comment ﬁeld of the Target List.
4. Check the visibility again with the Open Time visibility tool.

26

ASTRO-F Observer’s ManualChapter 4

FIS: Far-Infrared Surveyor
The FIS is designed primarily to perform the All-Sky Survey in four wavelength bands. The
instruments are operated such that data acquisition is continuously made as the telescope scans
the sky, resulting in sets of strip data of sky brightness. This operation can also be used forpointed observations in a Slow-Scan mode. The FIS is also equipped with a Fourier Transform
Spectrometer (FTS) that enables imaging spectroscopy over the full FIS wavelength range. FTS
observations are only made as pointed observations.

27

ASTRO-F Observer’s Manual

28

4.1

Hardware Specification

4.1.1

Overview

The speciﬁcations of the FIS instrument are summarized in Table 4.1.1. It provides four photometric bands between 50 and 180 µm, two broad bands and two narrow bands. Individual
detector systems are implemented for the two short and two long wavelength bands, respectively.
The FTS covers the entire FIS wavelength range with a resolution of ∼ 0.36 cm−1 (R = 450–150)
or ∼ 2.4 cm−1 (R = 75–23).
Table 4.1.1: Hardware Speciﬁcations of the Far-Infrared Surveyor (FIS)
Photometric mode
Band
N60
WIDE-S
WIDE-L
N160
Wavelength Range (µm)1 50–80
60–110
110–180
140–180
Central Wavelength (µm) 65
90
140
160
Band Width (µm)2
21.7
37.9
52.4
34.1
Detector
Monolithic Ge:Ga3
Stressed Ge:Ga
Array size
20 × 220 × 315 × 315 × 2
Operational Temperature∼ 2.2 K
∼ 2.0 K
Pixel size (arcsec)4
26.8
44.2
4
Pixel pitch (arcsec)
29.5
49.1
Readout
Capacitive Trans-Impedance Ampliﬁer (CTIA)
Spectroscopic mode
Spectrometer
Martin-Puplett type Fourier transform spectrometer
Wavelength (µm)
N/A
60–110
110–180
N/A−1
Resolution (cm )N/A
0.36 / 2.4
N/A
1 The detector responsivity is higher than 20 per cent of the peak.
2 As deﬁned by the bandwidth equation below
3 The SW detector was manufactured by NICT.
4 The values may change as more accurate telescope parameters become available.

The Central Wavelength is given to be a representative number of the band near the weightcenter of the band proﬁle. The band width is then deﬁned by the following equation for a ﬂat
spectrum, νFν =const.


r(ν)T (ν)Fν (ν)dνr(ν)T (ν)dν/ν=.
∆ν ≡r(νc )T (νc )Fν (νc )r(νc )T (νc )/νc

(4.1.1)Version 3.2 (November 29, 2005)

4.1.2

29

Optics and Filters

The FIS Optics
Figure 4.1.1 shows the bottom view looking out of the telescope (telescope and the sky are into
the page) from the FIS instrument. The light path is indicated by the dash-dotted arrows.
Optical elements are labeled. There are three moving parts, a shutter at the light entrance,
a ﬁlter wheel, and the moving mirror of the FTS. The ﬁlter wheel can be rotated to switch
between photometric and spectroscopic (FTS) modes. The shutter will be closed occasionallyto obtain instrumental dark levels.
Blocking Filters(σc~300cm-1 & 430cm-1)

Collimator Mirror

Bending Mirror

100mm

Input Aperture (with shutter)
Black-Body Source
MovableRoof-Top Mirror

Holeor
Input Polarizer

SCANNER

Mirror Transport Mechanism
Filter Wheel
Beam Splitter(Polarizer)

Dichroic Beam Splitter (σc~ 91cm-1)
or
Output Polarizer

Fixed
Roof-Top Mirror

SPECTROMETER

Camera Mirror

Camera Mirror

Low-Pass Filter (σc~90cm-1)
Low-Pass Filter (σc~70cm-1)

Weight: ~5.5kg

Calibration Sources

Low-Pass Filter (σc~220cm-1)
High-Pass Filter (σc~133cm-1)

Long Wavelength Detector(stressed Ge:Ga; < 1.8K)

Short Wavelength Detector(Ge:Ga; ~2.2K)

Figure 4.1.1: The optical paths within the FIS. Locations of the calibration lamps are indicated.

Filters and Photometric bands
Figure 4.1.2 describes the ﬁlter conﬁguration of the FIS photometric bands. The LW narrow
band ﬁlter was replaced in 2004 with a new ﬁlter with a cut-oﬀ at 70 cm−1 (143 µm). Thisextended the band proﬁle shortward and results in better eﬀective responsivity. Following thischange, we now denote the longest FIS band as “N160” instead of the previous name N170.
Calibration Lamps
Two sets of calibration lamps are installed in the FIS. The Cal-A system consists of two lamps
placed in front of the SW and LW detector arrays. They can be operated individually so thatthe optimal calibration procedure for each detector array can be realized. An important roleof the Cal-A lamp is to produce a pulse ﬂash to simulate the passage of a real point source.
The pulse ﬂashes are planned periodically (e.g., every minute) to monitor and correct the time
variation of the detector responsivity1 . Occasionally the lamps are continuously turned on toprovide a stable irradiation. In addition, the Cal-B system is a lamp located in the middle of
the FIS optical path. Its’ light passes through the beam splitter reaching both detector arrays.
This Cal-B system will be used as a backup for the Cal-A system. Both of these two sets of1

In the current calibration plan the duration of pulse is as long as 0.5 sec

ASTRO-F Observer’s Manual30
Incident photons

Blocking Filter ( c=430cm-1 & 300cm-1 )Dichroic Beam Splitter ( c=91cm-1 )LPF ( c=90cm -1)

LPF ( c=220cm-1 )LPF ( c=70cm-1 )(~HPF ( c=133cm-1)

N160
WIDE-L-14 m) (~-110 m)stressed Ge:Ga array

WIDE-S(110-60 m)

N60( -50 m)Ge:Ga array

for FTS mode
dichroic BS is replaced by output polarizer

Figure 4.1.2: The ﬁlter conﬁguration of the FIS photometric bands. LPF is the Low Pass Filter.
lamps are used solely for monitoring the relative responsivity variation of the detectors. The
absolute ﬂux calibration will be made with celestial objects.
In addition to the calibration lamp sets, there is a black body source on the inner wall. It is
used for calibration of the FTS.
Shutter
There is a shutter at the entrance aperture of the FIS. The instrumental dark level can bemeasured by closing the shutter. Leakage at the 1/200 (TBC) level has been observed. During
the All-Sky Survey we plan to close the shutter periodically to monitor the dark level.
The internal calibration sequence, calibration lamps and shutter operation for pointed observations are described in the AOTs (Astronomical Observation Templates).

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4.1.3

31

Detector System

ArraysThe FIS has two kinds of detector arrays, a set of two monolithic Ge:Ga arrays2 for the SWchannel (50–110 µm) and a stressed Ge:Ga array for the LW-channel (110–180 µm). They
are operated at ∼ 2.2 (SW) and ∼ 2.0 (LW) K respectively. Both the SW and LW channels
consist of two photometric bands; i.e, the FIS will have four bands in total, which are operated
simultaneously in scan mode.
Figure 4.1.3 shows the FIS detector arrays in detail on the Focal-Plane. The upper detectors
with three rows are the WIDE-S and WIDE-L bands, and the lower ones with two rows are theN60 and N160 bands. The FIS detector arrays are tilted with respect to the scan direction by
26.5 deg such that the interval between the neighbouring scan paths of the detector pixels is one
half of the physical pitch of the pixels. As noted above, the SW and the LW detector arrays
observe almost the same area on the sky. The slight diﬀerence in the sky coverage is due to thediﬀerence of the projected array size.

Figure 4.1.3: Focal-Plane layout of the FIS detectors.

Detector Readout Modes
The readout circuit of the FIS is a Capacitive Trans-Impedance Ampliﬁer (CTIA). The signal
from each pixel is accumulated in the capacitor as the detector observes far-infrared photons.
The voltage at the capacitor is then read out at appropriate intervals. The charge level in thecapacitor is periodically reset in order to avoid saturation. The accumulated signal between
two resets is called a Ramp. Diﬀerentiation of the Ramps provides the data that should be
proportional to the ﬂux of the incoming radiation. The reset operation of the FIS detectors is
2

The SW detector was manufactured by NICT.

32

ASTRO-F Observer’s Manualsynchronous, i.e., resets are inserted with a given time interval, regardless of the signal levelof each pixel. Although the reset interval is carefully optimized to maximize the instrument
performance, bright objects passing through the FIS FoV will possibly cause saturation of thedetector signal.
Three diﬀerent detector readout modes are considered.
Simple sampling mode is the basic operation sequence. Each sampled signal is downlinked
to the ground after A/D conversion.
Co-add mode samples the ramps at a higher rate, then the sums of every n samplings are√transfered to the ground. This can reduce the noise by a factor of ∼ n while preserving
the data size, (n is ﬁxed as 6 in the current design). This mode will most likely be
used in ﬂight unless it has unexpected side eﬀects. Hereafter we also refer to thismode as the Nominal mode.
CDS (Correlated Double Sampling) mode is designed to avoid saturation in very brightregions, such as the inner galactic plane. The detectors are reset at every two samplings,
and on board diﬀerentiation of these two samplings is retrieved as the signal. Therefore,
the downlinked data is already diﬀerentiated. Performance of this readout mode has
not been fully assessed. There is a possibility that observers will be asked tochange their plan according to the ﬂight performance evaluation.
Evaluation of the readout sequences has been made during the ground tests, but the ﬁnal
decision will be made during the PV phase. In principle we will use the the same operation modethroughout the survey. However, the possibility of switching the operation mode corresponding
to the sky brightness (galactic plane / other sky) has not been ruled out.
Bias Light for the LW detectors
The Ge:Ga detectors generally have a strong transient response, i.e. the detector does not
respond instantaneously to the incoming ﬂux but rather has a delay of a few – hundreds seconds
time scale (see Section 4.2.2). In the All-Sky Survey, a point source passes across a detector pixelin 1/4–1/7 seconds. This is too short for the detectors to reach a steady level corresponding to
the source ﬂux. Accordingly the point source sensitivity decreases. This eﬀect is particularly
serious for the LW detectors (stressed Ge:Ga detectors), and in the low ﬂux condition (dark sky+ weak sources).
In order to reduce the impact of this eﬀect, we decided to add a ‘bias light’. This takes
the form of a set of thin stainless wires put in front of the LW aperture. During the survey a
constant current will be used to give a steady photon ﬂux into the detector array. The point
source sensitivity in the All-Sky Survey is expected to recover by a factor of a few with this ‘bias
light’.
The SW detectors do not require such a mechanism.

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4.1.4

33

The Fourier Transform Spectrometer (FTS)

Note: The FTS mode is available to the Open Time users with several
constraints. Users are kindly requested to read the notes in Section 4.2 aswell as the Call for Proposals document.
The FIS provides the opportunity for imaging spectroscopy over its full wavelength range(60–180 µm) using the two wide-band arrays (WIDE-S and WIDE-L). The Fourier transform
spectrometer of the FIS is of the Martin-Puplett type (Figure 4.1.4). This polarizing type
Michelson interferometer uses input and output polarizers to create a linearly polarized beam,
and another polarizer to split the beams to a ﬁxed and moving mirror to make the interferogram.
The input and output polarizers are mounted on the ﬁlter wheel. These polarizers reduce theoptical eﬃciency of the instruments by 1/4 (ideal case) compared to the photometric mode. The
total optical eﬃciency is thus 10–20 per cent. The moving mirror, driven by an electromagnet,
can shift ±9.2 mm from its physical center. The optical center (zero path position) is displaced
by 4.6 mm from the physical center and the optical path length is asymmetric to the center.
Fixed Roof-top Mirror
Camera Mirror

Movable Roof-top
MirrorCollimeter
Mirror90LPF

Main Beam 430BKF&
Splitter(Polarizer) 300BKF

Bending
MirrorLong Wavelength
Detector (WIDE-L)
Short Wavelength
Detector (WIDE-S)

Liot Stop

Output Polarizer

Bending
MirrorInput Polarizer

218LPF

Bending Mirror
Camera Mirror

Figure 4.1.4: The optical diagram of the FIS Fourier Transform Spectrometer (FTS).
Two preset operation modes are provided for observers; Full-resolution mode and SED mode
(Figure 4.1.5). They are in principle diﬀerent only in the mirror scan path length, and accordingly in the resultant spectral resolution. The Full-resolution mode uses the full physical path
length, while the SED mode scans ±4.6 mm with respect to the optical center. The mirror
moving speed is always a constant 0.38 mm/sec. Taking an interferogram (scanning one-way)takes 48 sec in Full-Resolution mode and 12 sec in SED mode.
The spectral resolution, deﬁned as 1/L (L is the optical path length) is about 0.36 cm−1 for
the Full-Resolution mode and 2.4 cm−1 for the SED mode, respectively. Note that these values
assume apodization at the spectrum reconstruction. Better spectral resolution may be obtainedby optimizing the data reduction procedure. The speciﬁcation of the FTS mode is summarizedin Table 4.1.2.
In FTS mode, the WIDE-S and WIDE-L are operated simultaneously, with a higher sampling
rate than in the photometric mode in order to properly sample the interferogram. The narrow
band detectors, N60 and N160, are not used.

ASTRO-F Observer’s Manual34

Full-Resolution mode
9.2mm

= 1 / 2.76 cm

0.36 cm-1

9.2mm

0mm
0mm
48 sec-27.6mm-9.2mm

Optical PL

Mechanical PL

SED mode
4.25 mm
0 mm-4.25 mm

96 secDrive Speed = 0.383 mm/sec= 1 / 0.425 cm

2.4 cm-1

6.375 mm
2.125 mm
0 mm
48 sec96 secDrive Speed = 0.383 mm/sec

Optical PL

Mechanical PL

Figure 4.1.5: The mirror driving patterns of the FIS/FTS. There are two modes, Full-Resolution
mode and SED mode.

Table 4.1.2: Speciﬁcations of the FIS Fourier Transform Spectrometer (FTS) mode.
Wavelength (µm)
Operation mode
Mechanical Path Length (mm)Optical Path Length (mm)One scan time (sec)Resolution (cm−1 )Resolution λ/∆λ60–110 (WIDE-S), 110–180 (WIDE-L)
Full-ResolutionSED−9.2–+9.2+2.1–+6.3
27.6 mm
4.2mm
48
24
0.36
2.4
450–150
75–23

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4.1.5

35

Instrument Operation

The data acquisition process of the FIS is programmed into the hardware sequencer. When an
observation is started, a preset speciﬁed sequence pattern is repeatedly carried out and the data
are packed into telemetry packets and sent to the satellite’s main computer (Data Handling
Unit: DHU). The data sampling timing is as accurate as the FIS internal clock and is quite
stable. Detector level reset and calibration pulse are controlled by the FIS on-board softwarewith an accuracy as good as the readouts. They can alternatively be operated by commands.
Open and close of the shutter and continuous irradiation by the calibration lamps are madeby commands. The status of the shutter and calibration lamps are all recorded for the dataprocessing.
The survey observation sequence and the AOTs (Astronomical Observation Templates) aredesigned by combining these elemental operation units. Usually users are not aware of thedetails of the operation for the preparation of observations. However, such information may be
useful for the data reduction.

ASTRO-F Observer’s Manual

36

4.2

Pre-flight performance

The expected performance of the FIS instrument is discussed in this section. The information
is derived from the data taken during the FIS single component tests, the FPI evaluation tests
in combination with the IRC, and the performance evaluation tests in the ﬂight conﬁguration
during 2003–2005. Limitations of the test environment and data analysis mean that some values
noted here are still preliminary or have uncertainties. Users should bear this in mind for theobservation planning.

4.2.1

OpticsImaging Quality
It was impossible for the ASTRO-F team to carry out an end-to-end performance test for thetelescope and instruments in the cryogenic environment. Therefore, imaging quality throughsolely the FIS optics was measured. During the laboratory evaluation tests, a pin-hole mask
was placed at the focus of the instrument and the light from a blackbody source placed behind
was observed. The mask was smoothly moved along the short side of the detector (not titledas in the ﬂight conﬁguration) to simulate the scanning observation. Figure 4.2.6 shows the raw
image of the pin-hole passage observed by one of the rows (20 pixels) of the SW detector.

Figure 4.2.6: Pin-hole image observed by the SW detector. Signals from the 20 pixels in oneof the ﬁve rows are displayed. Each stripe indicates the time proﬁle of the output signal as the
pin-hole moves from top to bottom.
The signal proﬁle from a pixel as the pin-hole moves across it is presented in Figure 4.2.7.
The proﬁle is very symmetric with respect to the peak. In the SW detector, the width ofthe observed proﬁle is slightly larger than that expected from the optics. The reason for thisdiﬀerence is not well understood. Possible reasons are (1) cross-talk between the pixels, (2)unknown optical aberration, (3) the mask was slightly oﬀ-focus in the test conﬁguration.
The actual PSF of the LW detector is almost consistent with the model calculation.
Careful measurement of the PSF is planned during the PV phase.
Figure 4.2.8 shows the signal change of the LW detector as the pin-hole moves along thedetector array. When the image of the hole is located at the Column-3 pixels, the Column-4
pixels also output a signal. The possibility of electrical cross-talk may be ruled out. Concerning
optical ghosting, an aperture mask has been added in front of the detector mount to avoid

1

SW detector PSF

37

Measured(FWHM = 39)

Intensity [arbitrary]

Intensity [arbitrary]

Version 3.2 (November 29, 2005)

Calculated(FWHM = 32)

0.8
0.6
0.4
0.2
0

1

LW detector PSF

Calculated(FWHM = 52)

0.8
0.6
0.4
0.2
0

Pixel size

-120

-60

0

Position [arcsec]

Measured(FWHM = 55)

Pixel size

60

120

-120

-60

0

60

120

Position [arcsec]

Figure 4.2.7: The signal proﬁle for a pin-hole scan along a pixel in the WIDE-S (left) and WIDEL (right) array. It is recognized that the observed proﬁle (red) is broader than the expected one
from the model (blue). The physical size of the detector pixels is indicated in the plots.
reﬂection of the light. However, no pin-hole measurements were made after this installation andwe will not know whether the ghosting has been successfully removed until we will have theﬂight data.
The SW array also had a ghosting problem during the component evaluation tests due tothe reﬂection from a part of the arrays themselves. It has been successfully removed by adding
an aperture mask in front of the detector arrays to cover the reﬂective part.

Figure 4.2.8: The detector signal as a pin-hole moves along the detectors from upper-rightto lower left in Figure 4.1.3 in parallel to the detector rows. Columns 1,2,3 are the WIDE-L
detectors and 4 and 5 are of from N160. Ghosting features are indicated by arrows.

Spectral Responsivity
The relative spectral responsivity of the FIS system is presented in Figure 4.2.9. The responses
are normalized at their peaks. The ﬁlter transmittance as well as the eﬃciency of the opticalelements, and the detector spectral responsivity are taken into account. The ﬁlters and optics
were measured in the end-to-end system conﬁguration at room temperature. The narrow bandﬁlters for N60 and N160 were measured individually at the cryogenic temperature, since it was
known that these ﬁlters showed a temperature dependent variation.

ASTRO-F Observer’s Manual38

The detector spectral responsivity was evaluated using the FTS of the FIS itself, with ablackbody source at diﬀerent temperatures. Each pixel of the detector shows a diﬀerent spectral
response. This diﬀerence is larger in the LW detector, due to the non-uniform eﬀective stress on
the detector tips. The plots in Figure 4.2.9 are regarded as the ‘typical’ proﬁles. Pixel by pixel
proﬁles are being prepared for the calibration.

Relative Response

1.0

WIDE-S

0.8

WIDE-L

N160

N60

0.6
0.4
0.2
0.0
40

80

120

160

200

Wavelength (µm)
Figure 4.2.9: The spectral response curves of the four FIS photometric bands normalized at thepeak of each band. These proﬁles should be regarded as ‘typical’ for the bands.
The spectral responsivity data shown in the plots are also available on the ISAS user support
web page.
Leakage
A combination of two blocking ﬁlters at the entrance aperture are expected to reduce the shortwavelength light leakage to factors of 10−5 at 10 µm and 10−9 at 0.5 µm optimally. However,no end-to-end measurement has been made and the leakage could be as worse as 10−5 –10−6 at
0.5 µm.
Distortion
Figure 4.2.10 shows the detector shape projected onto the sky. Due to the fact that the FIS
aperture entrance is about 20 arcmin away from the center of the focal plane (boresight), andthat the FIS optics are an oﬀ-axis design, signiﬁcant distortion of the detector shape is observed(red lines). The shape without optical distortion is also plotted for comparison (blue). The
distortion will be further evaluated during the PV phase, and information necessary for thecorrection will be provided. It is not necessary for observers to take these eﬀects into accountin the observation planning.
Polarization
The dichroic beam splitter used in the FIS is known to have polarization. In Figure 4.2.11 thetransmittance and reﬂectivity of the ﬁlter is given. The reﬂection was measured at two diﬀerent

Version 3.2 (November 29, 2005)

39

FIS-SW Distortion Pattern [2005-June]
CORNER
1
2
4
3

InScan [arcmin]

3
4
5
1 2

3 4

5 6

7 8

2
1

15

1 2

3 4

5 6

7 8

9 10
1112
1314
1516
1718
1920

9 10
1112
1314
1516
1718
1920

CORNER
4

25

3

InScan [arcmin]

25

20

FIS-LW Distortion Pattern [2005-June]

WIDE-S

3
4
5
1

20

1
2
2

3

1

15

2

3

4

5

6

10
0

5

10
CrossScan [arcmin]

15

20

5

6

2
1

N60

10

4

0

5

7

8

9

7

10

8

11

9

12

10

13

11

14

12

13

WIDE-L
14

15

N170
15

10
CrossScan [arcmin]

15

20

Figure 4.2.10: The FIS detectors projected onto the sky, with (red) and without (blue) optical
distortion.
angles perpendicular to each other. A diﬀerence in performance at a few speciﬁc wavelengths at
approximately 160 and 180 cm−1 (corresponding to 63 and 55 µm, respectively) was observed.
Unfortunately the angle of the ﬁlter in the ﬂight module instruments is not accurately known,
and this diﬀerence will result in an uncertainty in the measurement of the polarized light.

Figure 4.2.11: The reﬂectivity and transmittance of the dichroic beam splitter. The blue andgreen curves are the reﬂectivity in the SW channel side. The two curves deviate around 160 and
180 µm.

ASTRO-F Observer’s Manual

40

4.2.2

Detector system

Dead pixels
The SW- and LW-arrays are known to have 4 and 7 dead pixels, respectively. In addition one
SW pixel has a large noise level and an unstable signal output. Figure 4.2.12 shows the positionsof the dead and badly performing pixels.

Figure 4.2.12: The positions of dead (blue cross) and badly performing (green cross) pixels in
the FIS detector arrays. Deﬁnitions of the pixel numbering indices are also presented.

Responsivity and Uniformity
Figure 4.2.13 shows the responsivity of all the SW detector pixels. The values are calculated fora νIν = constant spectrum at the central wavelength deﬁned in Table 4.1.1. Similar information
presented in a diﬀerent form is shown for the LW detectors (Figure 4.2.14). The variation (1σ)
is 25–30 per cent.

Figure 4.2.13: Responsivities of all pixels in the SW detector arrays. See ﬁgure 4.2.12 for thedeﬁnition of the pixel number.

Version 3.2 (November 29, 2005)

41

Figure 4.2.14: Pixel by pixel responsivity variation of the LW detector arrays.

Noise
After the results of the early FM performance tests, the FIS introduced two noise ﬁlters, theﬁrst in the cryogenic part near the detector and the other outside the cryostat. These ﬁlters
successfully reduced the noise signiﬁcantly.
There are noise components at speciﬁc frequencies in the SW detector (Figure 4.2.15). It
is suspected that the components are residuals of electrical interference from the other instruments that could not be removed by the ﬁlters. The frequency of the components move as the
instrument status changes. The power of these excess components is roughly one half of thetotal noise power. Attempts to reduce these noise components in the data reduction process arebeing considered.

Figure 4.2.15: Noise spectral proﬁle of a SW detector pixel. Time variation of the noise spectra
taken during the laboratory test under dark conditions. White stripes are excess noise components. We observe a number of components at speciﬁc frequencies. Some of them move with
time.

42

ASTRO-F Observer’s ManualRamp curve
Ideally the output signal of the FIS detector should be proportional to the number of incoming
photons, but this is not the case in reality. The actual ramp curves show deviations from theoptimal linear line even when the data is taken under constant incoming radiation. The SW
detector shows this eﬀect more prominently (Figure 4.2.16). As the slope of the ramp curve gives
the instantaneous ﬂux, any deviation from linearity means that the responsivity changes along
a ramp. Currently, the SW ramp curve becomes ﬂatter as an integration progresses. Therefore,
if one wishes to observe a faint target, it is more eﬃcient to insert a reset frequently and to use
the linear and steep part of the ramps (see Section 4.4.1). The non-linearity of the LW detector
is 5–10 per cent.
If the shape of the ramp curve is reproducible, and if the curve is monotonic with incoming ﬂux, it will be possible to correct this eﬀect. We have repeated many tests to check the
reproductivity of the ramp curves under various ﬂux /detector conditions. We conﬁrm that theramp shape is stable at least over a few days time scale if the detector driving parameters anddetector temperature remain unchanged. Ramps taken under such conditions can be corrected
almost perfectly. However, we also see a large discrepancy over several pixels or under strong
radiation. The current correction method results in a residual error of the order of 10 per centin ﬂux. More studies are needed.

Figure 4.2.16: Comparison of integration ramps under constant irradiation levels. The LW
detector (right) shows an almost linear response while the SW signal (left) deviates from linearity
as the integration progresses.

Saturation After Eﬀect
It is found that the saturation of the detector is followed by a decrease in the signal level(Figure 4.2.17). This eﬀect is more obvious in the LW channel. The eﬀect appears moresigniﬁcantly when the saturation is heavier. It seems that the eﬀect is already present when the
integrated signal has reached about half of the full saturation level. The recovery time scale isof the order of a minute. This eﬀect is still under quantitative analysis.
It is recommended to set the reset interval so as to not cause saturation of the detector,
although at the cost of sensitivity. The saturation level given in the following section approximately considers this eﬀect.
Transient response
It is known that the Ge:Ga type detectors show a strong transient response; the detector output
signal does not respond instantaneously to a change of incoming ﬂux but rather has a delay of
a few – hundreds seconds time scale.

Version 3.2 (November 29, 2005)

43

Figure 4.2.17: An example of Saturation after eﬀect for a LW detector pixel. The diﬀerential
signal for a constant irradiation is plotted against time after the saturation. In this ﬁgure an
attempt is made to ﬁt the data by a two component exponential function. More detailed analysis
is ongoing.
Figure 4.2.18 demonstrates how the transient response of the FIS detectors changes witha step change of incoming ﬂux. In order to simulate the real ﬂight conditions, a calibration
lamp is turned on weakly to provide a background level for the SW detector, and a Bias Light(Section 4.1.3) is used for the LW array. In both these cases, the time constants are 10–30 sec.
Since the ASTRO-F/FIS almost always observes the sky in scanning mode, an understandingof the eﬀect of the transient response to the observed data is essential. The most extreme caseis the All-Sky Survey, in which a point source is scanned in only 0.2–0.4 sec, much smaller
than the transient time constant, resulting in a lower output signal than expected by factor ofa few (Figure 4.2.19). This also aﬀects the detection limits. The eﬀect is not so severe for the
Slow-Scan mode for the pointed observations.
Ideally, the transient response should be understood and reproduced by modeling the chargetransfer in the detector element. However, strong non-linear physics prevent us from making aperfect model, although the analysis is continuing. For the All-Sky Survey data, the point source
ﬂux is planned to be corrected empirically by tables comparing the signals for the instantaneousand continuous ﬂux levels.

ASTRO-F Observer’s Manual

44

0.003

0.006
SW 17ch

LW 11ch0.005
differential output (V)differential output (V)0.00250.002

0.00150.001

0.0005
0

0.004
0.003
0.002
0.001

0

200

400
600
800
sampling @24Hz

1000

1200

0

0

100

200

300
400
500
sampling @16Hz

600

700

800

Figure 4.2.18: Examples of transient response for the FIS detectors for a step-wise change ofthe incident ﬂux.

Figure 4.2.19: Transient response to pulse irradiation. On the left side of the each ﬁgure thecalibration lamps are ﬂashed at very short periods simulating the passage of a point source in
the All-Sky Survey. While on the right, the lamps are turned on continuously for 5 sec. The
intensity of the calibration lamps was the same for both cases. Although 5 sec is not enoughtime for the detectors to reach a constant output level, it is obvious that the output signal for
ﬂashes is smaller than that of the 5 sec case. This ratio will be used for the ﬂux calibration ofpoint sources.

Version 3.2 (November 29, 2005)

45

Cross-talk
Cross-talk between the detector pixels will broaden the eﬀective FWHM of a point source image.
Moreover, it will increase the eﬀective number of glitches, and complicate the transient response.
This eﬀect may be more serious in the SW detector, which is a monolithic structure. The crosstalk eﬀect for point sources was measured by a pin-hole mask image, as explained in Section 4.2.1.
From this measurement, it appears that the non-stationary cross-talk is not signiﬁcant. The
measured FWHM of a point source image is consistent with a cross-talk level of about 20 per
cent. However, other explanations cannot yet be ruled out.
So far no clear evidence of cross-talk observed in the LW detectors (Section 4.2.1). Moreanalysis is required to a give clear answer to this phenomenon.
Spectral responsivity
The LW detectors of the FIS are of the stressed Ge:Ga type. The stronger the stress (pressure)on the detector element, the longer the wavelength limit can be extended. It is observed thatthe cut-oﬀ wavelength (the longest wavelength at which the detector is sensitive) varies fromone pixel to another, probably because of inhomogeneous stress among the pixels. Figure 4.2.20
demonstrates the pixel by pixel variation of the spectral responsivity of the LW channel. It isseen that the cut-oﬀ wavelength ranges from 165–185 µm. The N160 band extends longer thanthe WIDE-L band as the detectors are under even higher pressure.

Figure 4.2.20: Pixel by pixel variation of the spectral response of the LW detector.
The diﬀerence in wavelength proﬁle directly inﬂuences the photometric accuracy, especiallyfor cool targets such as molecular clouds. The error in photometry is less than a few per centfor sources warmer than 100 K, but is more than 10 per cent for a target of 30 K.

ASTRO-F Observer’s Manual

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4.2.3

FTS modeTechnically, the FTS mode of the FIS is very challenging. In addition to the various instrumental
behaviors noted earlier, there are more issues to be considered. Observers who are consideringobservations with the FTS mode should carefully read the following notes and investigate if theirobservations are feasible with the instrument. Please note that according to the performancemeasured during the PV phase, there remains the possibility that observations with this mode
may be canceled.
Transient Eﬀects on the FTS data
In the FTS mode, the detectors have to observe a rapidly changing signal from the interferogram.
Naturally, strong transient eﬀects appear in the output signal. Figure 4.2.21 plots the interferograms corresponding to the forward and backward motion of the mirror. Large deviations areseen.
For the moment we do not have perfect correction methods for the transients, and thecorrection of this eﬀect is rather empirical. It is known that the strength and time constantof the transient eﬀect is a function of the irradiating ﬂux level. Therefore, we ﬁrst observe theinternal calibration lamp whose SED is known. The brightness of the lamp is adjusted to thatof the target. The data from the calibration lamp is reduced without transient correction. By
comparing the results and the black body spectrum of the lamp in the calibration data, we canderive a correction factor between the observed and “true” spectrum. This correction factor isthen applied to the target spectrum.
0.1[V/sec]SW#23

forwardbackward

0.05
0
-0.05
-0.1
-0.15
-0.1

-0.05

0

0.05

0.1

Optical Path Difference [cm]0.1

[V/sec]0.05

forwardbackward

LW#24

0
-0.05
-0.1
-0.15
-0.2
-0.15

-0.1

-0.05

0

0.05

0.1

0.15

Optical Path Difference [cm]Figure 4.2.21: Comparison of forward (blue) and backward (red) interferograms. The diﬀerence
between the two interferograms is thought to be due to detector transients.

Version 3.2 (November 29, 2005)

47

Detector Response Inhomogeneity
As described above, pixels in the detector arrays do not have a uniform responsivity. Figure 4.2.22 indicates the relative responsivity variation of the detector pixels in FTS mode at
diﬀerent wavelengths. The majority of the pixels are within 20 per cent (SW) and 50 per cent
(LW) of each other. There are several pixels with extremely high responsivity; namely pixel
23 (and several nearby ones) of the SW and pixel 31 of the LW array. These high responsivity
pixels are used to estimate the initial center of the interferogram.
In addition to the responsivity variation, there is variation in the wavelength proﬁle (longwavelength cut-oﬀ) in the LW detectors.

Figure 4.2.22: The pixel by pixel variation of detector response in FTS mode. There are a few
pixels showing signiﬁcantly high values.

Spectrum Reproductivity
In order to evaluate the stability of the spectroscopic data, three spectra were taken during thelaboratory tests. The ﬁrst two were taken with a one hour interval, then the third one wasobtained after three hours. These were independently reduced and are overlaid in Figure 4.2.23.
Although the spectra of the SW channel show large amplitude fringes, we see that the three
spectra are almost equivalent to each other to within about 10 per cent.
Note that these spectra were taken one after the other in a continuous test session. It has
not yet been evaluated whether the same result holds true for the spectra taken between longintervals or after turbulence (e.g., SAA).
Fringes
The SW channel suﬀers from strong fringes in the spectra. Figure 4.2.24 shows an enlargementof the fringe patterns in the Full-Resolution mode and the SED mode respectively. The causeof the fringes is thought to be a reﬂection within the detector element, by the parallel surface ofthe chip (therefore it should also present in the photometric mode). Throughout the laboratory
measurements, the fringe pattern remains at the same position, implying that it can be easily
removed. However, it is not clear whether this will hold for the ﬂight data, especially for point
sources. Observations of sharp lines on the fringe shoulders will require special care.

ASTRO-F Observer’s Manual

48
0.001

0.00030.00080.000250.00060.000150.0002[V/sec/cm-1][V/sec/cm-1]

0.00020.00040.00010-0.00025 10 -5-0.0004
60

80

100

120

140

160

180

0
200

Wavenumber [cm -1]

Figure 4.2.23: Three FTS spectra taken one and three hours after the initial spectrum areoverlaid. They are almost equivalent. The bottom plot is the diﬀerence between the spectra. It
is seen that the diﬀerence looks random but is within about 10 per cent. The large deviations
beyond 180 cm−1 are due to excess noise.
0.0015

0.001

0.001
2.53 cm

0.001
2.53 cm-10.0008-10.00080.00080.00060.00060.00040.00040.00020.0002[V/sec/cm-1]

0.0004[V/sec/cm-1]

0.0006[V/sec/cm-1][V/sec/cm-1]

0.001

0.00050.0002

0
90

92

94

96

98

Wavenumber [cm -1]

0
100

0
110

112

114

116

118

0
120

Wavenumber [cm -1]

Figure 4.2.24: Fringe pattern of the FTS spectra in Full-Resolution mode (left) and SED mode(right). In Full-Resolution mode, the fringe pattern is resolved out and an “Airy pattern” ofthe Fabry-Perot interference is clearly seen.The blue lines in the bottom panel are the rms noiselevel. The three lines in the right panel are diﬀerent measurements, and match each otherindicating that the fringe pattern is well repeated.
Interference from the Cryocooler
It is known that the vibration from the cryocooler interferes the FTS mirror drive. The vibration
appears as a change of the mirror scan speed of up to 50 per cent, periodically with a frequencyof 15 Hz. As the detector readout is at a constant frequency, this means that the sampling of
the interferogram is at inhomogeneous intervals of the optical path. As the results the FFTtechnique cannot be applied. It may also have an eﬀect on the transient behaviour and change
the observed spectra. No clear inﬂuence on the resultant spectra has been observed at this stageof the data analysis.

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4.3

49

The FIS All-Sky Survey

The FIS is designed primarily to perform the All-Sky Survey in 4 photometric bands at wavelengths of 50–180 µm. The goal is to observe the entire sky on at least two independent orbits.
Although most of Phase 1 will be dedicated to observing time for the All-Sky Survey, supplemental survey operations will also be needed in Phase 2 to ﬁll in the gaps remaining from Phase
1. A dedicated team within the ASTRO-F Project will handle all responsibility for the operationand the data reduction of the FIS All-Sky Survey. Products will be distributed to the science
users. (see Section 6).

4.3.1

Operation

During the survey, the detectors are readout continuously with a constant sampling rate foreach array. At a ﬁxed interval (TBD, ∼ 2 seconds) the detectors are reset. Pulse ﬂashes of thecalibration lamps are inserted periodically every 1 minute while keeping the shutter open in orderto monitor the detector responsivity. About twice per orbit as well as before and after pointed
observations, the shutter is closed to measure the dark level, and the calibration lamp is turnedon continuously to monitor the detector responsivity in its steady state. The time constants of
the above sequence will be adjusted during the PV phase. The resets and calibration operations
will obviously cause gaps in the scanning data.

ASTRO-F Observer’s Manual

50

4.3.2

Expected Performance

Detection Limits
Table 4.3.3 tabulates the 5σ detection limits of the FIS survey for point sources and diﬀuse
radiation. The estimates are based on the simulation of Jeong et al. (2003) that assumed rather
ideal conditions.
The numbers are for a single scan with the detector array (not per pixel). Two readout
modes, namely Nominal and CDS modes (see Section 4.1.3) are considered. Note that theperformance of the CDS mode has not yet been fully conﬁrmed.
The estimates of sensitivity are based on the latest information of the detector performanceand the optical eﬃciency. Known detector characteristics, such as noise levels, responsivity,
dark current, transient eﬀects, spectral response, expected eﬀects of charged particle hits, etc.,
as well as the optical eﬃciency are taken into account. The numbers should be treated as thebest estimates of the FIS sensitivity, however, there may still be uncertainty of up to a factor
of a few. Time variation of the detector responsivity and noise are not accurately included.
Pixel by pixel variation of the sensitivity is as much as a factor of two. In particular, it isnoted that the long-wavelength cut-oﬀ of the LW detector is diﬀerent for diﬀerent pixels, due
to inhomogeneous stress on the array. This means that the pixel by pixel relative responsivity
changes as a function of the SED of a target.
The LW detector has a very long time constant for the transient response, which results ina lower sensitivity in the survey operation (see Section 4.2.2). The eﬀect of the bias light has
already been taken into account in the table, but not fully conﬁrmed.
It is possible that the high redundancy of data taken at high ecliptic latitude regions will
improve the detection limits if the data is suitably processed. Such a process will be applied forthe generation of the Faint Source Catalogue. The Bright Source Catalogue is constructed fromindividual scans and the sensitivities in this table should be adopted.
Table 4.3.3: The FIS All-Sky Survey 5σ detection limits per single scan.
N60
WIDE-S
WIDE-L
N160

Point Source (mJy)
1000
200
400
800

Diﬀuse Source (MJy/sr)
20
4
2
4

The actual detection limits at the longer wavelengths may be dominated by the sky confusion.
Sky confusion is not considered in the current sensitivity estimates. The observers are suggested
to refer to Appendix A for discussions of the the sky confusion. An independent estimate canbe found on the Observer Support page of the ISAS Web page (in Japanese).
Saturation Limits
Since the FIS detectors use charge integrating type circuits for data acquisition, the detector
signal will saturate when the accumulated charge between resets reaches a certain limit. Thismeans that the saturation does not take place spontaneously by a single bright source but
depends on the brightness distribution along the entire scan path between resets. Once thedetector saturates, all the information in that detector is lost until next reset. Impacts ofcharged particles could also cause saturation (Section 4.4.6).
Table 4.3.4 gives the saturation level of the FIS All-Sky Survey. The point source saturation
limit is deﬁned as where the ramp curve reaches a level where the reproductivity of the ramp

Version 3.2 (November 29, 2005)

51

curve shape becomes bad (e.g., more than 10 per cent) by a point source passage. The diﬀuse
saturation level is determined as where the ramp reaches the same level in a 2 sec integration.
Table 4.3.4: The FIS All-Sky Survey saturation limits.
Band
N60
WIDE-S
WIDE-L
N160

Point source (Jy)
500
100
100
300

Diﬀuse Source (GJy/sr)
1.0
0.2
0.1
0.5

Our simulation (Jeong et al. internal report) concludes that with a 2 sec reset interval, about
2 per cent of the sky will be lost to saturation. These areas will be mostly on the inner galactic
plane. Use of a shorter reset interval or the CDS mode is under consideration for these areas.

ASTRO-F Observer’s Manual

52

4.4

The FIS Pointed Observations: AOTs

4.4.1

FIS01: Photometry / Mapping of Small Area

The FIS01 AOT is designed for photometry of point sources and/or mapping of small areas up
to ∼ 25 × 10 arcmin2 .
Table 4.4.5: FIS01 summary.
Observing Mode
Band
Scan pattern
Sampling mode
Reset interval
Slow-Scan speed
Shift length

Fixed parameters
Photometry
N60, WIDE-S, WIDE-L, N160
Two round-trip scans with a cross-scan shift
User deﬁned parameters
Nominal or CDS
0.5 or 1 or 2 sec (for Nominal sampling mode)
8 or 15 arcsec/sec
70 or 240 arcsec[Instrument Operation]
The instrument is operated in the photometry/imaging mode; all four FIS photometric bands(N60, WIDE-S, WIDE-L, N160) retrieve data. The detector operation parameters are ﬁxed
at the recommended values provided by the instrument team. Users can specify only two
parameters; the readout mode (Nominal / CDS) and if the Nominal mode is selected, thereset interval.

[Scan Operation]
An observation consists of two sets of round-trip scans in the in-scan direction with a shift in
the cross-scan direction. The round trip scan is mandatory to ensure data redundancy. The
cross-scan shift is used to increase the redundancy or to observe a wider area of the sky. The
scan speed is either 8 or 15 arcsec/sec. These speeds are selected to maximize the performance(sensitivity) while enabling observations of a certain size area of sky.

[User speciﬁed parameters]
The following inputs can be given by the observer.
Detector readout mode (Nominal / CDS) CDS mode should only be used for the brightest targets that would normally saturate the detector in Nominal mode. See Section 4.1.3
for the details of the readout methods, and Sections 4.4.2 for the sensitivity and saturationlevels.
Reset interval This parameter is only relevant for the Nominal readout mode and is chosen
based on the brightness of the targets. Generally, a long reset interval improves sensitivity
in the ampliﬁer but at the same time increases the risk of saturation (please be remindedthat once a detector is saturated all information until the next reset is lost, and thatthe saturation can also be caused by charged particles). On the other hand, a shorter
reset looses more data points at the resets. For the SW detectors, intervals longer than
two seconds are not recommended even for faint sources since the responsivity drops due

Version 3.2 (November 29, 2005)

53

to nonlinear eﬀects. Here we present three choices for the reset interval in Table 4.4.5.
Shorter (0.25 sec) and longer (4 or 8 sec, only for the LW channel) reset intervals may be
added after performance evaluation during the PV phase.
By default, the SW and LW detectors (i.e., all four FIS bands) are always reset simultaneously. As the SW detector has a tighter constraint on the reset interval because of its
non-linearity, we recommend the following procedure to decide the reset interval.
1. In the case of a faint target where sensitivity is important, look at Table 4.4.9 andpick the reset interval such that the target ﬂux does not exceed the limit. As thesame reset interval has to be applied to both the WIDE-S and N60, the numbers for
the WIDE-S band, are usually suﬃcient.
2. If the target is bright and sensitivity is not essential, consult Table 4.4.8 and choose
the reset interval such that the target ﬂux does not exceed the saturation limits forany of the four FIS bands.
3. For very bright sources consider CDS mode. However, this readout mode is not yet
fully veriﬁed, and may have a large ﬂux uncertainty. A decision will be made in the
PV phase.
Scan parameters Two numbers have to be given to deﬁne the scan pattern; A scan-speed anda cross-scan shift length. A scan-speed of 8 arcsec/sec balances depth and sensitivity withscan area for data redundancy and is recommended for photometry of small numbers ofsources. A higher scan speed of 15 arcsec/sec is used to observe larger areas to a slightly
shallower sensitivity.
A cross-scan shift length of 70 arcsec is prepared for point source photometry, and 240
arcsec for small area mapping. In the latter case, it is expected that the observation will
be repeated on the successive orbit to extend the mapping area, with enough overlap toensure redundancy.
A round trip scan takes approximately tscan = (l/s + tacc ) × 2 seconds where l is the scan
length in arcsec, s is the scan speed and tacc is the time needed for scan acceleration. tacc
is roughly 30 sec (8 arcsec/sec) and 40 sec (15 arcsec/sec). Deadtime for a cross-scan
shift is roughly estimated as 35 sec for a 70 arcsec shift and 45 sec for 240 arcsec shift(see Section 3.2.2 for details). As the total observing time for a pointing opportunity islimited to 10 minutes, the scan length is uniquely determined once the scan speed and shift
length are set. Table 4.4.6 shows approximate scan lengths and corresponding observationeﬃciencies for combinations of scan speed and cross-scan shift lengths. These values will
be updated when more accurate numbers for the dead time become available. Data arealso retrieved during the acceleration. These data may be usable but it has yet to be
conﬁrmed.
Target position The satellite is operated such that the target position is located at the centerof the scanned area.

ASTRO-F Observer’s Manual

54

2

+

+

3

1

+

+

+

+

Lscan

+

Licomplete+

Lccomplete
Lstep
Figure 4.4.25: The scan sequence of the AOT FIS01. (1) First round-trip scan, (2) Step-Scan,(3) Second round-trip scan. The area indicated by the green hatch is scanned four times (tworound-trips) by all four detector arrays with more than two pixels of each array. The area inside
the orange line is scanned by four bands at least twice (one round-trip). The dimensions in theﬁgure are given in table 4.4.6 for each scan parameter set.

Table 4.4.6: Approximate estimates of scan lengths for diﬀerent scan parameter in FIS01.
(1)
Scan speed(arcsec/sec)
8
8
15
15

(2)
Lstep
(arcsec)
70
240
70
240

(3)
Lscan
(arcsec)
890
870
1520
1480

(4)
Licomplete(arcsec)
580
560
1210
1170

(5)
Lccomplete(arcsec)
420
240
420
240

(6)
Eﬃciency
0.74
0.73
0.68
0.66

(1)(2) Scan parameters.
(3) Scan length with a steady speed.
(4) Scan length observed by all four detector arrays with more than two pixels per band.
(5) Cross-scan length observed by both ﬁrst and second round-trip scans.
(6) Fraction of time used for taking usable data to the total observing time (600 sec).

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4.4.2

55

Expected Performance of FIS01

Detection Limits
As we explained in Section 4.4, photometry (imaging) observations by the FIS in pointed observation mode is always carried out with the Slow-Scan operation. Therefore the estimates inthis section are given by scaling the All-Sky Survey detection limits by some factors which arefunctions of the scan speed. However, because of the nonlinear responsivity, transient eﬀects,
and other complicated characteristics of the instruments, the scaling law is not always straightforward. Table 4.4.7 lists the sensitivities for the two scan speeds, 8 and 15 arcsec/sec, used for
the AOTs. Note that confusion is not included in these estimates.
The numbers in the table are for a single scan. Therefore one can imagine that four scansin FIS01 should improve the detection limit by factor of 2. On the other hand, this redundancy
can alternatively be used to increase the reliability of the results. In such a case, the detection
limit remain unchanged. We leave this at the discretion of the user. As the “standard” rule ofthumb, we propose to apply the numbers√ per one round-trip scan. So for FIS01, which carries
out two round trips, a sensitivity of 1/ 2 of the value in the table can be assumed.
Table 4.4.7: The 5σ detection limits for AOT FIS01. Numbers are per scan. See text for details.
Scan speed
Reset
N60
WIDE-S
WIDE-LN160

0.5 sec
90
18
24
48

Scan speed
Reset
N60
WIDE-S
WIDE-LN160

0.5 sec
2
0.4
0.18
0.36

Point Source (mJy)
8 arcsec/sec
1.0 sec 2.0 sec CDS 0.5 sec
60
45
1600
140
12
9
320
28
12
6
460
36
24
12
920
72
Diﬀuse Source (MJy/sr)
8 arcsec/sec
1.0 sec 2.0 sec CDS 0.5 sec
1.5
1
40
3
0.3
0.2
8
0.6
0.09
0.04
2.5
0.24
0.18
0.08
5.0
0.48

15 arcsec/sec
1.0 sec 2.0 sec
100
70
20
14
18
9
36
18

CDS
2500
500
700
1400

15 arcsec/sec
1.0 sec 2.0 sec
2
1.5
0.4
0.3
0.12
0.06
0.24
0.12

CDS
50
10
3.5
7

Saturation Limits
Table 4.4.8 gives the saturation levels for the AOTs FIS01 and FIS02. The saturation limits
are functions of the reset interval only. These limits are deﬁned such that the integration rampsof the high sensitivity pixels reach a level where the reproductivity of the ramp shape becomesbad. Beyond this level we still obtain data from the lower-sensitivity pixels but photometricaccuracy and quality of reconstructed image map may be signiﬁcantly degraded.
Linearity Limits
The linearity limit is only relevant for the SW detector which shows a large non-linearity in the
integration ramp. The limits are deﬁned such that the ramp reaches to the level where the slope(i.e. sensitivity) becomes approximately half of the value at the beginning of the ramp. It is
recommended to carry out any observations below this limit for eﬃcient measurement of faintsources.

ASTRO-F Observer’s Manual

56

Table 4.4.8: The saturation limits for FIS01 and FIS02. See text for details.
Point Source (Jy)Reset (sec) 0.5 1.0
2.0
N60
280 14070
WIDE-S
60
30
15
WIDE-L120 5015N160
300 130
60
Diﬀuse Source (GJy/sr)
Reset (sec) 0.5 1.0
2.0
N60
4.0 2.0
1.0
WIDE-S
0.8 0.4
0.2
WIDE-L1.0 0.42 0.12N160
2.5 1.10 0.5

CDS
3600
700
1500
3500
CDS
50
10
13
30

Table 4.4.9: The linearity limits for the SW detectors for FIS01 and FIS02.
Point Source (Jy)Reset (sec) 0.5 1.0 2.0
N60
60 30
15
WIDE-S13
6
3
Diﬀuse Source (GJy/sr)
Reset (sec) 0.5 1.0 2.0
N60
0.9 0.4 0.2
WIDE-S
0.2 0.1 0.05

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4.4.3

57

FIS02: Mapping of large area

The FIS02 AOT is designed for mapping of relatively large areas. An observation covers a long
strip of ∼ 1 degree.
Table 4.4.10: FIS02 summary.
Fixed parameters
Observing Mode PhotometryBand
N60, WIDE-S, WIDE-L, N160
Slow-Scan speed 15 or 30 arcsec / sec
Scan pattern
One round trip scan
User deﬁned parameters
Sampling mode
Nominal∗ or CDS
Reset interval
0.5 or 1.0 or 2.0 sec
Slow-Scan speed 15 or 30 arcsec/sec∗ = Co-add mode

[Instrument operation]
The operation of the instrument with this AOT is the same as FIS01, except that FIS02 carries
out a single round-trip only, and that the choice of scan speed is 15 or 30 arcsec/sec. 3[Scan operation]
The diﬀerence of this AOT from FIS01 is that an observation consists of only a single roundtrip scan. The scan speed is chosen from 15 or 30 arcsec/sec. Unlike FIS01, no cross-scan
shift is operated during a pointed observation in order to maximize the scan length. It is
expected that an adjacent area of the sky will be observed on a successive (or later) orbit.
It is strongly recommended to observe the same area of the sky at least twice with diﬀerent
pointing opportunities (i.e. total two round trips at any point) to acquire enough redundancy.
It is the observer’s responsibility to allocate the scans to cover the required area and/or to
obtain the redundancy. With the currently assumed deadtime for acceleration, the scan lengthis approximately 1 and 2 deg. depending on the scan speed.

[User speciﬁed parameters]
The following inputs have to be given by the observer.
Detector readout mode (Nominal / CDS) Same as FIS01.
Reset interval Same as FIS01. See section 4.4.2 and 4.4.2 for the sensitivity and saturationlevels.
Target position The satellite is operated such that the target position is located at the beginning of the scan strip.

3

30 arcsec/sec is a new addition for Open Time observations

ASTRO-F Observer’s Manual

58

4.4.4

Expected Performance of FIS02

Detection Limits
Table 4.4.11 lists the sensitivity for the two scan speeds, 15 and 30 arcsec/sec, used for the
AOTs. Note that confusion is not considered in these estimates.
The detection limits in the table are per single scan. Read also the notes in Section 4.4.2 forestimating the detection limits for multiple observations.
Table 4.4.11: The 5σ detection limits for AOT FIS02.
Point Source (mJy)
15 arcsec/sec
30 arcsec/sec
0.5 sec 1.0 sec 2.0 sec CDS 0.5 sec 1.0 sec 2.0 sec
140
100
70
2500
220
150
150
28
20
14
500
45
30
30
36
18
9
700
70
35
20
72
36
18
1400
140
70
40
Diﬀuse Source (MJy/sr)
Scan speed
15 arcsec/sec
30 arcsec/sec
Reset
0.5 sec 1.0 sec 2.0 sec CDS 0.5 sec 1.0 sec 2.0 sec
N60
3
2
1.5
50
4
3
3
WIDE-S
0.6
0.4
0.3
10
1
0.6
0.6
WIDE-L
0.24
0.12
0.06
3.5
0.35
0.17
0.10
N160
0.48
0.24
0.12
7
0.70
0.35
0.20
* Values for 15 arcsec/sec are identical to those in for the FIS01.
Scan speed
Reset
N60
WIDE-S
WIDE-L
N160

CDS
4000
800
1300
2600

CDS
70
15
7
14

Saturation Limits
Since the FIS detector saturation depends only on the reset interval, this is the same table as
for FIS01 (Table 4.4.8 on page 56).
Linearity Limits
Since the FIS detector saturation depends only on the reset interval, this is identical to FIS01(Table 4.4.9 on page 56).

Version 3.2 (November 29, 2005)

4.4.5

59

FIS03: FTS Spectroscopy

Observations with the FIS/FTS are very challenging. Due to the low sensitivity
and narrow dynamic range, careful selection of the target is needed. There arevarious instrument anomalies that have to be settled and corrected. Reduction of
the FTS data will consist of many iterative procedures and trial processing. Any
observers who plan to observe in this mode are kindly requested to carefully readthe following section as well as the performance section (Section 4.2) and to optimize
their observation plan to maximize the scientiﬁc output.
AOT description
The FIS03 AOT is designed for spectroscopic observations with the FTS.
Table 4.4.12: FIS03 summary.
Fixed parameters
Observing Mode Spectroscopy
Band
WIDE-S, WIDE-L
Scan pattern
Staring (no Step-Scan)
User deﬁned parameters
Sampling mode
Nominal (ﬁxed)
Reset interval
0.1, 0.25, 0.5, or 1.0 sec (see text)
Resolution
Full-resolution / SED[Instrument operation]
The instruments are set up for the FTS mode. The ﬁlter wheel is rotated to switch from theempty hole to the polarizer. Detector readout mode is switched to high-frequency sampling
mode. Only data from the WIDE detector arrays are acquired. During the observation, themoving mirror is driven. Two path lengths, i.e., spectral resolutions, are available (See Section 4.1.4). The spectral resolution is ∼ 0.36 cm−1 for the Full resolution mode and ∼ 2.4 cm−1for the SED mode. The SED mode provides better quality spectra than the Full-resolution mode
spectra after being smoothed to a corresponding resolution. There are ﬁve choices for the reset
interval, similar to the other FIS AOT’s. However, CDS readout mode is not available.

[Scan operation]
The FTS spectroscopic observation is operated in a staring pointing mode. To minimize the deadtime, No Step-Scan, Micro-Scan, nor Slow-Scan is carried out during a pointed observation. If
the observer needs accurate spatial-spectroscopic information, one must repeat the observationof the same area with a slightly shifted target position. Observations of a wider area should
repeat the observations on diﬀerent orbits.

[User speciﬁed parameters]
The following inputs have to be given by the observer.
Resolution mode Full-resolution or SED. See the description above.

ASTRO-F Observer’s Manual

60

Reset interval Four choices from 0.14 , 0.25, 0.5, and 1.0 sec are available. As the CDS mode forthe brightest targets cannot be used for the FIS03 AOT, a shorter reset interval of 0.1 sec
is alternatively provided. This shortest reset interval is available thanks to the 5–7 times
faster detector sampling rate in this operation mode. As with the other observing modes,
the longest reset interval that does not cause saturation of the detectors is preferable. If
the target ﬂux is uncertain, a shorter reset interval is safer.
Target position The target position is observed at the center of the WIDE-S array. A small
adjustment may be applied according to the in-ﬂight performance.
Expected Performance[Detection Limits]
The 5σ detection limits for the FTS observations in one pointing opportunity are shown in
Table 4.4.13. In fact, the sensitivity of the FTS has a wavelength dependency. The sensitivity
is better near the band center and drops toward the band edges, following the spectral responseof the system. The SW channel shows a further decrease in sensitivity below 80 µm due to adecrease in the interferometric eﬃciency. In addition, interference occurs within the detector
elements of the SW array (Section 4.2.3. This results in strong fringes in the spectra. Due tothese fringes, the sensitivity is a rapid function of wavelength. More detailed information will
be given later.
It should be noted that the sensitivity is signiﬁcantly lower than the numbers previously
announced (before Observer’s Manual ver. 2.0).
Table 4.4.13 summarizes the detection limit of the FIS03 mode for various wavelengths, andTable 4.4.14 shows the detection limits for major FIS observable far-infrared lines. Figure 4.4.26
shows the 5σ detection limits proﬁle for continuum and line observations.
Table 4.4.13: Detection limits of the FTS for the exposure of 500 sec (5σ).
Line(10−15 W/m2)
Wavelength
(µm)
60
90
110
110
140
170

4

Wavenumber(cm−1)
170
111
90
90
71
60

Detector

The actual reset interval is 0.14 sec.

SW
SW
SW
LW
LW
LW

Full
33
4.8
2.4
37
8.4
24

Continuum
Point (Jy) Diﬀuse (MJy/sr)
Resolution
Full SED
Full
SED
710 100 15000
2130
100
14
2200
320
63
9
1400
200
510
73
7300
1040
160
23
2000
280
440
63
4700
670

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Table 4.4.14: Detection limits of the FTS for major far-infrared lines (500 sec, 5σ).
Line

Detector[N III][O III][O III][O I][O I][N II][C II]

SW
SW
SW
SW
LW
LW
LW

Wavelength
(µm)
57.32
51.83
88.36
63.18
145.53121.80157.74

Wavenumber(cm−1)
174.47
192.99113.18158.27
68.72
82.10
63.95

Detection Limit(10−15 W m−2 )
580
400
19
40
11
12
15

Wavelength [µm]
LW

200

150

120

100

80

LW

[O III][O I][O III]

10-14

SED : 2.4cm -1
10-15

10

WIDE-S band

WIDE-L band

SW : 2.1x10-8 * PSF
LW : 5.9x10-8 * PSF[N II][C II]

100

10-13

2

100

Full-resolution : 0.36cm -1

[O I]

1000

50
10-12

SW

1000

Flux [Jy]

10000

60

10000

[N III]SW

150Line Flux [W/m ]Surface Brightness [MJy/sr]

100000

75

100

125

150

175

200

-1

Wavenumber [cm ]Figure 4.4.26: The 5σ detection limits proﬁle of AOT FIS03 mode. The axes are for diﬀuse
radiation (LW and SW detectors respectively), point source continuum ﬂux, and line emission,respectively from left to right.

ASTRO-F Observer’s Manual

62[Saturation Limits]

Table 4.4.16 shows the saturation level in the FTS mode. These estimates are based on thephotometric mode data scaled by factors of instrumental eﬃciency. The polarizer used in the
FTS system reduces the optical throughput by a factor of 4. Higher detector sampling rates
also help to push the saturation level up higher. The numbers still have large uncertainties andshould be treated as “of order” estimates.
Table 4.4.15: This table is not available any more. Caption remains to keep the table numbers.

Table 4.4.16: Saturation limits for FIS03.
Point Source (Jy)Reset (sec)
0.1 0.25 0.5
SW(WIDE-S) 1100 520 260
LW(WIDE-L) 2000 960 400
Diﬀuse Source (GJy/sr)
Reset (sec)
0.1 0.25 0.5
SW(WIDE-S)
13
6.7 3.4
LW(WIDE-L)
17
8.0 3.4

1.0
130
170
1.0
1.7
1.4

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63

Remarks and Constraints of the FIS Observations

1. Hits of charged particles
Impacts of charged particles on the detector arrays produce glitches or spikes. A particle hit
could saturate the detectors. In addition “tails” with decay time-scales of a few tens–few
thousands seconds are observed after glitches. They are regarded as a temporal variationin the detector responsivity. The frequency of impacts heavy enough to be recognized as
glitches are estimated to be 0.1–a few per minute per pixel.
2. The South Atlantic Anomaly (SAA)The SAA is a region where the radiation belt surrounding the Earth reaches to a lower
altitude than at other places. In the SAA, a signiﬁcant number of charged particles hitting
the detectors prevent ASTRO-F from any observation. The responsivity of the detectors
also change after the passage of the SAA, with a recovery time-scale of a few thousand
seconds. About a half of the orbits per day are aﬀected by the SAA. We will carry out
a Bias boost (increasing detector bias to wash out charge carriers generated in the SAA)
for curing the detectors. This reduces the temporal variation of the detector responsivity,
but we still have a signiﬁcant time variation. We will attempt to correct these eﬀects bymonitoring the detector response by calibration lamp pulses and in the data reduction
process.
3. Observation of extended sources
It is known that Ge:Ga detectors have various complicated characteristics such as transienteﬀects and after eﬀects of charged particle hits etc. These characteristics make the data
reduction and calibration diﬃcult. Although the data reduction of the pointed observations
are the users responsibility, close communication with the project team is inevitable. In
particular, observers who are interested in imaging observations should not be overly
optimistic. We are open to any observers wishing to participate in the data reduction
activity. We also plan to organize a user consortium for the FIS observations in particular.
4. Observation with FTS mode
Again, we would like to emphasize thats observation with the FTS (FIS03) is not easy.
Therefore following conditions are given for the Open Time Proposals:(a) The number of FTS observations in open time will be limited to no more than 5 percent of the total number of open time pointed observations.
(b) Proposers are strongly recommended to conﬁrm the feasibility of proposed observations with the condition that the signal-to-noise ratio be larger than 10 for any single
pointed observation.
(c) Successful proposers are required to work closely with the FIS/FTS team for the data
reduction. A dedicated contact point will be given.

64

ASTRO-F Observer’s ManualChapter 5

IRC: Infrared Camera
The Infrared Camera (IRC) on board ASTRO-F is designed to carry out deep imaging observations in pointed observing mode. It’s unique wide ﬁeld coverage of 10 × 10 arcmin2 is ideal
for survey-type observations. The IRC’s low-resolution spectroscopic capabilities in the imaging
ﬁeld are also well suited for multi-object spectroscopic surveys. A scanning operation has beendeveloped with the cameras enabling us to extend the ASTRO-F All-Sky Survey wavelength
coverage into the Mid-Infrared range.

65

ASTRO-F Observer’s Manual

66

5.1

Hardware Specification

5.1.1

Overview

The IRC consists of three cameras: NIR, MIR-S and MIR-L. Covering the wavelength ranges
of 2–5, 5–13, 12–26 µm, respectively. In Table 5.1.1, the speciﬁcations of each camera aretabulated. See also ﬁgure 2.4.5 and 2.4.6.
Table 5.1.1: Speciﬁcations of the Infrared Camera (IRC).
ChannelNIR
MIR-S
Detector
InSb(SBRC-189) Si:As (CRC-744)
Array
512 × 412
256 × 256
2
1
Imaging FoV (arcmin )9.5 × 10.0
9.1 × 10.0
Imaging Area (pixel2 )1
391 × 412
233 × 256
Pixel Size (arcsec)
1.46 × 1.46
2.34 × 2.34
Wavelength (µm)
1.7–5.5
5.8–14.1
Filters
N2, N3, N4
S7, S9W, S11
Dispersion Elements
NP, NG
SG1, SG2
1 Cross-scan × in-scan. Masked areas are excluded.

MIR-L
Si:As (CRC-744)
256 × 25610.3 × 10.2
246 × 2562.51 × 2.39
12.4–26.5
L15, L18W2 , L24
LG2

Each camera is equipped with three ﬁlters and two dispersion elements. The ﬁlter selection
is preset in the Astronomical Observation Templates (AOTs; Section 5.5) and cannot be freely
chosen by the observers.

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67

OpticsAll three cameras of the IRC are oﬀ-axis, refractive optical systems. Each camera uses 2–5
lenses. A Ge beam-splitter divides the light between the NIR and MIR-S cameras. Transmitted
light comes into the NIR channel. Chromatic aberration does exist in the cameras and is largest
in the MIR-S camera which uses only two lenses.
A ﬁlter wheel is placed at the iris of each camera. Three ﬁlters, two dispersion elements, and ablind mask are prepared for each camera. The blind position will be used for dark measurements
during the ﬂight operation.

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68

5.1.3

Filters and Dispersion Elements

Table 5.1.2 shows the ﬁlters and the dispersion elements of the IRC.
Each camera is equipped with three broad band ﬁlters and two dispersion elements. The
NIR camera covers three independent wavelength bands that very roughly correspond to thewell known K, L, and M bands. The two MIR cameras each have two narrow ﬁlters which cover
the shorter and longer half of the wavelength range of the cameras,and a wide ﬁlter that overlaps
the two narrow ﬁlters.
The two dispersion elements of the NIR camera provide diﬀerent spectral resolutions over asimilar wavelength range. In the MIR-S and MIR-L, each dispersion element covers about halfof the camera’s wavelength range. However, unfortunately one of the elements of MIR-L, LG1,
was degradated during the ground tests. Thus only LG2 will be used for observations, resulting
in a gap in the wavelength range corresponding to the LG1 element.
Table 5.1.2: IRC Filters and Dispersion Elements(1)
Channel

(2)
Name

(3)

(4)
Wavelength(µm)
1.7–2.7
2.7–3.7
3.7–5.5
1.7–5.5
2.5–5.0
5.8–8.4
6.5–11.6
8.6–14.1
5.5–8.3
7.4–13.0
12.4–19.4
13.9–25.3
20.4–26.5

(5)
(6)
(7)
Center Width
Resolution
(µm)
(µm)
(λ/∆λ)
N2
ﬁlter
2.43
0.68
—
N3
ﬁlter
3.16
1.12
—
NIR
N4
ﬁlter
4.14
1.22
—
NP
prism
—
—
22 @3.5 µm
NG
grism
—
—
135 @3.6 µm
S7
ﬁlter
7.3
2.6
—
S9W
ﬁlter
9.1
4.3
—
MIR-S
S11
ﬁlter
10.7
4.7
—
SG1
grism
—
—
47 @6.6 µm
SG2
grism
—
—
34 @10.6 µm
L15
ﬁlter
15.7
6.20
—
L18W1 ﬁlter
18.3
10.0
—
MIR-L
L24
ﬁlter
23.0
5.4
—
LG12
grism
— Not available —
LG2
grism
17.7–25.0——
27 @20.2 µm(4) Deﬁned as where the responsivity for a given energy is larger than 1/e of the peak.
(5) Central wavelength of the ﬁlter band.
(6) Eﬀective bandwidth.
(7) From Table 5.5.10. Dispersion power of prism depends on wavelength. This number
is an average of the values at both ends of the wavelength range (4).
1 Renamed from L20W. No change of the wavelength proﬁle itself.
2 LG1 was degraded during the ground test and will not be used for observations.

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69

Field of View (FoV)

Figure 5.1.1 shows the location of the FoV of the three IRC cameras on the Focal-Plane projected
onto the sky (See also Figure 2.4.8). As noted earlier, the NIR and MIR-S cameras observe thesame FoV on the sky simultaneously. The MIR-L camera looks to a position oﬀset by ∼ 20
arcmin from the other two cameras. This means that at least two pointed observations on
diﬀerent revolutions are needed to observe a particular position by all three cameras. Step-Scan
operation (Section 3.2.2) will not be used because of its large dead time.
Reference positions for the AOT conﬁgurations for pointed observations are indicated in bluein Figure 5.1.1. Two reference positions for imaging observations (AOT IRC00, IRC02, IRC03,
and IRC11; see Section 5.5), namely the centers of the NIR/MIR-S and MIR-L cameras, arelabelled as ‘N’ and ‘L’, respectively. Reference positions for spectroscopic observations (AOT
IRC04) are explained in the following section.
The numbers in Table 5.1.1 and Figure 5.1.1 were derived from the data after the ﬁnal
integration of the instrument into the FM cryostat and can be regarded as the ﬁnal pre-ﬂight
values. An uncontrollable alignment error shifted the NIR and MIR-S FoV slightly toward theslits, resulting in a smaller imaging area than the original design. On the other hand, the FoVof the MIR-L has been enlarged by the successful refurbishment of the camera optics after 2003.
Slit (MIR-L): 7"x0.4' (Ls)

Scan Direction

10.7'

5.0'

MIR-L
4.9'

(L/Lc)

10.2'

10.0'
0.4'

NIR
& MIR-S(N/Nc)

Telescope Axis
0.9'
5.0'
Slit (NIR): 3"x1 (Nh)Slit (NIR): 1'x1' (Np)

25.0'9.3'(NIR)
9.5'(MIR-S)

Slit (MIR-S): 5"x0.4'
Slit (NIR): 5"x0.8' (Ns)0.7'(NIR), 0.5'(MIR-S)

Figure 5.1.1: Focal-Plane layout of the IRC cameras. Reference positions for pointed observations are indicated in blue.

5.1.5

Slits for Spectroscopy

The dispersion elements of the IRC are set into the ﬁlter wheel so that all the light in the FoV isdispersed. A spectrum is obtained in the direction parallel to the scan path (in-scan direction),
which runs along the vertical direction in Figures 2.4.8 and 5.1.1. Slits are provided for eachcamera (Figure 5.1.1) in order to avoid contamination by nearby sources / diﬀuse radiation. The
slits are primarily designed for extended sources and it should not be assumed that they can beused to guide a point source into the slit, except for the NIR camera which has an aperture forpoint sources. The current design is the following.

ASTRO-F Observer’s Manual70• The slit for the NIR camera consists of three parts of diﬀerent widths. The left most(closest to the imaging area) has a 5 arcsec width and will be mainly used for simultaneous
observations of diﬀuse light with the MIR-S camera. This slit position is labeled as ‘Ns’ for
IRC04 AOT observation parameter. Both the NP (low-resolution prism) and NG (highresolution grism) will be used with this slit. The middle 1 × 1 square part, referred to as‘Np’ is for spectroscopy of point sources. The aperture is large compared to the absolute
pointing accuracy of the satellite (designed to be better than 30 arcsec) to ensure that thetarget can be accurately guided into the area. Note that for observations of faint sources,
confusion due to galaxies may be a serious problem. The NG (grism) is assumed to be
used with this aperture. The rightmost (outer) part (‘Nh’) has a 3 arcsec width and isused for the highest resolution spectroscopy of diﬀuse radiation with the NG (grism).
• The MIR-S has a slit of 5 arcsec width for diﬀuse light. As this slit overlaps with innermost
slit for the NIR camera, it is also referred to as ‘Ns’. We assume that the point source
density in the mid-infrared range is low enough to avoid serious confusion, such thatspectroscopy of point sources can be made in the imaging ﬁeld.
• The MIR-L has a 7 arcsec slit for diﬀuse light similar to that of the MIR-S camera. This
position is referred to as ‘Ls’.

5.1.6

Detectors

The detectors of the IRC are provided by Raytheon IRCoE. The NIR camera has an InSb
detector array of 512 × 412 pixels (SBRC-189), in which 391 × 412 pixels are used for imaging
and the other 121 × 412 pixels are for slit spectroscopy. The two detectors for the MIR cameras
are both Si:As type (CRC-744) with a format of 256 × 256 pixels. There are some, but few dead/ hot pixels in the detectors.
It is known that the detector performance sensitivity changes with temperature. All detector arrays are equipped with thermometers to monitor the module temperature during theiroperation.

Version 3.2 (November 29, 2005)

71

5.2

Pre-flight performance

5.2.1

OpticsField of View (FoV)
As already explained in the previous section, The FoV of the IRC cameras diﬀer slightly from theoriginal design speciﬁcation values. The MIR-L has wider FoV of 10.3 (cross-scan) × 10.2 (inscan) arcmin2 . The FoV of the NIR is diﬀerent for diﬀerent ﬁlters depending on the alignmentof the ﬁlters. A summary is given in Table 5.2.3.
Table 5.2.3: Final FoV size of the NIR camera.
Band
N2
N3
N4

FoV (arcmin2 )9.2 × 10.0
9.5 × 10.0
9.5 × 10.0

Point Spread Function (PSF)
The expected performance of the ASTRO-F telescope was to be diﬀraction limited at ∼ 6 µm.
This means that the PSF is more extended than the diﬀraction limits in the NIR camera. The
PSF of each band was measured by putting an aperture mask at the focus of the cameras. The
PSF was measured by putting a pin-hole mask at the focus of each camera. The FWHM of themeasured PSF is larger than the designed value except for the L18W and L24 bands where theimages are diﬀraction limited. The PSF becomes elongated at the edges of the imaging area.
Table 5.2.4 shows the Full-Width Half Maximum (FWHM) of the PSF for the imaging bands ofthe IRC cameras. Note that because of diﬀerences between the laboratory measurements andﬂight conﬁguration, the real value may be slightly diﬀerent from those listed in the table.
Table 5.2.4: Full-Width Half Maximum (FWHM) of the PSF for the IRC imaging bands measured in the ground tests.
(1)
Channel
NIR
MIR-S
MIR-L

(2)
Band
N2
N3
N4
S7
S9W
S11
L15
L18W
L24

(3)
FWHM
(pix)
1.85
1.47
1.34
1.22
1.39
1.52
2.13
2.22
2.45

In Section 5.5.1 we discuss the optimized photometric aperture size and the detection limits.

ASTRO-F Observer’s Manual72
Chromatic Aberration

The NIR camera has chromatic aberration of 0.3–0.4 mm, which is in fact larger than thedesigned value. The most likely cause is the uncertainty of the optical characteristics of theelements under cryogenic temperature. The mean amount of aberration is comparable to thatof the diﬀraction limited optics at 2 µm and thus eﬀects the image quality.
Optical Transmittance
The throughout transmission function (RSRF: Relative Spectral Response Function) of the IRC
broad band ﬁlters are shown in Figures 5.2.2, 5.2.4, 5.2.6. The ﬁlters and optical components
as well as the detector response are taken into account to produce the curves1 . These RSRFs
can be directly integrated over spectra given as Fλ to obtain the in-band ﬂuxes.
Figures 5.2.3, 5.2.5, 5.2.7 show the RSRFs of the dispersion elements. In the ﬁgures they
are given as per photon.
This data, in digital format, are also available on the Observer’s support page.

1

These RSRFs have been prepared in collaboration with Dr. Martin Cohen to ensure a common calibrationof the IRC with other missions and ground-based instruments, including Subaru/COMICS, in the framework of
the absolute calibration network provided by him and his colleagues.

Version 3.2 (November 29, 2005)

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Figure 5.2.2: The Relative Spectral Response Function of the IRC/NIR Camera for Fλ .

Figure 5.2.3: The Relative Spectral Response Function of the IRC/NIR dispersion elements per
photon.

74

ASTRO-F Observer’s Manual

Figure 5.2.4: The Relative Spectral Response Function for Fλ of the IRC/MIR-S Camera.

Figure 5.2.5: The Relative Spectral Response Function of the IRC/MIR-S dispersion elements
per photon.

Version 3.2 (November 29, 2005)

75

Figure 5.2.6: The Relative Spectral Response Function of the IRC/MIR-L Camera for Fλ .

Figure 5.2.7: The Relative Spectral Response Function of the IRC/MIR-L dispersion element
per photon. LG1 will not be used for astronomical observations due to degradation. Data shown
here is 2nd order light subtracted.

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ASTRO-F Observer’s Manual

Leakage
Leakage has been investigated for the MIR ﬁlters. Figures 5.2.8 shows the ﬁlter transmittance
proﬁles of the ﬂight backup ﬁlters. All ﬁlters show an indication of short wavelength leakage
around 1/2 of the main passband. In addition, long wavelength leakage is also found in the S7and S11 ﬁlters. The contribution of these short wavelength leaks in the photometry for stars> 3000K is about 0.6 per cent in the S7 and S9W band, 1–2 per cent in the S11 and L24 bands,
and 3–4 per cent for the L15 and L18W bands.

Figure 5.2.8: Possible leaks in the IRC/MIR ﬁlters. Leaks are indicated by blue or red masks.

Distortion
Image distortion of the IRC FoV is illustrated in Figure 5.2.9 (Kim et al. 2005; Fujishiro etal. 2005) for three ﬁlter bands. They are calculated by optical simulation software with theparameters reﬂecting the results of the laboratory measurements.
Distortion is largest in the NIR camera with a shift of up to 10 pixels near the edge. It is
about 7 pixels in the MIR-S camera. The MIR-L camera shows little distortion, 1.5 pixels at
the most. The image distortion will be also measured during the PV phase and a correction will
be applied in the data reduction.
Ghosting
Ghosting is found in the NIR and MIR-S cameras. Figure 5.2.10 shows NIR and MIR-S imagesof a pin-hole mask at the aperture. The pin-holes are at regular grid intervals and are seen as
bright spots. The ghosting images of the holes appear between the bright spots.
The cause of the ghosting has been understood as reﬂection inside the beam-splitter. The
relative position between the true image and the ghost image is predictable. The energy ratioof the ghost image with respect to the true image is about 0.7 per cent for the NIR bands andVersion 3.2 (November 29, 2005)

77

Figure 5.2.9: Image distortion pattern of the IRC cameras. From left to right, NIR (N2),MIR-S (S7), and MIR-L(L18W). Distortion is emphasized by the arrow lengths. The maximum
distortion is about 10 pixels in NIR, 7 pixels in MIR-S, and 1.5 pixels in MIR-L, respectively(Kim et al. 2005; Fujishiro et al. 2005).
about 0.75 per cent for S7 and S9W bands. The S11 band shows the highest ratio of 3.8 per
cent, due to the characteristics of the A/R coating of the beam splitter.
Since the ghosting is predictable, it is, in principle, possible to remove it at the data reduction
stage. Nevertheless, observations of faint sources nearby strong sources will require special care.
No obvious ghosting was detected in the MIR-L camera.

Figure 5.2.10: Ghost image of the NIR (right) and MIR-S (left) cameras. These images are
obtained by putting a pin-hole mask at the entrance aperture. The bright spots are pin-holes,
and weak spots between the bright spots are the ghost images. The slits are on top of each
image. Extended features seen in the MIR-S data are due to the measurement conﬁguration
and will not appear in the ﬂight data.

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78

5.2.2

Detector system

Bad Pixels
There are some dead / bad pixels reported in each camera of the IRC. The numbers are muchless than 0.1 per cent in the NIR and MIR-S, and are about 0.3 per cent in the MIR-L camera
2 . The bad pixels are usually independent or clustered over a several pixels. Only the MIR-L
shows a relatively large cluster near the middle and at one of the corners (Figure 5.2.11). The
one in the corner is probably due to eﬀects from the detector readout circuit.

Figure 5.2.11: Bad pixels on the MIR-L detectors. The dimensions of the image correspondto the FoV geometry projected onto the sky. A large clustering at the bottom-right corner isprobably due to eﬀects from the detector readout circuit.
Dithering observations are essential to minimize the eﬀect of bad pixels and are included in
AOT IRC02 and 03 Section 5.5). It is the users’ responsibility to repeat observations in diﬀerent
pointing opportunities for IRC00 and 04 (and probably 03) observations.
NoiseTransient eﬀects of the detectors are expected to be very small, and may be negligible in mostcases. The stability of the responsivity after irradiation by charged particles (especially in SAA)
has not been evaluated quantitatively.
Linearity
The linearity of the IRC detectors has been measured. Plotted in Figure 5.2.12 are the deviations
from linearity of the detector output signal against the readout signal in ADU. It is seen thatthe three detectors follow a similar non-linear pattern and stay within 10 per cent deviation up
to ∼ 12000 ADU. This results in a dynamic range of over four orders in magnitude. Current
investigation shows that the non-linearity correction can improve the linearity to the few percent level.

2

Here we deﬁne bad pixels as those signals in which the dark image deviates by 3σ from the average.

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Figure 5.2.12: Deviation from linearity plotted against the readout signal level in ADU for thethree IRC detectors.

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80

5.3

The IRC Instrument Operation

5.3.1

Pointed Observations

Exposure cycle
The minimum observation unit of the IRC is called an exposure cycle. An IRC pointed observation consists of an n times repeated exposure cycle with operations such as a ﬁlter change andMicro-Scan inserted between them. The operation sequence of the exposure cycle is designedby the IRC team to optimize the performance of the instrument and observation eﬃciency. The
current (and expected to be the ﬂight) design of an exposure cycle is illustrated in Figure 5.3.13.
The unit time length for a detector driving is ∼ 0.58 seconds, which corresponds to the time
for the MIR-S/L detectors to access all pixels in the array. The NIR detector takes 4 times
longer. Time durations in the exposure cycle are always a multiple of this unit time.
One exposure cycle takes about 72.5 sec, for both the NIR and MIR channels. The threecameras are operated synchronously. However, use of this exposure cycle is diﬀerent in the NIR
and MIR channels. The NIR detector carries out one short (2.3 sec) and one long exposure (51sec) in a cycle, while the MIR detectors do one short (0.6 sec) and three long (19 sec) exposures.
These short exposures correspond to one detector driving unit time. The short exposures areuseful for bright stars that may saturate the detectors. However, the photometric accuracy ofthe short-exposure is probably much worse than the long-exposure. Fowler 4 sampling is taken
for the longer integrations of all detectors, to reduce the readout noise by a factor of two.
Two images from the NIR channel and four images from each MIR channel are produced in
one exposure cycle. They are stored in the frame memory buﬀer, then divided into telemetry
packets and down-linked to the ground. The on ground data reformatting software will reconstruct the frame memory buﬀer and eventually a set of images, which are then passed to theobservers.
Signal level
Frame1RST

Frame2
Fowler4

RST-INT

NIR~2.3[s]Time~51[s]~72[s]MIR-S/L

Fowler4

RST RST

Fowler4
RST

Fowler4
RST
Time

RST-INT~19[s]~0.6[s]Frame1

Frame2

Frame3

Frame4

Figure 5.3.13: An illustration of the operation of one exposure cycle. A pair of short and longexposures are carried out for the NIR camera while one short and three long exposures areprepared for the MIR channels. RST indicates a reset of the detector.

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A pointed observation is carried out by repeating the exposure cycle with Micro-Scans andﬁlter changes inserted between them. These operations cause dead time due to the operations
themselves and stabilization of the satellite attitude. Current estimates of the dead time are
20–80 seconds, depending on the operation and the performance of the attitude control system.
Additional time for sending the data to the satellite’s data processor (DHU) is also needed. These
operations and dead-times are taken into account in the design of the Astronomical Observation
Templates (AOT; section 5.5).
During a pointed observation, the IRC on-board controller communicates with the Attitudeand Orbit Control Unit (AOCU) via the DHU. When the IRC requests a Micro-Scan for dithering, it asks the AOCU to carry out the operation. The AOCU drives the satellite, then waits
for the stabilization. When the AOCU decides that the satellite has stabilized, with a preset
criteria, it sends back a signal to the IRC (via the DHU) to start the next exposure cycle. Thissequence continues until the AOT is completed. If, by some reason, the attitude of the spacecraft
exceeds the allowed range of the angle with respect to the Earth, the AOCU stops the pointing
mode attitude and returns back to the survey attitude, regardless of the observation. The IRC
will continue its programmed exposure sequences but of course those data acquired during the
maneuver are useless and will not be provided to the users. Current AOTs assume that one
pointed observation could last longer than 10 minutes and add extra exposure cycles at the end.

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5.3.2

The IRC Scan Observation

The two MIR cameras are able to operate in a “Scan Mode”. This operation mode will be used
for the IRC All-Sky Survey. An AOT to use this mode for pointed observations is also available.
The scan operation is realized by operating a part of the two-dimensional detector arrays.
Namely, one line of a detector array, perpendicular to the scan direction, is repeatedly readout
as the satellite scans the sky resulting in an image strip along the path. The NIR array has atechnical diﬃculty for such operation from the detector structure and will therefore not be usedin this observation mode.
Due to the strong constraint of the data down-link rate, it is impossible to transfer individual
pixel data to the ground in this mode, so some on-board preprocessing will be carried out . The
data from every four physical pixels are summed up to form one virtual pixel. By adjusting
the detector readout interval and scan speed, the shape of a virtual pixel can be controlled.
The current plan is to compose an almost square shape virtual pixel from four physical pixels
for the All-Sky Survey, i.e., the pixel size will be about 9.4 arcsec2 for the MIR-S channel, and10.0 × 9.4 arcsec2 for the MIR-L channel, respectively. The same sampling interval is applied
for the Slow-Scan observations in pointed mode, resulting in a very elongated virtual pixel size.
A smaller virtual pixel size, 2 × 2 mode, could be implemented technically and such a highresolution operation may be carried out if the necessary down-link rate can be guaranteed.

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The IRC All-Sky Survey

The IRC All-Sky Survey will be performed in parallel with the FIS All-Sky Survey, by using theS9W and L18W ﬁlters on the MIR-S and MIR-L, respectively. Since the scan path will shift by
4 arcmin per orbit in the cross-scan direction at the ecliptic plane, and the FoV size of the IRC
detector is 10 arcmin, every sky position will ideally be observed at least twice. At present, thefollowing spatial resolutions are considered for the IRC All-Sky Survey. The virtual pixel size
may be changed depending on the availability of the down-link rate.
Table 5.4.5: Current Speciﬁcation of the IRC All-Sky Survey
Camera Filter Virtual pixel size
MIR-S
S9W
9.4 × 9.4 arcsec2
MIR-L
L18W 10.0 × 9.4 arcsec2
cross-scan × in-scan

The ASTRO-F team is responsible for the planning and the data reduction of the IRC AllSky Survey, and will distribute catalogues to science users, as in the case for the FIS All-Sky
Survey.

5.4.1

Operation

During the survey observation, the detector is sampled continuously in the Scan mode. At
every 9 sec interval, the charge in the capacitor is reset. A reset causes a gap in the scan strip
corresponding to 2–5 sampling intervals. The width of the gap depends on any reset anomaly
in ﬂight.

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5.4.2

Expected Performance

Detection Limits
Expected 5σ detection limits of the IRC All-Sky Survey are summarized in Table 5.4.6. They
are estimated based on the latest performance evaluated in the ground tests. The IRC detection
limits are generally inﬂuenced by the background radiation, especially the Zodiacal light in the
Mid-Infrared range. In the case of All-Sky Survey, detector noise dominates the detection limits
so that only one set of numbers are given.
Table 5.4.6: 5σ Detection limits for the IRC All-Sky Survey
Point source Diﬀuse source∗(mJy)(MJy/sr)
MIR-S
S9W
80
25
MIR-L L18W
130
32
∗ Detection limits for diﬀuse source are per virtual pixel.
CameraFilterSaturation Limits
The saturation limits are given as the ﬂux level at which the detector signals deviate fromlinearity by 10 per cent. In fact, it is rather a rapid change from these levels to the absolutesaturation and it is reasonable to use these levels as the saturation limits.
Table 5.4.7: Saturation Limits for the IRC All-Sky Survey
CameraFilter

MIR-S
MIR-L

S9W
L18W

Fsat
(Jy)
36
93

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85

The IRC AOTs

The Astronomical Observation Templates (AOTs) of the IRC are summarized in Table 5.5.8.
Five AOTs are deﬁned for the instrument.
Table 5.5.8: IRC AOTs and user parameters.
AOT
00
02
03
04
11

PurposeDeep imaging
Imaging with 2 ﬁlters/ch.
Imaging with 3 ﬁlters/ch.
Spectroscopy
Slow-Scan# of ﬁlters
per ch.
1
2
3
1 or 2
1

DitheringNo
3 pos/ﬁlter
2 pos/ﬁlter
No
N/A

Filter3 choices
N/A (ﬁxed)
N/A (ﬁxed)
2 choices
3 choices

Parameters
Pointing
N, L
N, L
N, L
Ns, Np, Nh, Nc, Ls, Lc
Scan speed

The procedures for the AOT operations for IRC00, 02, 03 and 04 are shown in Figure 5.5.14.
As explained in the previous section ( 5.3.1) an AOT consists of a combination of exposure cycles(indicated as “Exp. C”), Filter Wheel rotations (“W”) and Micro-Scan operations (“M”). They
are optimized such that the maximum exposure time is obtained. The AOTs are designed touse more than 600 sec in the case that extra observation time is available. However, such extra
time cannot be guaranteed at present.

Figure 5.5.14: Observation sequences of the AOT IRC00, 02, 03, and 04. Yellow boxes indicatedas “Exp. C.” with ﬁlter / dispersion element names are exposure cycles. Orange boxes with“M” are Micro-Scan operations including stabilization, and light-blue boxes with “W” are Filter
Wheel rotations. Dead time for a Filter Wheel change depends on the relative position of theelements. The Green area on the right side is the extra observation time which is not guaranteed.

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5.5.1

Detection Limits and Saturation Limits: General Remarks

In the following sections, the expected detection limits for corresponding AOTs are given. These
are 5σ detection limits per single pointing opportunity, estimated based on the latest performance evaluated in the ground tests. Realistic exposure times in the operation sequence areconsidered, i.e., dead time due to Micro-Scan and ﬁlter change are taken into account. Repeating an observation of the same region with the same AOT improves the sensitivity by a factor√of n, where n is the number of pointing opportunities.
Detection limits in the imaging mode assumes aperture photometry of a well isolated source.
The photometry aperture is optimized for each band such that the signal (photons from thesource) to noise (dark + readout + sky background) ratio is maximized. The dark and readout
noises are estimated from laboratory measurements. Table 5.5.9 lists the best aperture radius
and the energy from the source falling into that radius. Table 5.5.10 summarizes the estimatedperformance parameters of the spectroscopic mode used to calculate the detection limits.
Table 5.5.9: Photometry radius and corresponding encircled energy fraction.
(1)
Camera
NIR
MIR-S
MIR-L

(2)
Filter
N2
N3
N4
S7
S9W
S11
L15
L18W
L24

(3)
Aperture radius(pixel)
1.90
1.90
1.93
0.63
0.77
0.77
1.00
1.20
1.67

(4)
Encircled energy(per cent)
70
65
62
27
32
31
34
33
38

(3) The optimized aperture radii that maximizes the S/N of the source.
(4) The fraction of energy within the aperture.

Table 5.5.10: Performance of the IRC Spectroscopic Mode
(1)
Camera(2)
Element

NIR

NP
NG
SG1
SG2
LG1
LG2

MIR-S
MIR-L

(3)
Resolution(λ/∆λ)
22@3.5 µm
135@3.6 µm
47@6.6 µm
34@10.6 µm
27@20.2 µm(4)(5)
Dispersion Res. bin(µm/pixel)(pixel)
0.058
2.7
0.010
2.7
0.057
2.5
0.101
3.1
— N/A —
0.153
4.9

(6)
(7)
Slit width(arcsec) (pixel)
3.0
2.05
3.0
2.05
5.0
2.13
5.0
2.13
7.0

2.92

(4) Length of the resolution element (∆λ) in pixel.
(6)(7) Slit width for the diﬀuse source spectroscopy.
Detection limits are given for two regions on the sky; at the Ecliptic Pole and the Ecliptic
Plane. The diﬀerence is simply due to the diﬀerence in the photon noise level of the zodiacal
emission. Detection limits for intermediate positions are provided by the Instrument Perfor-

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87

mance Tool available on-line. It is noted that the current estimates have an uncertainty due tothe laboratory measurement setup with the assumed telescope performance. We believe that thevalues are reliable to ∼ 30 per cent accuracy under the expected performance of the spacecraft
and the telescope.
The saturation limits are given as the ﬂux level at which the detector signals deviate fromlinearity by 10 per cent. In fact, it is a rather rapid change from these levels to the absolutesaturation and it is reasonable to use these levels as the saturation limits.

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5.5.2

IRC00: Deep Imaging Mode

Recommended Usage
IRC00 mode is prepared for deep imaging observations of high visibility regions where morethan 6 pointed observations can be easily obtained. This mode is also available for the OpenTime users under this constraint.
Description
The mode is designed to maximize the exposure time during a pointing opportunity. The
exposure cycle is simply repeated until the end of the pointed observation. No ﬁlter change
nor dithering operation is inserted. It is expected that observations with other ﬁlter(s) anddata redundancy by dithering are taken on additional pointing opportunities. It is strongly
recommended to observe the same sky position with at least three independent pointings withthe same ﬁlter setting. The necessary number of pointings is then multiplied by the number
of ﬁlters. This means that this AOT is only applicable in very high visibility regions, at highecliptic latitudes. It is recommended for observers to use this mode only in areas of the sky thatcan be observed more than 6 times (three dithers, two ﬁlters) by each of NIR/MIR-S or MIR-L.
Due to operational constraints it is very diﬃcult to ensure such high visibility for Open Timeand thus use of this mode has to be planned with great care.
Parameters
There are two parameters for this mode. The ﬁrst parameter speciﬁes the ﬁlter combination.
To simplify the operation and data handling, three preset ﬁlter combinations are deﬁned inTable 5.5.11.
Table 5.5.11: Filter combinations of IRC00
Set
a
b
c

NIR
N2
N3
N4

MIR-S
S9W
S7
S11

MIR-L
L18W
L15
L24

The second parameter designates the FoV reference position to be pointed to the target.
Table 5.5.12: Target position parameter for IRC00
Code
Description
N
Target position is placed at the center of the NIR/MIR-S cameras.
L
Target position is placed at the center of the MIR-L camera.
# Same as IRC02 and IRC03.

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89

Expected Performance

Detection Limits
The expected 5σ detection limits of the IRC00 observing mode are given in Table 5.5.13.
Table 5.5.13: Detection limits for AOT IRC00(1)
Camera

NIR
MIR-S
MIR-L

(2)
Filter

N2
N3
N4
S7
S9W
S11
L15
L18W
L24

(3)
(4)
Point source(µJy)
Pole Plane
7.9
8.7
3.7
4.3
7.2
8.8
35
51
29
47
38
63
68
120
84
160
180
330

(5)
(6)
Diﬀuse source(MJy/sr)
Pole Plane
0.033
0.036
0.014
0.017
0.026
0.032
0.062
0.091
0.050
0.081
0.064
0.11
0.093
0.17
0.091
0.17
0.16
0.30

(3)–(6) 5 σ detection limits for a pointed observation: 9 exposure cycles for one ﬁlter in apointing opportunity.
(3)(5) Pole = at Ecliptic pole.
(4)(6) Plane = at Ecliptic plane.
(7) Estimated as the 10 per cent linearity range for an exposure time of 2.3 sec and 0.6 sec for
NIR and MIR-S/L, respectively.

Saturation Limits
The saturation limits of the IRC00 mode are given in Table 5.5.14. These limits are deﬁned as
the 10 per cent linearity range for the short exposure (2.3 sec and 0.6 sec for NIR and MIR-S/L,
respectively) of the exposure cycle.
Table 5.5.14: Saturation limits for IRC00
Camera

NIR
MIR-S
MIR-L

Filter
N2
N3
N4
S7
S9W
S11
L15
L18W
L24

Fsat
(mJy)
780
250
360
1800
1000
1800
2500
3200
23000

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5.5.4

IRC02: Imaging with Two Filters

Recommended Usage
IRC02 mode is prepared for general purpose imaging observations. This mode can be used atany sky position even in low visibility regions.
Description
A sky position is observed with two ﬁlters each of which consist of three images with small
position shifts by dither. i.e., this AOT can provide a self-standing data set in just a single
pointing. The net exposure time per ﬁlter is about 1/3 of IRC00.
Parameters
The ﬁlters used in this AOT are ﬁxed as the following. The ﬁlter set can cover the whole IRC
wavelength range except for that of N2 (wide bands are not included).
Table 5.5.15: Filter combination of IRC02
CameraFilterNIR
N3 & N4MIR-S
S7 & S11

MIR-L
L15 & L24

A parameter specifying the reference to the target position must be given for this AOT.
Table 5.5.16: Target position parameter for IRC02
Code
Description
N
Target position is placed at the center of the NIR/MIR-S cameras.
L
Target position is placed at the center of the MIR-L camera.
# Same as IRC00 and IRC03.

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Expected Performance

Detection Limits
The expected 5σ detection limits of the IRC02 observing mode are given in Table 5.5.17.
Table 5.5.17: Detection Limits for IRC02(1)
Camera(2)
Filter

NIR

N3
N4
S7
S11
L15
L24

MIR-S
MIR-L

(3)
(4)
Point source(µJy)
Pole Plane
6.4
7.4
13
15
60
82
66
110
120
210
310
570

(5)
(6)
Diﬀuse source(MJy/sr)
Pole Plane
0.025
0.029
0.046
0.056
0.11
0.16
0.11
0.18
0.16
0.29
0.29
0.51

(3)–(6) 5 σ detection limits for a pointed observation: 3 exposure cycles per ﬁlter in a pointing
opportunity. Numbers for N2, S9W, and L18W are only for reference. These ﬁlters are notused in the IRC02.
(3)(5) Pole = at Ecliptic pole.
(4)(6) Plane = at Ecliptic plane.

Saturation Limits
The saturation limits of the IRC02 mode are given in Table 5.5.18. These limits are deﬁned as
the 10 per cent linearity range for the short exposure (2.3 sec and 0.6 sec for NIR and MIR-S/L,
respectively) of the exposure cycle.
Table 5.5.18: Saturation Limits for IRC02
CameraFilter

NIR

N3
N4
S7
S11
L15
L24

MIR-S
MIR-L

Saturation Limit
(mJy)
250
360
1800
1800
2500
23000

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5.5.6

IRC03: Imaging with Three Filters

Recommended Usage
IRC03 is used for general purpose observing. It observes a target with all three ﬁlters ofeach camera. This mode is not recommended for sky areas of low visibility where two pointed
observations cannot be guaranteed.
Description
IRC03 is similar to the IRC02 mode but diﬀers in the number of ﬁlters and dithering positions
in one pointing opportunity. The target sky position is observed by three ﬁlters of each camera.
Instead, only one Micro-Scan (2 dithered images) is applied for each band. To obtain reliable
results it is recommended to have suﬃcient redundancy by observing the same sky position atleast twice on diﬀerent orbits. The net exposure time per ﬁlter is thus 2/9 of the IRC00.
Parameters
IRC03 uses all ﬁlters of the IRC.
Observers should be reminded that the wavelength coverage of the S9W and L18W ﬁlters
overlap with the two other ﬁlters in the same camera. The advantage of using this mode insteadof IRC02 has to be carefully considered from the scientiﬁc viewpoint of each observation.
Table 5.5.19: Filter combination of IRC03
CameraFilterNIR
N2 & N3 & N4

MIR-S
S7 & S9W & S11

MIR-L
L15 & L18W & L24

The position reference has to be given as with IRC00 and 02.
Table 5.5.20: Target position parameter for IRC03
Code
Description
N
Target position is placed at the center of the NIR/MIR-S cameras.
L
Target position is placed at the center of the MIR-L camera.
# Same as IRC00 and 02.

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Expected Performance

Detection Limits
Expected 5σ detection limits of the IRC03 observing mode is given in Table 5.5.21.
Table 5.5.21: Detection limits for IRC03(1)
Camera

NIR
MIR-S
MIR-L

(2)
Filter

N2
N3
N4
S7
S9W
S11
L15
L18W
L24

(3)
(4)
Point source(µJy)
Pole Plane
17
18
7.8
9.1
15
19
74
110
62
100
81
130
140
260
180
330
380
690

(5)
(6)
Diﬀuse source(MJy/sr)
Pole Plane
0.069
0.076
0.030
0.035
0.056
0.069
0.13
0.19
0.11
0.17
0.14
0.22
0.20
0.36
0.20
0.36
0.35
0.63

(3)–(6) 5 σ detection limits for a pointed observation: 2 exposure cycles per ﬁlter in a pointing
opportunity.
(3)(5) Pole = at Ecliptic pole.
(4)(6) Plane = at Ecliptic plane.

Saturation Limits
The saturation limits of the IRC03 mode are given in Table 5.5.22. These limits are deﬁned as
the 10 per cent linearity range for the short exposure (2.3 sec and 0.6 sec for NIR and MIR-S/L,
respectively) of the exposure cycle.
Table 5.5.22: Saturation Limits for IRC03
Camera

NIR
MIR-S
MIR-L

Filter
N2
N3
N4
S7
S9W
S11
L15
L18W
L24

Saturation Limit
(mJy)
780
250
360
1800
1000
1800
2500
3200
23000

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5.5.8

IRC04: Spectroscopic mode

Recommended Usage
IRC04 is for spectroscopic observations with the IRC. Observation with this mode can be madein regions of moderate to high visibility which allow two or more pointed observations at eachtarget position.
Description
In this mode, the IRC dispersion elements are used to take spectra of the targets. The MIR-S
cameras always use two dispersion elements to cover their full wavelength range. However, due
to the loss of the LG1 grism, the MIR-L camera uses LG2 only. This results in an unavoidable
gap between 13–18 µm. The wavelength coverage of the NP (moderate resolution) and NG (highresolution) are similar to each other, and the user should choose one of these two for this AOT.
As we explain in section 5.1.5, a slit is provided in each camera in order to observe diﬀuse
radiation. The NIR camera also has an entrance aperture (slit) for a point source, to enableconfusion-less spectroscopy. It is used with the NG in high-resolution mode (see below). The
point source density in the FoV of the MIR-S/L cameras is expected to be low enough that they
will be able to observe point sources directly in the imaging ﬁeld.
A short exposure image in the N3, S9W, and L18W ﬁlters, is taken by each camera respectively, for pointing alignment.
No dithering operation is carried out during one pointing in this AOT. Observers are highly
recommended to observe the same area of the sky on two or more independent orbits to ensure
data redundancy.
Parameters
The ﬁrst parameter speciﬁes the dispersion element for the NIR camera.
Table 5.5.23: Dispersion element selection in IRC04
Set
a
b

NIR
NP (N3)
NG (N3)

MIR-S
SG1 & SG2 (S9W)
SG1 & SG2 (S9W)

MIR-L
LG2 (L18W)
LG2 (L18W)

The second parameter gives the position reference for this AOT. It is more complicated thanother AOTs because of the presence of the slits. Np and Nh are intended to be used with theNG, and others with the NP. Only the Ns slit can be used with both elements.

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Table 5.5.24: Target position parameter for IRC04
Code
Nc
Ns
Lc
Ls
Ns
Np
Nh

Description
Available for Set a (NP)The target position is placed at the center of the NIR/MIR-S (FoV).
The target position is placed at the common slit of the NIR/MIR-S.
The target position is placed at the center of the MIR-L (FoV).
The target position is placed at the slit of the MIR-L.
Available for Set b (NG)The target position is placed at the common slit of the NIR/MIR-S.
The target position is placed at the point source aperture of the NIR.
The target position is placed at the high-resolution slit of the NIR.

95

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5.5.9

Expected Performance

Detection Limits
Table 5.5.25 and 5.5.26 summarizes 5 σ detection limits of the IRC04 observing mode. These
numbers are for 8 exposure cycles, i.e. a single pointing opportunity. Line detection limits are
given for integrated line ﬂuxes of an isolated line that can be treated as a point source (byaperture photometry technique at the line position).
Continuum detection limits for point source (slitless or with Np) are per pixel in the direction of dispersion. In other words, they are the ﬂux that provides a S/N = 5 signal after
being integrated over the pixels along the spatial direction of one row in the spectral dispersion
direction. Averaging over the dispersion direction can improve the detection limits with a costof spectral resolution. q On the other hand, continuum detection limits for diﬀuse source (withslit) are given per pixel. Integrating over the slit length can improve the sensitivity. If you
integrate the spectrum
over 0.8 arcmin (about 33 pix), the NP detection limits at the Eclipticpole can be 0.26/ (33) = 0.045 [MJy/sr].
These numbers are for the monotonic limit at the wavelength band center. The spectral
response functions (RSRF; Figures 5.2.3, 5.2.4, 5.2.6) have to be taken into account for the limits
at other wavelengths. Moreover, the spectral resolution changes as a function of wavelength in
the case of NP, the prism dispersion element, and this also has to be considered for any detailed
estimate of the performance. In this case, the resolution is approximately about 1/4 of thecenter value (at 3.5 µm) at 2 µm and about twice at 5 µm.
Table 5.5.25: Line Detection Limits for IRC04
Camera

NIR

MIR-S

MIR-L

Element

λcenter(µm)

Position

NP
NP
NG
NG
SG1
SG1
SG2
SG2
LG1
LG2
LG2

3.5
3.5
3.6
3.6
6.6
6.6
10.6
10.6

Nc
Ns
Np
Ns/Nh
Nc
Ns
Nc
Ns

20.2
20.2

Lc
Ls

Point Source Diﬀuse Source(10−17 W m−2 )Pole Plane Pole
Plane
0.18
0.26
—
—
—
— 0.13
0.13
0.22
0.27
—
—
—
— 0.19
0.19
1.3
1.8
—
—
—
— 0.92
0.94
1.3
2.0
—
—
—
— 0.630.68— N/A —
2.2
3.9
—
—
—
— 0.92
1.1

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Table 5.5.26: Continuum Detection Limits for IRC04
Camera

NIR

MIR-S

MIR-L

Element

λcenter(µm)

Position

NP
NP
NG
NG
SG1
SG1
SG2
SG2
LG1
LG2
LG2

3.5
3.5
3.6
3.6
6.6
6.6
10.6
10.6

Nc
Ns
Np
Ns/Nh
Nc
Ns
Nc
Ns

20.2
20.2

Lc
Ls

Point Source(mJy)
Pole Plane
0.05
0.08
—
—
0.48
0.60
—
—
1.8
2.5
—
—
2.0
3.3
—
—
— N/A —
5.5
9.8
—
—

Diﬀuse Source(MJy/sr)
Pole
Plane
—
—
0.26
0.27
—
—
2.8
2.8
—
—
3.8
3.9
—
—
3.5
3.9
—
7.0

—
8.4

Saturation Limits
Table 5.5.27: Saturation Limits for IRC04
Band

NIR
MIR-S
MIR-L

Line limit

NP
NG
SG1
SG2
LG1
LG2

(10−14 W/m−2 )
7.4
11
55
48
91

Continuum limit
Point source Diﬀuse source(Jy)(GJy/sr)
2.2
15
24
160
76
230
77
270
— N/A —
220
680

ASTRO-F Observer’s Manual

98

5.5.10

IRC11: Slow-Scan mode

Recommended Usage
IRC11 is recommended for observations of large areas to moderate depth.
Description
In this mode, the detectors are operated following the same procedure as the All-Sky Survey.
As described in Section 5.3.2, the NIR detector is not used in this mode. Observers can choosea pair of ﬁlters from the MIR-S and MIR-L camera from three presets.
The scan pattern of IRC11 is identical to that of FIS02. One round trip scanning is carriedout in a single pointing opportunity. Two scan speeds, 15 or 30 arcsec/sec are available. See
Section 4.4.3 for details.
Parameters
The following three ﬁlter sets are available.
Table 5.5.28: Filter combinations of IRC11
Set
a
b
c

NIR
N/A
N/A
N/A

MIR-S
S9W
S7
S11

MIR-L
L18W
L15
L24

The second parameter designates the FoV reference position to be pointed to the target. As
for FIS02, the position is regarded as the beginning of the scan.
Table 5.5.29: Target position parameter for IRC11
Code
N
L

Description
Target position is placed at the center of the NIR/MIR-S cameras.
Target position is placed at the center of the MIR-L camera.

The scan speed has to be selected for the third parameter. Current estimates of the scan
length in one pointing are approximately 1 and 2 deg for 15 and 30 arcsec/sec, respectively.
Table 5.5.30: Scan speed available for IRC11
Scan speed

15 or 30 arcsec/sec

Version 3.2 (November 29, 2005)

5.5.11

99

Expected Performance

Detection Limits
Expected 5σ detection limits of the IRC11 observing mode are given in Table 5.5.31.
Table 5.5.31: Detection limits for IRC11: Slow-Scan mode
(1)

(2)

Scan speed (arcsec/sec)

MIR-S
MIR-L

S7
S9W
S11
L15
L18W
L24

(3)

(4)
(5)
(6)
Point source (mJy)
15
30
Pole Plane Pole Plane
11
12
22
23
8
9
14
16
8
10
16
18
13
19
24
30
12
19
21
29
44
65
80
100

(7)
(8)
(9)
(10)
Diﬀuse source (MJy/sr)
15
30
Pole Plane Pole Plane
3.8
4.1
7.3
7.7
2.3
2.7
4.5
4.9
2.4
2.9
4.5
5.1
3.4
4.8
6.2
7.7
2.6
3.8
4.6
6.0
8.4
11.1 15.6
19

Saturation Limits
Table 5.5.32: Saturation limits for IRC11: Slow-Scan mode
(1)

(2)

Scan speed (arcsec/sec)
S7
MIR-S
S9W
S11
L15
MIR-L
L18W
L24

(3)
(4)
Point source
(Jy)
15
30
3.3
6.5
1.8
3.6
3.0
6.0
4.3
8.6
4.4
8.8
18
23

(5)
Diﬀuse source(MJy/sr)
15 & 30
220
130
130
200
150
530

100

5.6

ASTRO-F Observer’s ManualNotes and Restrictions for the IRC Observations

1. Bright sources We believe that bright sources that saturate the detectors will not cause
permanent damage. However, very bright sources are known to cause after eﬀects (i.e.
latent images, etc.), degrading the performance of the detectors for more than an hour.
Therefore, some constraint for observing ﬁelds with very bright sources may be given for
IRC observations, according to the results obtained in the PV phase.
2. Eﬀects of South Atlantic Anomaly (SAA) We believe that the passage through the
SAA does not cause serious damage to the detectors. However, detector performance
may be signiﬁcantly degraded immediately after the SAA. Current observation schedulingallocates IRC observations relatively close to the SAA region (say, 5 minutes after thepassage). The real SAA region and its inﬂuence to ASTRO-F observations will be clear
only after the PV phase. Observation of targets very close to the SAA may have to bereconsidered or may even be canceled.

Chapter 6

Data Reduction and ProductsThis section describes the data handling policy, current scope of the data reduction support anddata products. Please note that the data reduction software are under development and only apreliminary plan is explained here. Details of the data handling and data reduction guide will
be prepared in the future as User manuals.

101

ASTRO-F Observer’s Manual

102

6.1

Basic Policy

The basic policy of the ASTRO-F data reduction is the following.
• Processing of the Large Area Surveys, including the All-Sky Survey, is to be carried out in aunique and uniform manner by a dedicated team organized by the ASTRO-F Project withthe products being distributed to the users. Released products are diﬀerent depending on
the diﬀerent surveys.
• The reduction of pointing observation data is the observer’s responsibility. The ASTRO-F
team will provide the necessary information to handle the data, e.g., calibration data andsoftware to correct instrument related characteristics. Astronomical analysis such as point
source extraction should be carried out by commonly available software.
We welcome any users to participate in the data reduction activity at any level of the work,from giving quantitative reports of their data reduction results to proposing new algorithms to
the programs.

Version 3.2 (November 29, 2005)

103

6.2

Large Area Surveys

6.2.1

All-Sky Survey

6.2.2

ASTRO-F Point Source Catalogues

We plan to release several catalogues with diﬀerent detection limits, completeness and reliability,
during the progress of the data reduction. Release dates in Table 6.2.1 are internal for the
ASTRO-F Project team members and are also for purposes of quality control. After one yearof internal use, the data will be released to the entire world astronomical community. In thefollowing, BOS is the beginning of survey which starts two months after the launch, and EOS
is the end of survey, estimated to be about 1.5 years after the launch.
A draft of catalogue contents are listed in Table 6.2.1.
Table 6.2.1: ASTRO-F Point Source Catalogues
Catalogue
FIS Flux of Known Sources

Team Release
BOS + 2 month –EOS + 2 month

Description
FIS ﬂuxes of known sources such as
IRAS point sources. Released incrementally during the survey.
Bright Source Catalogue (BSC) EOS+ ≤ 1 year
Uniform (but relatively shallow) detection limit over the entire sky.
High redundancy resulting in higher
reliability.
Faint Source Catalogue (FSC)
BSC + ≤ 1.5 year Data redundancy is used to improvethe detection limit. Generated only
in regions where such improvementis eﬀective.
BOS: Beginning of Survey; EOS: End of Survey = He boil-oﬀ[FIS Flux of Known Sources ]This catalogue will be incrementally provided between BOS + 2 month and EOS + 2
month. The main purpose of this catalogue is a data consistency check with the IRAS andMSX catalogues but it will also provide new ﬂux information beyond 100 µm for brightsources. Hence all the sources in these catalogues should already be known. The catalogue
will be used as a prototype for the Bright Source Catalogue.
[Bright Source Catalogue ]The Bright Source Catalogue is regarded as the main product of the ASTRO-F All-Sky
Survey, and is expected to be released before EOS + 1 year. The catalogue will provide a
uniform quality data set over the entire sky. Sources are extracted from each single data
scan and then cross-checked. This means that the high-ecliptic latitude regions where a lotof data redundancy is available will have better reliability but no advantage in sensitivity..
[Faint Source Catalogue ]This catalogue will be released within 1.5 years after the release of the Bright Source
Catalogue but may be provided incrementally. In contrast to the Bright Source Catalogue,
and rather similar to the IRAS Faint Source Catalog, all available scans for each area ofthe sky will be ﬁrst integrated into a data set, then source extraction will be applied tothis merged data. This will improve the detection limits in the more redundant areas,

ASTRO-F Observer’s Manual

104

i.e., the detection limit will not be uniform over the entire sky. This catalogue may be
produced only where such processing is deemed to be eﬀective.
Even deeper catalogues may be produced upon the data quality, improvement of the dataprocessing, and scientiﬁc requirements.
Note that the above explanation was originally envisaged for the FIS All-Sky Survey. We
plan that the IRC All-Sky Survey will follow the same policy and will be released together with
the FIS catalogues.
The catalogues are expected to contain millions of sources, resulting in a data size of a few to
ten Gigabytes. The data will be distributed primarily via the Internet with a suitable interface
for data query. It is planned that the ASTRO-F database will have a standard interface to
Virtual Observatory (VO). Subset contents of the catalogue may also be distributed by DVD orother appropriate media.

6.2.3

Image Maps

Release of scientiﬁc quality image maps from the ASTRO-F All-Sky Survey data has not beenoﬃcially decided by the project. This is because we are not convinced that the stability of
the instrument performance, especially the dark current level, is good enough to produce usefulimage maps. A ﬁnal decision will be made after the PV phase. We expect that imagelets of small
patches of the sky can be created relatively easily but large area maps will be more diﬃcult.
Preparation for the creation of the image maps has been started, initially aiming to process
pointing observation data and survey data of small sky areas with high redundancy.
Table 6.2.2 shows the preliminary plan for image maps with the ASTRO-F All-Sky Survey
data.
Table 6.2.2: ASTRO-F Image Map Products
Products
Crude Map

TBD

Area

Imagelets

1 deg × 1 deg

Image Atlas (ver.1)

10 deg × 10 deg

Image Atlas (ver.2)

10 deg × 10 deg

Description
For ASTRO-F team internal use
only. This map will be used for theveriﬁcation and testing of the data.
Images of interesting small extended
sources and regions.
Destriped image maps of wide ﬁelds(all-sky) of scientiﬁc quality.
Same as ver. 1 but the Zodiacal
light is subtracted.

Processed Raw Data
The raw data of the ASTRO-F survey (with the appropriate data processing and calibration)
will be available after EOS+5 year when the calibration has been stabilized.

Version 3.2 (November 29, 2005)

6.3

105

Pointed Observation Data

In principle, the users will receive the following data set.
• Raw data in appropriate FITS format.
• Automatic Processed results from the data reduction pipeline (hereafter Pipeline output)
in the same FITS format. Especially in early stage of the mission, the automatic processing
may not be suﬃcient enough to create scientiﬁc quality data. Therefore this output should
be regarded as a reference for the user’s own data reduction.
• Data reduction Tool Kit package. The package will consist of modules for automatic dataprocessing. A Pipeline with user friendly interface and additional support tools.
• Calibration data appropriate for the data.
Data will be primarily distributed on-line via ISAS’s Data ARchive and Transmission System(DARTS). The Tool Kits will be provided via the User Support web page. Details will be given
later.

6.3.1

FIS Data Reduction: AOT FIS01, 02

Users will receive raw and Pipeline processed data in the format of time series detector output.
All information relevant for the data reduction will be stored in one FITS ﬁle along with thedata itself. In addition, a simple co-add image created from the data will be attached.
In the early stages (Phase 1), the pipeline for pointed observations is expected to be quite
similar to that of the All-Sky Survey. i.e., it does not use data redundancy provided by theround-trip scan eﬃciently. Such functions will be implemented in later releases, possibly by
beginning of Phase 2. The goal for the data quality of the pipeline output is minimum scientiﬁc
grade; not wrong but with large uncertainties.
Users may want to process the data by themselves with the Tool Kit, following the Pipeline
processing steps interactively by adjusting the module parameters.
We will also provide an image construction tool. The minimum function of the tool will beto create a simple co-add map and optionally some kind of destriping algorithm. Data analysis
with the images is expected to be made with existing software packages or users’ programs. The
ASTRO-F project will not support this part of the data reduction.
The software programs are currently developed on IDL 6.1.

6.3.2

FIS Data Reduction: AOT FIS03

As repeated many times, observation with this AOT is challenging and the data reduction is
not an exception. Figure 6.3.1 illustrates the data reduction procedure for FTS mode. Oneof the most critical parts is to determine the centre positions of the interferogram for every
detector pixel. The current estimates of the scope is that we will need many cut-and-try eﬀorts.
Transients are another critical issue. For the time being an empirical approach is applied forthe transient correction. More sophisticated methods will be investigated.
In general, we foresee that substantial experience will be necessary to handle the FIS data,
especially for the FTS observations; and therefore, close contact with specialists in the project
team will be crucial. There is a plan to organize a “user consortium” to discuss various technicalproblems and contribute activity to the data reduction.

ASTRO-F Observer’s Manual106

Figure 6.3.1:

6.3.3

IRC Data Reduction: AOT IRC00, 02, 03

IRC imaging mode data reduction is expected to be rather more straightforward following asimilar procedure to ground base observations. In addition there will be instrument speciﬁc
corrections such as removing ghost images and spikes by charged particle hits.
One of the issues to be concerned with is to ensure good quality dark frame and ﬂat-ﬁelddata. Both will be measured by the instrument team in the calibration dedicated observations
through the PV phase and during the observing phases. Therefore, the calibration data will be
updated and improved during the mission. Users may want to reprocess the data with updated
calibration data by themselves.
Our support is limited to correction and calibration of each observation. We do not plan tosupport photometry of the calibrated images or mosaicing.
Currently the system is developed in the IRAF environment. Exporting it to IDL is TBD.

6.3.4

IRC Data Reduction: AOT IRC04

Data reduction of this mode is similar to that of the imaging mode for the basic processing of
the dark subtraction, despiking, or bad pixel masking etc. In this AOT an image is also takenwith one ﬁlters per each camera. Source positions are extracted from this image and applied tothe spectroscopy images to clip each target.
Several relevant technical issues for the spectroscopic observations are listed below.
• For reliable spectroscopic observations we need accurate and wavelength dependent ﬂatﬁelding.

Version 3.2 (November 29, 2005)

107• Wavelength calibration is automatically made by the relationship between imaging position and the spectroscopy data. It also includes distortion correction of the image (Section 5.2.1).
• There may be confusion from nearby targets in the case of observations in crowded regions. Currently these are not decomposed but rather masked in the overlap region forthe measurements.
• Measurement of the spectrum is carried out in an aperture on the spectroscopic image.
The software attempts to ﬁt an aperture automatically but the user may want to adjust
this step interactively.
The tools are being written in IDL.

6.3.5

IRC Data Reduction: AOT IRC11

Data taken by this mode is processed by the pipeline developed for the All-Sky Survey. As
for FIS02 (and 01), the pipeline output in the early stage of the mission may not use the data
redundancy (round trip scan) eﬃciently.
Data reduction of this observing mode is still in the investigation phase and details will begiven later.

ASTRO-F Observer’s Manual108

REFERENCES
Fujishiro, N., et al., 2005, Proc. SPIE, in press
Jeong, W.-S., et al., 2003, PASJ 55, 717
Kim, U., et al., 2005, Proc. SPIE, in press
Watanabe, M., 2001, “ASTRO-F Visibility Tool”, ASTRO-F internal report, 2001 Sep. 12 (in Japanese)

Appendix A

Confusion Limits and ASTRO-F
A.1

Introduction

The absolute sensitivity of a given space born instrument / telescope system is governed by 2
basic classes of noise; instrumental & photon noise and sky confusion noise. Instrumental noise
is intrinsic to the system and is contributed to by such factors as readout noise, dark current
ﬂuctuations, ﬂat-ﬁelding uncertainties etc. On the other hand, sky confusion noise is observational dependent and can arise from both the superposition of sources in crowded ﬁelds andfrom extended structures which vary in surface brightness on scales of the telescope and instrument resolution. At infrared wavelengths, the major components are the sky confusion due tointerplanetary dust bands (zodiacal light, in the range 5 < λ < 60 µm), the sky confusion due todust emission from irregular interstellar clouds at high galactic latitudes known as the “infraredgalactic cirrus” (Low et al. 1984), confusion due to stars at near infrared wavelengths and theconfusion due to the ﬂuctuation of the extragalactic background built up by the superpositionof individual faint sources below the resolution of the telescope beam.
Confusion can cause centroid position shifts and ﬂux uncertainties leading to positional
errors and spurious sources, therefore careful consideration must be given to the treatment ofconfusion noise and the confusion limit, as in reality, the confusion noise is a convolution of theobservational phenomenon and the observing instrument.

A.2

Confusion due to Infrared Cirrus

The ﬂuctuations in the surface brightness of extended structure on similar scales as the resolutionof the telescope/instrument beam can produce spurious events that can be easily mistaken for
genuine point sources, since the existence of a source is usually simply derived from the diﬀerencein signal between the on source and some background position. Such extended structure isobserved in whispy neutral interstellar dust in the Milky Way that is heated by the interstellar
radiation ﬁeld and is known as the infrared cirrus (Low 1984). Cirrus emission peaks at far IR
wavelengths (100 µm) but was detected in all 4 IRAS bands (Helou and Beichman 1990). The
galactic cirrus is a function of galactic latitude and is serious for wavelengths longer than 60 µm.
Quantatively, this uncertainty in brightness between on source and background positions canbe expressed as the second order structure function C(θ), of the sky brightness, whereB(x + θ) + B(x − θ) 2
| >x ,(A.2.1)
2is the mean diﬀerence squared over all on/oﬀ pairs, where B(x) is the on source positionbrightness at x, and B(x±θ) are the oﬀ source reference position brightnesses with a separation ofθ from the source (2 oﬀ source positions to correct for any gradient in the background brightness
C(θ) =< |B(x) −109

ASTRO-F Observer’s Manual

110emission). Fluctuations on scales of orderθ will aﬀect the measurements, the noise due to these
ﬂuctuations being given by σ (Jy)= Ω C(θ), where Ω is the solid angle of the measuring
aperture such that the limiting noise due to the ﬂuctuation of the sky brightness is given by
the value of C(θo ) at the resolution limit θo of the instrument. Thus in the case of backgroundﬂuctuations with structure, the measured confusion noise is a function of both the angularseparation of the source measurements and the aperture of the instrument.
Gautier et al. (1992) analysed the spatial structure of the IRAS infrared cirrus at 100 µm
using the Fourier transform power spectrum of its’ ﬂuctuations and found the well known relation
between the power, P and scale length, d or its’ reciprocal, the spatial frequency, k,
P = Po (do αk) = Po ( )α ,dko

(A.2.2)

with a power Po on a scale do or corresponding spatial frequency ko . Empirically, a power
spectrum of α ≈ 3 and relation between the reference power and average sky brightness, Po ∝<B(100 µm) >3 was found. Relating the cirrus power spectrum to the structure function, thenoise scales as,σ ∝ Po1/2 (d 1− α) 2,
do(A.2.3)

Helou & Beichman (1990) assume the relation Po ∝< B >3 is true for all wavelengths(replacing < B(100 µm) > with < B(λ) >) and taking the scale d as the resolution of atelescope (of diameter D), ∼ λ/D express the 1σ cirrus confusion noise in equation A.2.3 as,λD −2.5 < Bλ > 1.5σ(λ)= 0.732()2.5 (
)
(
) ,
1mJy
100µm
0.7m
1MJysr−1(A.2.4)Thus the cirrus confusion noise will depend on both the variation of the surface brightnessof the background structure and the resolution of the telescope, with the noise becoming lesssigniﬁcant for larger aperture sizes such that the next generation of space telescopes (Herschel,
SPICA) should not be so severely aﬀected by cirrus confusion over most of the sky. Herbstmeier
et al. (1998) have shown that assuming the sky brightness ﬂuctuations and instrumental noise2 ≤ σ 2 + 2σ 2 , thus assuming
are statistically independent then the total measured noise is σtot
instsky
equality a lower limit for the sky brightness noise can be estimated. Kiss et al (2001) haveanalysed 40 sky regions with the ISOPHOT instrument on ISO concluding that the cirrus noise
was consistent within a factor of 2 to the formulation of Helou & Beichman (1990).
Applying equation A.2.4 to the ASTRO-F FIS instrument and assuming a range of average
brightnesses for < Bλ > gives us the results for the cirrus confusion noise in Table A.2.1.
Note that Jeong et al. (2004) have carried out simulations of the cirrus background for the
ASTRO-F FIS wide bands (WIDE-S & WIDE-L). They have concluded that the formula ofHelou & Beichman (1990) systematically underestimates the cirrus noise in the dark sky case(< 1 MJy/sr) and overestimates the cirrus noise in the bright sky case (> 3 MJy/sr). Therefore
the limits in Table A.2.1 should be treated as lower and upper limits accordingly. Table A.2.2
shows the corresponding estimates extrapolated to the IRC bands,

A.3

The Extragalactic Source Confusion Limit

The galaxy confusion limit is deﬁned as the threshold of the ﬂuctuations in the background skybrightness below which sources cannot be discretely detected in the telescope beam ∼ λ/D.
Thus the ﬂuctuation noise arises from the same origin as the galaxies that one is aiming todetect! If we assume galaxies are distributed as a power law in ﬂux, S, down to some detection
limit Slim ,Version 3.2 (November 29, 2005)

111

Table A.2.1: Estimated confusion limits due to galactic cirrus in ASTRO-F FIS bands assuming
the model of Helou & Beichman (1990).
Band
N60
WIDE-S
WIDE-L
N160

Sensitivity
5σ (mJy)
600
200
400
800

Confusion Limit (5σ, mJy) for < Bλ >0.5 MJy/sr 3 MJy/sr
10 MJy/sr
0.37
5.49
33
0.77
11
68
3.6
54
330
4.3
63
388

Table A.2.2: Estimated confusion limits due to galactic cirrus in ASTRO-F IRC bands from theASTRO-F Large Area Survey (NEP) document.
Band
N2
N3
N4
S7
S9W
S11
L15
L18W
L24

Sensitivity (600s)
5σ (µJy)
7.9
3.7
7.2
32.8
26.1
37.1
68.4
87.3
180.8

Confusion Limit (5σ, µJy) for < Bλ >0.5 MJy/sr 3 MJy/sr
10 MJy/sr
<1
<1
<1
<1
<1
<1
<1
<1
<1
<1
<1
<1
<1
<1
<1
<1
<1
1.6
<1
1.1
6.9
<1
1.6
9.5
<1
3.1
19

ASTRO-F Observer’s Manual112

N (> Slim ) = Nlim (

S −α
) ,
Slim

(A.3.5)where α,is the slope of the integral source counts (where α = 1.5 for a Euclidean Universe)
and Nlim is the number density at the limiting sensitivity Slim . Assuming that the counts ﬂatten
at some faint ﬂux, S0 , i.e. α(S0 ) = 0, then the intensity of the background (in Jy/sr) up to somemaximum ﬂux, Smax , corresponding to these sources is given by, Smax

I=

S0

S

dN
dS,
dS

(A.3.6)

The ﬂuctuations contributed by sources below the detection limit Slim are given by thesecond moment of the diﬀerential source counts dN/dS, σ in MJy/sr,
 Slim

σ2 =

S0

S2

dN
dS,
dS

(A.3.7)

Assuming the power law distribution of sources given in equation A.3.5 and A.3.7 can beevaluated to give,2
σ 2 = Nlim Slim

S0 2−αα(1 − (
)
),
2−α
Slim

(A.3.8)For the Euclidean case the dominant sources contributing to the background intensity arethose just below the detection limit Slim (Matsuhara et al. (2000), Lagache et al. (2000)).
However the strong evolution detected in the galaxy population steepens the source counts andproduces super Euclidean slopes in excess of α > 1.5 and the sources around the detection limit
also contribute signiﬁcantly to the ﬂuctuations in the background intensity.
Rigorous theoretical deﬁnitions of confusion have been presented by Scheuer (1957) andCondon (1974). Hogg (2001) has highlighted the more practical aspects of galaxy confusion
noise limitations. An analytical derivation broadly following Franceschini et al (1989) is given
below. Note that the clustering of sources will complicate the confusion noise (e.g. as in thecase of radio sources, Condon (1974)) although here, for clarity, we do not treat this eﬀect (seeFranceschini et al (1989), Takeuchi (2003) for discussions on the eﬀect on the confusion limit of
the clustering of sources).
Assuming that the sources are distributed randomly over the sky described by a power law
form N (S) ∝ S −α and a corresponding diﬀerential distribution given by n(S) = kS −γ whereγ = α − 1, then the detector response to a source of ﬂux S, at a position (θ, φ) from the axis ofa detector beam of proﬁle (point spread function) f (θ, φ) is given by x = Sf (θ, φ). Hence themean number of responses with amplitudes between x, x + dx in a solid angle dΩ is given by;R(x) =

n(x/f (θ, φ))dΩ,f (θ, φ)Ωb

(A.3.9)where Ωb is the solid angle of the beam in steradians. Note that for the power law distributionof sources discussed above, equation A.3.9 can be rewritten as,−γ

R(x) = kxΩbf (θ, φ)γ−1 dΩ = kx−γ Ωe ,(A.3.10)where Ωe = f (θ)γ−1 dΩ is the eﬀective beam size (Condon 1974). Taking the second momentof the R(x) distribution (the variance) gives the ﬂuctuation of the response,σ ;
 xc

σ2 =

0

x2 R(x)dx,(A.3.11)

Version 3.2 (November 29, 2005)

113

where xc is a cut oﬀ response introduced to stop the variance from diverging at bright source
ﬂuxes. More practically the confusion limit xc (corresponding to a cut oﬀ ﬂux Sc ) is set to somefactor of the confusion noise such that xc = qσ, where the factor q limits the largest response
included in the evaluation of the confusion noise σ (values of ∼5 are assumed in the calculationsof Franceschini et al. (1991)). The diﬀerence in assuming a cut oﬀ in the response as opposed
to a cut oﬀ in ﬂux is that weak contributions from strong sources are not neglected, as evena strong source far from the axis of the beam may contribute signiﬁcantly to the point spreadfunction of the beam.
Assuming for clarity in the calculations, a circular Gaussian beam proﬁle, f (Θ) = f ((θ/θo )2 ) =2e−4(θ/θo ) ln2 such that dΘ = 2θdθ/θo2 where θo is the Full Width at Half Maximum (FWHM) of
the beam, integrating equation A.3.11 over the solid angle of the beam gives;
σc2 (xc ) = πθo2 P (xc )(A.3.12)

or


θo =

σc2
,
πP (xc )(A.3.13)

where P , eﬀectively the power in the ﬂuctuations, is given by,
 xc

P (xc ) =0x2 dx.n(x/f (Θ))e4Θln2 dΘ,(A.3.14)

Thus the confusion limit can be directly related to the FWHM of the instrument beam. For
the simple power law representation of the distribution of extragalactic sources given previously
and the deﬁnitions of equations A.3.9 and A.3.10, the confusion limit is given by,σc =kΩe (3−γ)/2
x
,
3−γ c

(A.3.15)

Therefore, the confusion noise limit will be a complex function of the beam size θo , the source
counts N (S), the cut oﬀ in ﬂux Sc or response xc and the factor q. For the assumed symmetric(2/(γ−1)).
Gaussian beam proﬁle, σc ∝ θoNote that the beam size θ (or FWHM, θo ) ∝ λ/D where λ is the observation wavelengthand D is the telescope diameter. Therefore the confusion due to faint galaxies will be worse atlonger wavelengths and smaller telescope diameters. Since the confusion noise is related to themean number of responses (the source density) and the cut oﬀ response q/xc , a useful, practicalbenchmark for the confusion limit can be set by limiting the number sources per beam before thebeam becomes confused. Ideally the confusion limit would be determined by the ﬂux at which
the source density becomes 1 source per beam although more realistically a limit of between
1/20-1/50 sources per beam (20-50 beams per source) is assumed (e.g. Hogg 2001).
In Table A.3.3 and Table A.3.4 the source confusion limits for the ASTRO-F FIS and IRC
wavebands respectively are tabulated assuming a confusion threshold of 20 beams per source
and a concordance (i.e. ﬂat, dark energy dominated) cosmological world model (Ho = 72
km s−1 Mpc−1 , Ω = 0.3, Λ = 0.7) for the evolutionary model of Pearson (2005). These models
are consistent with the estimations of the confusion limits from the early release Spitzer results(Dole et al. 2004). From Table A.3.3 it is clear that source confusion will not be a problem for
the All Sky Survey but may become critical for slow scan (pointed) observations with the FIS.
In the case of the IRC, the longest wavelengths (15–24 µm) will suﬀer from source confusion forobservations with pointings of the order of a few.

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Table A.3.3: Estimated confusion limits (N.B. in mJy) due to point sources for ASTRO-F FIS
bands.
Band
N60
WIDE-S
WIDE-LN160

5σ Sensitivity (mJy)
All Sky Survey Slow Scan Mode
600
33
200
11
400
8.3
800
17

Source Confusion Limit (mJy)
20 beam/source
3
7
45
50

Table A.3.4: Estimated confusion limits (N.B. in µJy) due to point sources for ASTRO-F IRC
bands.
Band
N2
N3
N4
S7
S9W
S11
L15
L18W
L24

Sensitivity (600s) (µJy)
5σ (µJy)
7.9
3.7
7.2
32.8
26.1
37.1
68.4
87.3
180.8

Source confusion limit (µJy)
20 beam/source
<1
<1
<1
<1
<1
2
16
66
134

REFERENCES
Condon, J.J., 1974, ApJ 188, 279
Dole H. et al., 2004, ApJSS, 154, 93
Franceschini, A., Toﬀolatti, L., Danese, L., De Zotti, D., 1989, ApJ 344, 35
Franceschini, A., Toﬀolatti, L., Mazzei, P., Danese, L., De Zotti, D., 1991, A&ASS 89, 285
Helou G., BeichmanC.A., 1990, The confusion limits to the sensitivity of sub-mm telescopes, In: From
Ground-Based to Space-Borne Sub-mm Astronomy, Proc. of the 29th Liege International Astrophysical Coll., ESA pub. 117
Gautier, T.N., Boulanger, F., Perault M., Puget J.L., 1992, AJ 103, 1313
Herbstmeier, U. et al., 1998, A&A 332, 739
Hogg, D., 2001 AJ, 121, 1207
Jeong, W.S., et al., 2005, MNRAS 357, 535
Kiss, Cs., Abraham, P., Klaas, U., Juvela, M., Lemke, D., 2001, A&A 379, 1161
Lagache, G., & Puget, J.L., 2000, A&A 355, 17
Low, F.G., et al., 1984 ApJ, 278, L19
Matsuhara, H. et al., 2000, A&A 361, 407
Pearson C.P., 2005, MNRAS 358, 141
Scheuer, P.A., 1957, Proc. Cambridge Phil. Soc. 53, 764
Takeuchi T.T., Ishii T.T., 2004, ApJ, 604, 40
Vaisanen, P., Tollestrup, E.V., Fazio, G.G., 2001, MNRAS 325, 1241

Appendix B

ASTRO-F Cookbook for Open Time
Observations
B.1

Introduction

In order to aid observers with the preparation of proposals for Open Time, we present in thissection two worked examples for observation and submission of Open Time proposals with
ASTRO-F. We show
1- Imaging of extragalactic ﬁelds with IRC + FIS
2- Spectroscopy of a galactic target with IRC
In general the preparation of Open Time proposals follows the algorithm below using dedicated tools available via a web interface. In addition we expect users to iterate their observations
using the tools before their ﬁnal submission in order to obtain the best possible observation
strategy.
1- Target selection and visibility check2- Preparation of the ”Target list”
3- Validation check of the ”Target list”
4- Run Duplication Check Tool
5- Run Instrument Performance Tool (IPT)
6- Submission via Submission Tool115

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Table B.2.1: Extragalactic target ﬁelds for observations with ASTRO-F
Name
ISO FIRBACK MARANO
UKIDSS-UDS
SST GOODS-SVF
HST-N

B.2

Equinox
J 2000J 2000J 2000J 2000R.A.
03:11:00.0
02:21:20.0
16:09:52.0
12:36:49.9

Dec.
−54:45:00.0−04:30:00.0+54:55:00.0+62:12:58.0

Example 1: Imaging of extragalactic fields with the IRC
and FISIn this worked example we consider a program to image several deep extragalactic ﬁelds with
the ASTRO-F IRC & FIS instruments. We will simulate the preparation of the proposal fromscientiﬁc background to proposal submission.

B.2.1

Scientific BackgroundStudies with the Infrared Space Observatory (ISO), of the Hubble Deep Fields have revealed
star formation rates at least comparable to, or higher than those of optical/UV studies (Mannet al. 2002). At submillimetre wavelengths, surveys with SCUBA on the JCMT have revealed alarge (>3000 deg−2 at S850 > 2 mJy) strongly evolving population of sources with bolometricluminosities > 1012 Lo and star formation rates of ∼ 300–1000 M yr−1 with a median redshiftof 2.4 (Chapman et al. 2003). The star formation rate at z∼1 requires signiﬁcant evolution in
the IR galaxy population from the current epoch. Deep observations with the Spitzer Space
Telescope have conﬁrmed this strong evolution in the galaxy population out to ISO redshifts
and furthermore provided insight into the higher redshift Universe in the so called redshiftdesert z∼ 1–3 (Papovich et al. 2004). To connect the local and intermediate redshift IRAS/ISO
Universe to the higher and high z Universe observed by Spitzer and SCUBA, comprehensive
multiwavelength imaging is required throughout the extragalactic populations. Observations
with ASTRO-F can observe the full infrared wavelength range in 13 bands from 2–180 µm, ﬁll
in the area in the Spitzer wavelength desert between 8–24 µm to much higher sensitivities thanobtained in the previous ISO surveys (Elbaz et al 2002). Our targets are the sources responsible
for the strong evolution seen in the ISO counts. Although at fainter ﬂuxes the sources responsible
for the counts are thought to be luminous infrared galaxies (LIRGs) the identity of the sources atbrighter ﬂuxes at the upturn in the counts are not so clear. Depending on the various ﬂavours of
evolutionary scenario, they could be M82-like starburst galaxies or more luminous ARP 220-likeultraluminous galaxies (ULIRGs). We would like to test the latter possibility. Therefore ourtargets will be any ULIRG sources lying in the redshift range 0.5–1 since the redshift distributionof the ISO sources has a median value at around z∼ 0.8. For our program we will chose a set ofwell known ﬁelds where there is already a lot of follow up data. These are the ISO-FIRBACK
ﬁeld, UKIDSS-UDS ﬁeld, the SST-GOODS-SVF ﬁeld and the HST-N ﬁeld. These ﬁelds are
summarized in Table B.2.1.

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117

Selection of Target Field and Target Visibility

A Visibility Tool is provided that gives an approximate estimation of the open time visibility ofthe target ﬁeld during the Phase 2 of the ASTRO-F mission. Note that only 30% of the total
number of pointings with ASTRO-F is reserved for open time which translates roughly into 6
orbits every 2 days or almost a 2 degree interval on the ecliptic longitude. In addition to thisany orbits not scheduled for the guaranteed time programs are available for scheduling for open
time. However, visibility for speciﬁc targets is quite heavily constrained and an extremely strong
function of ecliptic latitude. The visibility given by this tool is calculated under the following
assumptions.
1- The instrument is of 10 arcmin width and at the boresight.
2- The nominal scan path along the great circle is considered with ±1 degree oﬀset control.
3- The Moon and SAA orbits have been removed.
4- The maximum number of pointed observations at any single pointing has a ceiling set at50.
Note that ASTRO-F has the capability to observe ±1 degree either side of the canonical scanpath (known as the oﬀset control), thus there is some limited ﬂexibility in the available numberof pointings per target although this number cannot be guaranteed in practice. Therefore thenumbers given by the visibility tool should be considered a guideline as to the absolute maximum
observing opportunities for any given target.

Figure B.2.1: The Visibility Tool user front end.
In general, if the number of targets we want to observe is small we can manually input theseco-ordinates into the Visibility Tool, alternatively we can provide the full target list as a batch
ﬁle: preparation of this list is explained in the next section. Figure B.2.1 shows the batch &interactive query page for the Visibility Tool with one of our ﬁelds as an entry. We have selected
equatorial co-ordinates in this case. Alternatively, selecting batch mode we will be prompted fora target List ﬁle to upload. The results of the Visibility Tool for all our ﬁelds (i.e. using batchmode) are shown in Figure B.2.2. The input data from Table B.2.1 is shown on the right hand

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Figure B.2.2: Results from the Visibility Tool. The input data is shown on the right handside and the output on the left. The output lists the co-ordinates in ecliptic format which ismore suitable for the ASTRO-F orbit and the number of pointings available in Phase 2a (ﬁrst
6 months after Phase 1), Phase 2b (4 months following Phase 2a) and the total number ofpointings available for open time observations in the entire Phase 2 of the mission. The ‘+’ &‘-’ signs refer to the scan direction: corresponding to the FoV moving from North to South in
Ecliptic coordinates for the ‘+’ case and in the opposite direction for the ‘-’ case.
side and the output on the left.
The output lists the query input on the right hand side and on the left, the target co-ordinatesin the ecliptic system, which is more suitable for the interpretation of the visibility and the total
number of pointings available for open time in Phase 2 of the mission. The number of pointings
are also broken down into categories of total number of orbits available in Phase 2a (ﬁrst 6
months of Phase 2) and Phase 2b (next 4 months of Phase 2). The ‘+’ & ‘-’ signs refer to thescan direction: corresponding to the FoV moving from North to South in Ecliptic coordinates
for the ‘+’ case and in the opposite direction for the ‘-’ case.
The visibility (i.e. number of times the ﬁeld can be observed during phase 2) of our ﬁelds
are 61, 18 ,25 & 26 respectively during Phase 2. At an ecliptic latitude of −17.5, the UDS ﬁeld
has rather low visibility with ASTRO-F and it may be unlikely that we will be able to carry out
all the pointings with the FIS and IRC instruments as desired. Therefore, we shall drop thisﬁeld and concentrate on the remaining three ﬁelds.
The output from the visibility tool should be saved as an ASCII ﬁle as it will need to be
uploaded as part of the ﬁnal submission process for our proposal.

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119

Preparation of Target List

The objective of the IRC observations is to detect the emission from the mid-infrared dust
features (PAH features in dusty ULIRGs to z ∼ 1). Thus to fully sample these features overthis redshift range we will require as many bands as possible. The IRC has 3 AOTs for imaging,
IRC00 is for deep imaging in a single band and requires 3 pointed observations to ensure adequate
dithering (moreover 2 bands are recommended so the total number of pointings required will
increase to 6). Since we are more interested in spectral coverage we should choose either IRC02
or IRC03. IRC03 will give us the maximum number of ﬁlters (all of them) so we select this.
IRC03 has only one micro-scan for each band and therefore we need to observe the same ﬁeld
twice for suﬃcient redundancy. We conservatively set the number of pointings for each IRC
AOT to 3, which will will check and conﬁrm at a later stage with the Instrument PerformanceTool.
For the FIS, we wish to observe the emission from the dust hump from the ULIRGs in ourtarget ﬁelds out to a redshift of around unity. We need to cover the IRC ﬁeld of view (FoV)
described above (10 arcmin × 10 arcmin), therefore we do not require a particularly long scanbut rather need to scan an area somewhat wider than the FIS detector array. Therefore weselect AOT FIS01 over FIS02 since the former gives us the option of a cross scan shift. The
FIS01 AOT has several parameters that we need to set. The ﬁrst is the scan speed and thesecond is the cross-scan shift. Furthermore, setting these parameters determines the scan length.
AOT FIS01 includes 2 round trip scans with a cross-scan shift. Therefore there will be unevencoverage of our FoV with the central part being covered by the full 2 round trip scans while
the edges will be covered by a √single round trip scan only, thus the central region will have abetter sensitivity by a factor of 2. The cross-scan shift can be either 70 or 240 arcsec, a smaller
step will increase the cross-scan area covered on both scan trips (7–4 arcmin respectively) but
decrease the overall area covered by the scan. We need to cover the 10 arcmin area of the IRC
FoV so we choose a scan step of 70 arcsec with a speed of 15 arcsec/s (resulting in a scan lengthof 20 arcmin and a cross scan coverage of 7 arcmin) and make 2 slightly overlapping pointings
with FIS01 to completely cover the IRC FoV with the FIS. A tool (iris sky) is provided to aid
us with the positioning of the instruments over our ﬁelds. With iris sky we can project our
instrument arrays onto our target ﬁelds (on images from IRAS and 2MASS). With iris sky it isalso possible to create our target list in its’ entirety although we shall construct ours manuallyin this case. Figure B.2.3 shows how to input the ﬁeld and AOT details into the iris sky tooland Figure B.2.4 shows the output with the array ﬁeld of view with from top left to bottom
right, the NIR/MIR-S camera, MIR-L camera and the 2 FIS scans respectively overlaid on thetarget ﬁeld.
Once the AOTs have been selected we can go ahead and prepare the target list for ourobservations. An ASTRO-F target list contains 1 line per target per AOT per ﬁeld of view.
We are observing our target ﬁelds with the IRC MIR-L camera, the IRC NIR/MIR-S camera
and twice with the FIS, therefore we have 4 lines in our target list per target ﬁeld. We have 3
target ﬁelds, therefore we have a total of 12 lines in our target list. The target list is shown in
Table B.2.2. The target list consists of a unique I.D. in ascending but not necessarily contiguousorder, followed by some identiﬁer for the target name which should if possible be a resolvable
by systems such as SIMBAD, NED. The next 3 ﬁelds are reserved for the coordinate systemand target coordinates. For both the IRC03 AOT and the FIS01 AOT, the target positionis at the centre of the FoV (note that for the FIS02 AOT, it would be at the beginning of
the scan). Next we need the AOT, which will be either IRC03 or FIS01. Following the AOT
are the speciﬁc parameters for each given AOT. For the IRC imaging, we have to specify 2parameters, the ﬁlter combination and the camera with which we want to observe. For the
IRC03 AOT the ﬁlter combination is ﬁxed - all of them - so we can ﬁx the ﬁrst parameter as a.

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Figure B.2.3: The steps to set up the iris sky tool. Selecting new entry opens a dialogue box.
The target position can then be entered as well as the AOT details.
The second parameter is either L or N depending on whether we are observing with the MIR-L
or NIR/MIR-S respectively. Since we are observing in all 9 bands we have 2 lines in our target
list, one for each IRC FoV. For the FIS, it is a little more complex, we have 3 parameters to set
for the AOT, the reset interval, scan speed and the cross-scan step. We have already discussed
the latter pair of parameters and have set them to 15 arcsec/s and 70 arcsec respectively. The
choice of reset interval (0.5, 1 or 2s) depends on what we are observing. In general a longer
reset interval improves sensitivity but increases the chance of saturation. For our case, we do
not expect extremely bright sources and would like to get right down to the galaxy confusion
limit, hence we select the longest reset interval of 2s.
In the following ﬁeld we input the number of pointings for each AOT, for our case thiswill be 3 for each IRC03 AOT and 1 for each FIS01 scan. We then have to prioritize ourtargets. ASTRO-F Open Time Observations are divided into 3 priorities, A , B, C . Our A& B targets, barring any unforeseen events, should be observed, while the priority C targets
should be considered as back-ups. In our target list, the number of priority A , B, C targets
should be around the same. For our observations we prioritize by target ﬁeld although we could
just as easily prioritize the FIS observations over the IRC observations for all our ﬁelds insteadfor example. We choose to prioritize our ﬁelds by their visibility ﬁrst and then by possible
duplication conﬂict (see the discussion on duplication later) with the HST ﬁeld as our back uppriority C target due to a duplication conﬂict.
ASTRO-F supports the concept of “symbiotic observation”. This option in the penultimate
ﬁeld of the target list entry (ﬁeld 10) enables the observers to specify that more than 2 pointings

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121

Figure B.2.4: Using the iris sky tool to overlay the arrays and FoV of the instruments on ourtarget ﬁelds. From top left to bottom right with the NIR/MIR-S camera, MIR-L camera andthe 2 FIS scans respectively overlaid on the target ﬁeld.

are required to complete an observation. This may be set, for example, for observations at thesame position but with diﬀerent instruments, or for mapping large areas. When this parameteris given, however, it will constrain the observation plan signiﬁcantly and increases the risk thatthe observations will not be scheduled at all. If we schedule any of our observations as symbiotic
we would group them together with a Group ID number and an observational style as “1;o” .
The second parameter being either o or a , where o means one of the observation is acceptable
and a means either all or none is acceptable. For our targets, we could still get useful information

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Table B.2.2: Target List for submission
1, ”IRC-ISO”, ”J2000”, 03:11:00.0, -54:45:00.0, ”IRC03”, ”a;L”, 3, A, , ”ISO FIRBACK MARANO”
2, ”IRC-ISO”, ”J2000”, 03:11:00.0, -54:45:00.0, ”IRC03”, ”a;N”, 3, A, , ”ISO FIRBACK MARANO”
3, ”FIS-ISO”, ”J2000”, 03:10:53.73, -54:45:45.78, ”FIS01”, ”2.0;15;70”, 2, A, , ”ISO FIRBACK MARANO”
4, ”FIS-ISO”, ”J2000”, 03:11:05.7, -54:44:09.5, ”FIS01”, ”2.0;15;70”, 2, A, , ”ISO FIRBACK MARANO”
5, ”IRC-SST”, ”J2000”, 16:09:52.0, +54:55:00.0, ”IRC03”, ”a;L”, 3, B, , ”SST GOODS-SVF”
6, ”IRC-SST”, ”J2000”, 16:09:52.0, +54:55:00.0, ”IRC03”, ”a;N”, 3, B, , ”SST GOODS-SVF”
7, ”FIS-SST”, ”J2000”, 16:09:47.01, +54:55:33.08, ”FIS01”, ”2.0;15;70”, 2, B, , ”SST GOODS-SVF”
8, ”FIS-SST”, ”J2000”, 16:09:56.88, +54:54:27.54, ”FIS01”, ”2.0;15;70”, 2, B, , ”SST GOODS-SVF”
9, ”IRC-HST”, ”J2000”, 12:36:49.9, +62:12:58.0, ”IRC03”, ”a;L”, 3, C, , ”HST-N”
10, ”IRC-HST”, ”J2000”, 12:36:49.9, +62:12:58.0, ”IRC03”, ”a;N”, 3, C, , ”HST-N”
11, ”FIS-HST”, ”J2000”, 12:36:40.3, +62:14:08.7, ”FIS01”, ”2.0;15;70”, 2, C, , ”HST-N”
12, ”FIS-HST”, ”J2000”, 12:36:59.4, +62:11:47.3, ”FIS01”, ”2.0;15;70”, 2, C, , ”HST-N”

even with a partial observation so we do not want the extra constraint of symbiotic observation,
therefore we will leave the ﬁeld blank (not forgetting the ”comma” ﬁeld separator ). Finally,
any notes are appended on the end of each target list line. Our ﬁnal target list is shown in
Table B.2.2.

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B.2.4

123

Target List Validation

Once the target list has been completed, we can proceed towards the submission stage. However
there are still a few steps to complete. We should run our target list through both the Target
List Validation Tool and then we need to submit the target list to the observation Duplication
Check Tool.
The Target List Validation Tool takes a target list and makes a simple check on the target
list format, e.g. correct number of ﬁelds, AOT in expected form etc. It does not perform any
test on the scientiﬁc validation, e.g. on position or sensitivity etc. We can submit the Target
List as a batch ﬁle in a similar manner to the Visibility Tool as in Figure B.2.5, in fact this(target list) format is accepted by all tools during the submission process. Figure B.2.6 shows
the output of the Target List Validation Tool. Any errors would be highlighted in red and asummary given at the bottom of the output.
The target list will be automatically validated when we upload it during the ﬁnal proposal
submission stage so there is no need for us to save the output of the Validation Check Tool.

Figure B.2.5: Input to the Target List Validity Checker

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Figure B.2.6: Output from the Target List Validation Tool. Any Errors would have beenmarked in red and the number of errors would have been listed at the bottom of the output in
the Pointings summary.

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125

Duplication Check Tool

Duplicate observations are decided on a basis of position, instrument/AOT and the number ofpointings. Figure B.2.7 shows the input screen for the Duplication Check Tool. Full results orsummary output can be selected, where selecting full results would give details of the conﬂicting
guaranteed time (Mission Programmes) observations for each duplicated observation in ourtarget list. Putting our target list into the Duplication Check Tool and selecting summary
for the out put, we ﬁnd that indeed there is a conﬂict of position or instrument for 5 of ourobservations in 2 of our ﬁelds (see Figure B.2.8), however since the AOT is diﬀerent it is notstrictly a duplicate observation, the TAC will have to look into the scientiﬁc justiﬁcation todecide whether our observation is unique.
The output from the Duplication Check tool should be saved as an ASCII ﬁle as it will need
to be uploaded as part of the ﬁnal submission process for our proposal.

Figure B.2.7: Input screen for the Duplication Check Tool. The input is the target list as usual.
Full results or summary output can be selected.

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Figure B.2.8: Results from the Duplication Checker Tool. We see that 5 of our observations
are in conﬂict with the guaranteed time on ASTRO-F. The green denotes a conﬂict in position
whilst the red denotes a conﬂict in position and instrument. Duplication can be made on
position, instrument, AOT and the number of pointings.

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127

Table B.2.3: Observation Details ﬁle for the Instrument Performance Tool. Columns are ID
number, Target name, Band, ﬂux in mJy, signal/noise, number of pointings (blank) and type ofsource point/diﬀuse/line.
1,”IRC-ISO”, L15, 90,5., ,P
1,”IRC-ISO”, L18W, 200,5., ,P
1,”IRC-ISO”, L24, 300,5., ,P
2,”IRC-ISO”, N2, 5,5., ,P
2,”IRC-ISO”, N3, 5,5., ,P
2,”IRC-ISO”, N4, 10,5., ,P
2,”IRC-ISO”, S7, 50,5., ,P
2,”IRC-ISO”, S9W, 50,5., ,P
2,”IRC-ISO”, S11, 50,5., ,P

B.2.6

Estimation of Sensitivity for IRC

The typical ﬂux densities for the colder ULIRGs in the mid-infrared are around one to a few 100µJy in the 15–24 µm range dropping to 10’s µJy in the 7–12 µm range (see Figure B.2.11). A
performance estimator tool is available to calculate the number of pointings required for a given
speciﬁc AOT, target sensitivity and desired signal to noise. The Instrument Performance Tool
requires the Target List and an Observation Details ﬁle to be combined as in Figure B.2.9. The
Observation Details ﬁle includes an identiﬁer for each separate AOT, the target name, AOT,
band, required sensitivity in micro-Jy, signal to noise, number of pointings and whether we areobserving a point, diﬀuse source, lines etc. One of either, sensitivity, signal/noise, number ofpointings is omitted in our input list and this is what will be calculated by the Instrument
Performance Tool. The Observation Details ﬁle for the Instrument Performance Tool, for the
FIRBACK-Marano ﬁeld is shown in Table B.2.3.
Putting these values into the Instrument Performance Tool, the output in Figure B.2.10 isreturned and from this we ﬁnd that we would require around 3 pointings in each band of the IRC
using AOT IRC03. Note that the ﬁeld that has been calculated by the IPT is marked with anasterisk for clarity. The output from the Instrument Performance Tool also gives us a warning
for the L18W band that we require more than 1 single pointing to complete this AOT. The 3
pointing 5 σ detection limits are overlaid on a typical z ∼ 0.8 cold ULIRG in Figure B.2.11. Notforgetting that the IRC MIR-L does not share the same ﬁeld of view as the IRC NIR/MIR-S,
this means that to cover our chosen target ﬁeld we require 3 pointings × 2 FoV = 6 pointings
in total with the IRC.

B.2.7

Estimation of Sensitivity for FIS

For the FIS, we wish to observe the emission from the dust hump from the ULIRGs in our target
ﬁelds out to a redshift of around unity. For this we need to cover the IRC ﬁeld of view describedin the previous section (10 arcmin × 10 arcmin) down to a sensitivity of around 20 mJy and
50 mJy in the short and long wavelength FIS bands respectively. The Observation Details ﬁlefor the Instrument Performance Tool, for the FIRBACK-Marano ﬁeld is shown in Table B.2.4
assuming a single pointing in FIS01. From the Instrument Performance Tool (Figure B.2.12)
we see that we can obtain our desired sensitivity (except in the N60 band but we shall live withthat) in a single pointing of the FIS under FIS01 with a scan speed of 15 arcsec/s. We also geta warning for the 2 wide bands, notifying us that we are dangerously near the confusion limit
for point sources and diﬀuse cirrus emission for these observations. Note that since we input the
number of pointings as a parameter for the FIS observations, the Instrument Performance Toolhas marked the calculated ﬂux with an asterisk. The resulting sensitivity from the Instrument

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Figure B.2.9: Input screen for the Instrument Performance Tool. Two ﬁles are required: thetarget list and an observation details ﬁle.
Performance Tool is shown overlaid on a typical z ∼ 0.8 cold ULIRG in Figure B.2.11.
The output from the Instrument Performance Tool should be saved as an ASCII ﬁle as it
will need to be uploaded as part of the ﬁnal submission process for our proposal.

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Figure B.2.10: Output from the Instrument Performance Tool for the IRC input ﬁleTable B.2.4: Observation Details ﬁle for the Performance Estimator Tool. Columns are ID
number, Target name, Band, ﬂux in mJy (blank), signal/noise, number of pointings and typeof source point/diﬀuse/line
3,”FIS-ISO”, N60, ,5.,1 ,P
3,”FIS-ISO”, WIDE-S, ,5.,1 ,P
3,”FIS-ISO”, WIDE-L, ,5.,1 ,P
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100
ULIG LIR=1012Lo
10

Flux {mJy}

1
0.1
0.01
0.001

ULIG z=0.8
IRC AOT-03 (Ecliptic Pole, 3 pointing, 5 sigma)
FIS AOT-01 (slow scan, 15arcsec/s)

0.0001
1

10

100

103

wavelength {µm} (observed frame)

Figure B.2.11: Estimated required sensitivities to detect dusty cold ULIRGs to redshift 0.8. We
assume 3 pointings with AOT IRC03, i.e. observations in all 9 ﬁlter bands and single FIS01
observation utilizing a scan speed of 15arcsec/s

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Submission of Proposal

After the duplication check, we can ﬁnally submit our proposal via the web interface. The
proposal submission has several stages and we will need to upload various output ﬁles saved
from the tools we used earlier. Each proposal has a 5 character abbreviation which will be
used to identify each proposal. The next step is to upload the Target List which will be automatically veriﬁed (syntax only) on uploading. Following this we will be required to upload theresults we saved from the output of the Visibility Tool, Duplication Check Tool and Instrument
Performance Tool respectively. Once all these ﬁles have been uploaded we will have to enter
the speciﬁc proposal information (name, address, telephone number, etc.) and general proposal
information (Co-Is, category, title, abstract). The Scientiﬁc Justiﬁcation must be submitted
in PDF format and should not exceed 6 pages in total including scientiﬁc rationale, objective,references, ﬁgures and tables, and technical feasibility of the proposal. A conﬁrmation page ofthe submission is created automatically. Shortly after, users will receive a formal conﬁrmationby e-mail with the cover page of their proposal.

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B.3

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Example 2: Spectroscopy of Planetary Nebulae with the
IRC

In this worked example we consider a program to perform spectroscopy on a few Planetary
Nebulae with the ASTRO-F IRC instrument. We will simulate the preparation of the proposal
from scientiﬁc background to proposal submission.

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Table B.3.5: Suitable Planetary Nebula for pilot spectroscopic study with ASTRO-F
Name
NGC7027
SMP83

B.3.1

Equinox
J 2000J 2000R.A.
21:07:01.7
05:36:21.0

Dec.
+42:14:11.0−67:18:14.0

Scientific BackgroundWe consider 2 planetary nebulae as a pilot study for spectroscopy with ASTRO-F. These objects
have extremely rich and highly ionized spectrum due to their hot central star. Analysis of thespectra can give insight into the evolution and origin of the ionizing central stars. Our ﬁrst
target is NGC7027, a very well known, bright planetary nebula with a rich spectrum that has
previously been observed by ISO (Bernard-Salas et al., 2001). The second target is the fainter
SMP83 planetary nebula in the Large Magellanic Cloud, also observed by Spitzer (Bernard-Salas
et al., 2004). We will perform both high resolution spectroscopy in the near-infrared range whichis unique to ASTRO-F, since Spitzer has no capability in this wavelength range. In addition
we shall carry out lower resolution spectroscopy over the entire wavelength range of the IRC inorder to obtain the dust spectrum of our sources. The target positions are given in Table B.3.5.

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B.3.2

Target Visibility

Using the Visibility Tool to check the availability of our targets (entering the co-ordinates using
the interface in FigB.2.1) we ﬁnd that both targets have reasonable visibility. The result of the
Visibility Tool is shown in Figure B.3.13. The visibility of SMP83 is extremely high with 86
observable opportunities due to its location almost on the South Ecliptic Pole where as NGC7027
at a more moderate latitude is only visible on 15 orbits. This emphasizes the unique visibility
constraints of the ASTRO-F Open Time observations.
The output from the visibility tool should be saved as an ASCII ﬁle as it will need to be
uploaded as part of the ﬁnal submission process for our proposal.

Figure B.3.13: Results from the Visibility Tool. The input data is shown on the right handside and the output on the left. The output lists the co-ordinates in ecliptic format which ismore suitable for the ASTRO-F orbit and the number of pointings available in Phase 2a (ﬁrst
6 months after Phase 1), Phase 2b (4 months following Phase 2a) and the total number ofpointings available for open time observations in the entire Phase 2 of the mission. The ‘+’ &‘-’ signs refer to the scan direction: corresponding to the FoV moving from North to South in
Ecliptic coordinates for the ‘+’ case and in the opposite direction for the ‘-’ case. Note that
our second target SMP83 has an extremely high visibility of over 86 due to its location almoston the South Ecliptic Pole at an ecliptic latitude of −87 degrees while NGC7027 has a more
moderate visibility of 15 at an ecliptic latitude of 55 degrees. This result emphasises the strong
dependence of the visibility of ASTRO-F observations on ecliptic latitude.

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Table B.3.6: Target List for submission
1,”NGC7027”,”J2000”, 21:07:01.7, +42:14:11.0,”IRC04”,”b;Nh”,2,A,””,”NGC7027 high resolution”
2,”NGC7027”,”J2000”, 21:07:01.7, +42:14:11.0,”IRC04”,”a;Ns”,2,B,””,”NGC7027 lower resolution”
3,”NGC7027”,”J2000”, 21:07:01.7, +42:14:11.0,”IRC04”,”a;Ls”,2,C,””,”NGC7027 lower resolution”
4,”SMP83”,”J2000”, 05:36:21.0, -67:18:14.0,”IRC04”,”b;Nh”,2,A,””,”SMP83 high resolution”
5,”SMP83”,”J2000”, 05:36:21.0, -67:18:14.0,”IRC04”,”a;Ns”,2,B,””,”SMP83 lower resolution”
6,”SMP83”,”J2000”, 05:36:21.0, -67:18:14.0,”IRC04”,”a;Ls”,2,C,””,”SMP83 lower resolution”

B.3.3

Preparation of Target List

We intend to carry out high resolution spectroscopy with the IRC grism (NG) in the nearinfrared as our ﬁrst priority. Following this, as our second priority, we will observe the entire
dust spectrum over the full wavelength range of the IRC. Note that unfortunately, due todegradation, the LG1 channel is not available for spectroscopic observations so we will have a
gap of a few microns in the mid-infrared. The spectroscopic IRC AOT is IRC04. This AOT has
2 parameters, the ﬁrst is the ﬁlter set and the second is the position of the slit. The NIR camera
has 3 slits. A common slit with the IRC MIR-S camera for spectroscopy of diﬀuse sources, anentrance aperture for point source spectroscopy and another slit for high resolution spectroscopy
with the NG element. In addition, the MIR-L camera also has a slit for spectroscopy of diﬀuse
sources. Since we wish to carry out both high resolution spectroscopy with the NG element
and lower resolution spectroscopy with the NP/SG1/SG2 elements, plus the longer wavelength
coverage provided by the LG2 element on the MIR-L camera, we require at least 3 pointing per
ﬁeld of view. In addition since no dithering is carried out during this AOT, we must observe
each sky area at least twice for data redundancy. Therefore we are looking at a minimum of
6 pointings per target and 3 lines per target in the Target List. Later we shall see with theInstrument Performance Tool that our fainter target (the high visibility SMP83) requires 2
pointings per ﬁeld of view per AOT. The target list is shown in Table B.3.6. As in the previous
example, with have a unique I.D. in ascending but not necessarily contiguous order, followed bysome identiﬁer for the target name which in the case of many of our targets should resolvable
the NED/SIMBAD systems. The IRC04 AOT has 2 parameters to set, the ﬁlter set and the slit
position. The ﬁlter set is determined by whether we wish to use the NP (ﬁlter set a) or NG (ﬁlter
set b) element of the NIR camera. We can use the iris sky tool to visualize our observations andFigure B.3.14 shows the iris sky tool with the outmost slit of the IRC camera positioned over
the target (NGC7027). In this ﬁgure, the entrance aperture for point sources and the common
inner slit for diﬀuse sources.
Next we have to choose the slit parameter. For the high resolution spectroscopy, this parameter is Nh : the target position is placed at the high resolution slit of the NIR camera. For the
lower resolution spectroscopy, the parameter is Ns & Ls for the NIR/MIR-S and MIR-L cameras
respectively. We set the number of pointings as 2 for both the brighter source and for our faintersource which we will conﬁrm with the Instrument Performance Tool later. For this example,
we prioritize our targets by the high and low resolution spectroscopy, giving the high resolution
observations with NG a priority A and the lower resolution spectroscopy with NIR/MIR-S andMIR-L priorities B & C respectively. Again, as in the previous example we choose not to link
our observations together as symbiotic since we can still get decent scientiﬁc results with either
the high or low resolution spectroscopy.

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Figure B.3.14: Using the iris sky tool to overlay the IRC NIR high resolution spectroscopy slit
over the Planetary Nebula NGC7027.

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137

Target List Validation

The Target List syntax is validated in an identical manner to the previous imaging example.
The target list will also be automatically validated when we upload it during the ﬁnal proposal
submission stage so there is no need for us to save the output of the Validation Check Tool at
this stage.

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B.3.5

Duplication Checker

Duplicate observations are decided on a basis of position, instrument/AOT and the numberof pointings. Putting our target list into the duplication checker we ﬁnd that indeed there isa conﬂict of position for one of our ﬁelds (see Figure B.3.15). Interestingly, the high latitude
target, SMP83, actually in the LMC has escaped conﬂict with the ASTRO-F LMC LS guaranteed
program. The output from the Duplication Check tool should be saved as an ASCII ﬁle as it
will need to be uploaded as part of the ﬁnal submission process for our proposal.

Figure B.3.15: Results from the Duplication Checker Tool. We see that one of our targets
are in conﬂict with the guaranteed time on ASTRO-F. Duplication can be made on position,
instrument/AOT and the number of pointings.

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Table B.3.7: SMP83 line strengths
Wavelength (microns)
7.652
10.52
12.82
13.11
14.33
15.55
18.73
24.33
25.89
33.46
34.84

Line
[Ne vi][S iv][Ne ii][Ar v][Ne v][Ne iii][S iii][Ne v][O iv][S iii][Si ii]Flux 10−17 W/m2
8.53
14.7
1.55
1.23
14.2
10.3
3.92
15.0
43.7
3.47
2.76

Table B.3.8: Observation Details ﬁle for the Performance Estimator Tool. Columns are ID
number, Target name, Band, ﬂux in W m−2 (blank), signal/noise, number of pointings and typeof source point/diﬀuse/line
1,”SMP83”,NG, ,5.,1,L
2,”SMP83”,SG2, ,5.,1.,L
3,”SMP83”,LG2, ,5.,1.,L

B.3.6

Estimation of Sensitivity for IRC

For our ﬁrst target (NGC7027) all the lines are brighter than 10−15 W m−2 (Bernard-Salas etal., 2001), therefore this is well above the IRC detection limits and should be easily observablein a single pointing. The second target (SMP83) presents fainter lines of the order of 10−17
W m2 , which are closer to the IRC line detection limits (compare for example the LG2 detection
limit of ∼ 1 × 10−17 W m2 to the S III line given in Table B.3.7 at 18.73 micron of intensity
3.9 × 10−17 W m−2 . For the shorter wavelength lines, as a ﬁrst approximation, we can scale theline ﬂuxes of NGC7027 by the distance to the LMC. Under this assumption, we could expect theH2 line at 2.8 microns to be of the order of ∼ 3.0 × 10−17 W m−2 and the Br alpha line at 4.05
microns to be ∼ 92.0×10−17 W m−2 . Using the Instrument Performance Estimator Tool (Target
List plus Observation Details File) we can estimate the capability of the IRC in spectroscopy
mode. Table B.3.8 shows the Observation Details File to be input into the Performance Tool.
For this example we leave the ﬂux ﬁeld (out of the ﬂux, S/N, number of pointings ﬁelds) blankand initially set the number of pointings to 1. The result is shown in Figure B.3.16 and we cansee the IRC is sensitive enough to detect these lines in SMP83 in a single pointing! Although
we still require 2 pointings for dithering. The output ﬁeld that is calculated by the Instrument
Performance Tool is marked by asterisks for clarity.
The output from the Instrument Performance Tool should be saved as an ASCII ﬁle as it
will need to be uploaded as part of the ﬁnal submission process for our proposal.

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Figure B.3.16: Output from the Instrument Performance Tool for the IRC input ﬁle for spectroscopy. The calculated ﬁelds are marked by asterisks.

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141

Submission of Proposal

After all the tools have been used and the Target List checked, we can ﬁnally submit our proposal
via the web interface in an identical manner as that described for the previous imaging example
in SectionB.2. Remember that we will need to upload the previously saved ASCII ﬁles from theoutputs of the Visibility, Duplication & Instrument Performance Tools.

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Chapman S.C., Blain A.W., Ivison R.J., Smail I.R., 2003, Nature, 422, 695
Elbaz D., Cesarsky C.J., Chanial P., Aussel H., Franceschini A., Fadda D., Chary R.R., 2002, A&A, 384,
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Mann R.G., et al. 2002, MNRAS, 332, 549
Papovich C. et al., 2004, ApJSS, 154, 70
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