1333 and 3623 – two asteroids with large amplitude lightcurves

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Spectroscopy and atmospheric transparency

The observation of celestial bodies using different types of ground-based telescopes is possible in the regions of electromagnetic spectrum for which the atmosphere is transparent. There are two spectral windows which allow the observation: the optical (V) up to the mid-infrared(the near-infrared 0.8 – 2.5m interval is denoted as NIR) and the radio one. The X-rays and ultra-violet wavelengths are blocked due to absorption by ozone and oxygen, while the far infrared radiation is blocked mainly due to absorption by water and carbon dioxide.
While in the optical wavelength region the atmosphere is almost completely transparent, in the near-infrared there are absorption bands of water vapors making some regions like 1.4-1.5 µ m and 1.8-2.0m poorly transparent (Fig. 2.2). Because of the effects of the atmosphere, ob-servations with space telescopes, such as the Hubble and Spitzer telescopes, are very valuable.
Another important difference between the V and NIR spectral intervals is the fact that the sky is brighter in the NIR region. For example in the J, H, K filt ers1 the estimated sky back-ground has 15.7, 13.6, respectively 13 mag/arcsec2. Additional, important variations of the sky background could be observed in the intervals of tens of arc minutes of the sky.
These issues in the NIR part require additional observing techniques and processing methods (described in Chapter 4) comparing to observation in the V part of the spectrum.

Spectroscopy for asteroids

The knowledge of the surface mineralogy of individual asteroids and groups of asteroids can be inferred through the spectroscopy. The solar light reflected from the asteroids contains essential information regarding the optical properties of the materials found at the asteroids surface. The spectral interval 0.8 – 2.5 µm is very important to discriminate between different mineralogy of silicate-based compounds. Silicate minerals identifica tion is based on the presence of broad bands of absorption around 1 and 2 µm. These bands are due essentially to the presence of olivine and pyroxene (or mixtures) on the surface of the asteroid.

IRTF Telescope and the SpeX instrument

Several large telescopes are equipped with a spectrograph. Some examples among those sup-porting research programs for planetary sciences are: the NASA InfraRed Telescope Facility (IRTF), the European Southern Observatory (ESO) Very Large Telescope (VLT), the ESO New Technology Telescope (NTT) and Telescopio Nazionale Galileo (TNG).
The NIR spectra presented in this thesis are obtained with NASA IRTF (Fig. 3.1a), a 3.0-meter telescope located on the top of Mauna Kea – Hawaii. It was built initially to support the Voyager missions, but today at least 50% of the observing time is devoted to planetary sciences. The IRTF hosts 6 facility instruments:, SpeX (Fig. 3.1b), NSFCAM2, CSHELL, MIRSI, Apogee, Moris. These instruments allow imaging, polarimetry, low and high resolution spectroscopy in the near to mid infrared (0.8 – 30)m.
SpeX – the most used instrument by planetologists from NASA IRTF telescope, is a low to medium resolution spectrograph and imager in the (0.8-5.5)m. It provides spectral resolu-tions of R ≈ 1000 – 2000 across 0.8 – 2.4m, 2.0 – 4.1m, and 2.3 – 5.5m, using prism cross-disperser [Rayner et al., 2003]. Single order long slit modes are also available. A high throughput prism mode is provided for 0.8 – 2.5m spectroscopy at R ≈ 100.
SpeX employs a 1024×1024 Aladdin3 InSBb CCD array for acquiring the spectra, while image acquisition could be made with a 512×512 Alladin2 CCD InSb array.
Two interfaces are used to manage the instrument and the spectrograph, GuideDog interface (Fig. A.1) is dedicated to pointing and tracking the object and BigDog (Fig. A.2) interface is used for spectrograph setup and spectra acquisition.
Observations on IRTF can be performed from anywhere in the world using an internet con-nection via VNC (Virtual Network Connection) protocol. The observing runs for this work were conducted remotely from Meudon-Paris (France), more than 12 000 Km away from Hawai [Birlan et al., 2004b, Bus et al., 2002]. Due to different time zones, for the observers in Meudon, the observing time occurred during daylight hours: a full hawaiian night session started at 5 a.m. and ended at 5 p.m. – Paris local time.
Using the equipment provided at Centre d’Observation à Dista nce en Astronomie à Meudon (CODAM), team had the control remotely of both the instrument/guider system and the spec-trograph set-up and spectra acquisition [Birlan et al., 2004a, 2006]. A permanent and constant audio/video link with the telescope operator was essential in order to administrate possible service interruptions, thus another interface was used to keep the audio-video link open (via Polycom ViewStation video-conference system both on Meudon and Mauna Kea). All soft-ware was re-initialized at the beginning of each night.

Planning the observations

The typical cycle of astronomical observations on world-class telescopes imply the following steps: 1) issue received with the call of proposal for observers; 2) targets selection; 3) proposal submission and evaluation; 4) observations; 5) data reductions and analysis; 6) publications and dissemination of the results.
Generally, the targets are selected based on a desired scientific criterion, which in general reduce their number up to few tens. Observational time is obtained after a severe selection of the best proposals made by the IRTF time allocation committee.
Scheduling the observing time for asteroids requires an ephemerides (the position of astro-nomical objects on the sky) calculator, such as: http://ssd.jpl.nasa.gov/horizons. cgi or http://www.imcce.fr. However, for the large observing programs that targets many objects an additional scheduler is required. It is the case of the program Physical prop-erties of low delta-V Near-Earth Asteroids for which I designed a planning software, available online at: http://m4ast.imcce.fr/lowdv.php.
The tool selects targets based on the following criteria:
• « delta-V » – the available propulsion required remains an engineering design constrain; typically « delta-V » should be lower than 7 km/sec for the initial rendezvous and should have additional 1 km/sec for return.
• H – the absolute magnitude, determines the diameter of the target and should be restricted to consider the kilometers size objects.
• the apparent magnitude and proper motion of the object should be selected in agreement with the telescope capabilities.
• the altitude at the moment of the observing time should correspond to a low airmass.
For example, among the objects accessible for observation with IRTF telescope on May 18-19, 2008 were: (5620) Jasonwheeler, ( 1943) Anteros, (143651) 2003 QO104, and (433) Eros.

Data reduction procedures

The data reduction procedures for the observational data consist in obtaining the flux as a func-tion of wavelength from the CCD images. Usually these images are in .fits 1 format.Additional information regarding the CCD images for astronomy can be found in the book « Electronic Imaging in Astronomy Detectors and Instrumentation » [McLean, 2008].
The calibration files are:
Bias – in the « no-signal » condition, the CCD electronics system will always produce a small positive readout signal for each pixel. This electronic signature is therefore known as the bias level. This can be easily measured by taking a zero second exposure time. Multiple bias frames can be averaged to reduce the random readout noise by averaging them.
Dark – Dark-current levels, due to thermal noise, are determined by long exposures with the CCD shutter closed. To minimize this effect CCDs are generally cooled to low temperatures. To remove this noise, an exposure is taken of similar length as the useful im-ages,with the dome and shutter closed. These dark images can also be used to find dead or hot pixels. If dark frames are used, the CCD bias is contained within them and separate bias corrections are not necessary. Similarly, multiple dark exposures can be averaged to reduce the random readout noise by averaging them. This is called a « master dark file ». Dark current is more significant in infrared arrays, and it may not be linear and scalable from different exposures [McLean, 2008].
Flat – Sensitivity variations from pixel to pixel arise as the result of fabrication processes and also due to optical attenuation effects such as microscopic dust particles on the surface of the CCD. A flat field image to correct for this effect is usually obtained by observing inside of the telescope dome (if it is matt or white) or place a huge white card on the dome. The dome is illuminated with a projector lamp. In this case the telescope is completely out of focus which ensures that the field is uniformly illumin ated. For faint objects it is the light of the sky that dominates, and so it is better to try to use the sky itself as a flat-field.
In other cases, as in photometry for instance, the flat field co uld be done using a sky region in the day light time (at the beginning and end of the night).
Arc lamp – « arc » images are used to determine the pixel to wavelength correspondence, more exactly to make the wavelength calibration. Typically, the lamps used contain helium, neon, xenon, argon or a combination thereof. The emission lines from the spectrum of the arc lamp are at known wavelengths and can be identified. For th e IRTF/SpeX a lamp with argon is available (Fig. 3.2). Standard star – A solar-like standard star spectrum taken at similar airmass is required to correct the atmospheric effects and to remove the signature of the Sun’s spectrum in order to have only the signature of the asteroid surface. The G2 stars are used with magnitudes (usually between 5 to 12) that allow to obtain a high SNR (signal to noise ratio) spectrum with a short integration time (a few seconds). If the star is too bright it will saturate the CCD, while a fainter star will require an unacceptably long integration time.

Table of contents :

I INTRODUCTION 
1 Why asteroids? 
1.1 The place of asteroids in the structure of the Solar System
1.2 The Discovery Of Asteroids
1.3 Distribution and diversity of asteroids
1.4 Asteroid brightness and albedo
1.5 My contribution to asteroids discovery
2 Why spectroscopy? 
2.1 Diffraction gratings and prisms
2.2 Spectroscopy and atmospheric transparency
2.3 A simple application
2.4 Spectroscopy for asteroids
2.4.1 Reflectance versus emission
2.4.2 Spectral features
II TECHNIQUES FOR ASTEROID SPECTROSCOPY 
3 Observing techniques 
3.1 IRTF Telescope and the SpeX instrument
3.2 Planning the observations
3.3 Data reduction procedures
4 Spectral analysis techniques 
4.1 Interpretation
4.1.1 Taxonomy
4.1.2 Spectral comparison – Comparative planetology
4.1.3 Space weathering effects
4.1.4 Band parameters
4.2 Algorithms
4.2.1 Taxonomic classification
4.2.2 Curve matching
4.2.3 Computing the space weathering effects
4.2.4 Application of the Cloutis model
5 M4AST – Modeling of Asteroids Spectra 
5.1 Spectral database
5.1.1 Structure of M4AST database
5.1.2 The content
5.1.3 M4AST database via the Virtual Observatory
5.2 The interface
5.2.1 Database interface
5.2.2 Modeling tool interface
5.2.3 Updating the database
5.3 Testing of M4AST
5.3.1 Results
5.3.2 Discussions regarding misinterpretations of spectra
III OBSERVATIONS AND RESULTS 
6 Spectral properties of near-Earth asteroids 
6.1 Log of observations
6.2 S-type Near-Earth Asteroids
6.2.1 (1917) Cuyo
6.2.2 (8567) 1996 HW1
6.2.3 (16960) 1998 QS52
6.2.4 (188452) 2004 HE62
6.2.5 2010 TD54
6.2.6 (164400) 2005 GN59
6.3 Spectral properties of two primitive NEAs
6.3.1 (5620) Jasonwheeler
6.3.2 2001 SG286
6.4 Discussion
7 Spectral properties of Main Belt Asteroids 
7.1 Log of observations
7.2 (9147) Kourakuen – a V-type asteroid outside Vesta family
7.3 A binary asteroid: (854) Frostia
7.4 1333 and 3623 – two asteroids with large amplitude lightcurves
7.4.1 (1333) Cevenola
7.4.2 (3623) Chaplin
7.5 Asteroid pairs: (10484) Hecht, (31569)1999 FL18
IV CONCLUSIONS AND PERSPECTIVES 
8 Conclusions and perspectives 
A The GuideDog and the BigDog interfaces 
B List of publications 
B.1 First Author
B.2 Co-Author
B.3 Conferences and Workshops

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