search, characterization, and properties of brown dwarfs
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2009
Search, Characterization, And Properties Of Brown Dwarfs Search, Characterization, And Properties Of Brown Dwarfs
Ramarao Tata University of Central Florida
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Search, Characterization, and Properties of Brown Dwarfs
by
Ramarao TataM.S. University of Central Florida, 2005
B.Tech. Sri Venkateshwara University, 2002
A dissertation submitted in partial fulfillment of the requirementsfor the degree of Doctor of Philosophy
in the Department of Physicsin the College of Sciences
at the University of Central FloridaOrlando, Florida
Fall Term2009
Major Professor:Eduardo L. Martın
ABSTRACT
Brown dwarfs (BD) were mere theoretical astrophysical objects for more than three
decades (Kumar (1962)) till their first observational detection in 1995 (Rebolo et al. (1995),
Nakajima et al. (1995)). These objects are intermediate in mass between stars and planets.
Since their observational discovery these objects have been studied thoroughly and holisti-
cally. Various methods for searching and characterizing these objects in different regions of
the sky have been put forward and tested with great success. Theoretical models describing
their physical, atmospheric and chemical processes and properties have been proposed and
have been validated with a large number of observational results.
The work presented in this dissertation is a compilation of synoptic studies of ultracool
dwarfs(UDs)1.
• A search for wide binaries around solar type stars in upper scorpio OB association
(Upper Sco) do indicate (the survey is not yet complete) a deficit of BD binaries at
these large separations (< 5AU).
• Twenty six new UDs were discovered at low galactic latitudes in our survey from
archival data and a novel technique using reduced proper motion.
1 We use the term “ultracool dwarfs” as the mass of most of the objects mentioned is unknown, which isrequired to classify an object as a brown dwarf. We define objects later than M7 as ultra cool dwarfs
iii
• Six field UDs were discovered by spectroscopic follow-up of the candidates selected
from a deep survey.
• Optical interferometry was used to independently determine the orbit of the compan-
ion of HD33636 which was initially determined using Hubble Space Telescope(HST)
astrometry and radial velocity found. Some inconsistency in the HST determined orbit
and mass.
• Optical linear polarization in UDs was used to investigate the dust properties in their
atmospheres. A trend in polarization as predicted by theoretical models was validated,
and atmospheric dust grain sizes and projected rotational velocities for these objects
were estimated
Comprehensive studies of UDs are proving to be crucial not only in our understanding of
UDs but also for star and planet formation as brown dwarfs represent their lower and upper
mass boundaries, respectively.
iv
To my beloved wife, Nicole Tata, for her support and motivation.
v
Acknowledgments
I would like to acknowledge the support of all my collaborators in each of the projects.
1. Chapter2: Jerome Guilet and Rohit Deshpande have contributed in this project with
data reduction and analysis.
2. Chapter3: CTIO facility was utilized for some of the observations. NOAO travel
funding is also appreciated. This project was a collaborative effort led by Phan-Bao
Ngoc.
3. Chapter4: I thank the VLTI staff for granting us the opportunity to test our new idea.
Carla Gil and Eder Martioli’s collaboration is acknowledged.
4. Chapter5: I thank the WHT staff for making these observations possible. This project
is a collaboration between, Eduardo Martin, Sujan Sengupta, Maria Rosa Zapatero,
Phan-Bao Ngoc, and Herve Buoy and was led by me.
I also acknowledge the extensive use of publicly available astronomy resources such as
NOAO-IRAF, Simbad, Vizier, Aladin and STARALT.
vi
TABLE OF CONTENTS
LIST OF FIGURES . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . viii
LIST OF TABLES . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . ix
CHAPTER 1 INTRODUCTION . . . . . . . . . . . . . . . . . . . . . . . . . . 1
CHAPTER 2 A SEARCH FOR BROWN DWARF COMPANIONS OF SOLAR-
TYPE STARS IN UPPER SCO . . . . . . . . . . . . . . . . . . . . . . . . . . . 6
2.1 BACKGROUND . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 6
2.2 METHOD . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 7
2.3 RESULTS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 19
CHAPTER 3 DISCOVERY OF NEW NEARBY L AND LATE-M DWARFS
AT LOW GALACTIC LATITUDE FROM THE DENIS DATABASE . . . 27
3.1 BACKGROUND . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 27
3.2 METHOD . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 28
3.3 RESULTS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 39
vii
CHAPTER 4 USING VLTI OPTICAL INTERFEROMETRY TO VERIFY
HST-ASTROMETRY DERIVED MASS AND ORBITAL PARAMETERS
OF THE COMPANION OF HD33636 . . . . . . . . . . . . . . . . . . . . . . . 47
4.1 BACKGROUND . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 47
4.2 METHOD . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 49
4.3 RESULTS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 55
CHAPTER 5 OPTICAL LINEAR POLARIZATION IN ULTRA COOL DWARFS:
A TOOL TO PROBE DUST IN ULTRA COOL DWARF ATMOSPHERES 64
5.1 BACKGROUND . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 64
5.2 METHOD . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 65
5.3 RESULTS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 75
CHAPTER 6 DISCUSSION . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 82
LIST OF REFERENCES 87
viii
LIST OF FIGURES
2.1 Reduced and stacked I band image of one of the sources V1121Oph . . . . . 9
2.2 Reduced and stacked H band image of one of the sources V1121Oph . . . . . 10
2.3 DENIS catalog magnitude vs I band instrumental magnitude for sources in
the field of view of one of the targets . . . . . . . . . . . . . . . . . . . . . . 11
2.4 2MASS catalog magnitude vs H band instrumental magnitude for sources in
the field of view of one of the targets . . . . . . . . . . . . . . . . . . . . . . 12
2.5 I band magnitude distribution of all the sources in the area of interest (1.5
arcmin box) around all of the targets . . . . . . . . . . . . . . . . . . . . . . 13
2.6 H band magnitude distribution of all the sources in the area of interest (1.5
arcmin box) around all of the targets . . . . . . . . . . . . . . . . . . . . . . 14
2.7 Separation distribution of all the sources in the area of interest (1.5 arcmin
box) around all of the targets . . . . . . . . . . . . . . . . . . . . . . . . . . 15
2.8 Initial selection Color-Magnitude Diagram . . . . . . . . . . . . . . . . . . . 16
2.9 Optical spectra of Upper Sco BD candidates . . . . . . . . . . . . . . . . . . 21
2.10 Color-Magnitude Diagram of our candidates and current status of the project 23
ix
2.11 Color-Magnitude diagrams showing all the candidates . . . . . . . . . . . . . 25
2.12 Color-Magnitude diagrams showing all the candidates . . . . . . . . . . . . . 26
3.1 (MJ , I - J) color-magnitude diagram for late-M and L dwarfs with known
trigonometric parallaxes . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 29
3.2 (I - J, J - K) color-color diagram for our 29 late-M and L dwarf candidates . 31
3.3 I-band reduced proper motions versus I - J . . . . . . . . . . . . . . . . . . . 38
3.4 Spectra of the 26 M8-L5.5 ultracool dwarfs are plotted by increasing spectral
type . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 41
3.5 Number of ultracool dwarfs per 2.5 pc spectrophotometric distance bin over
4,800 square degrees . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 45
4.1 A schematic representation of a 2-element interferometer. . . . . . . . . . . 51
4.2 AMBER data reduction work flow . . . . . . . . . . . . . . . . . . . . . . . 53
4.3 One of the averaged frame of HD33636 . . . . . . . . . . . . . . . . . . . . . 54
4.4 Squared visibility vs wavelength plot . . . . . . . . . . . . . . . . . . . . . . 58
4.5 Flux vs wavelength plot . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 59
4.6 Closure phase vs wavelength plot . . . . . . . . . . . . . . . . . . . . . . . . 60
4.7 Model fit for the visibility measurements . . . . . . . . . . . . . . . . . . . . 61
4.8 Apparent orbit of HD33636 and its companion. . . . . . . . . . . . . . . . . 62
x
4.9 Angular separation vs Orbital inclination . . . . . . . . . . . . . . . . . . . . 63
5.1 Reduced data frame . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 67
5.2 Best model fit of the observed data . . . . . . . . . . . . . . . . . . . . . . . 77
5.3 Spectral energy distribution of 2MASS J17312974+2721233 . . . . . . . . . . 78
xi
LIST OF TABLES
2.1 Ultracool dwarf candidates . . . . . . . . . . . . . . . . . . . . . . . . . . . 17
2.1 Ultracool dwarf candidates . . . . . . . . . . . . . . . . . . . . . . . . . . . 18
2.1 Ultracool dwarf candidates . . . . . . . . . . . . . . . . . . . . . . . . . . . 19
3.1 DENIS ultracool dwarf candidates . . . . . . . . . . . . . . . . . . . . . . . 33
3.1 DENIS ultracool dwarf candidates . . . . . . . . . . . . . . . . . . . . . . . 34
3.2 DENIS ultracool dwarf candidates measured proper motions . . . . . . . . . 35
3.2 DENIS ultracool dwarf candidates measured proper motions . . . . . . . . . 36
3.3 26 nearby L, and late-M dwarfs . . . . . . . . . . . . . . . . . . . . . . . . . 42
3.3 26 nearby L, and late-M dwarfs . . . . . . . . . . . . . . . . . . . . . . . . . 44
5.1 Target list . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 74
5.2 Polarization measurements . . . . . . . . . . . . . . . . . . . . . . . . . . . . 76
5.3 Model fit . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 79
5.4 IRAC photometry of 2MASS J17312974+2721233 . . . . . . . . . . . . . . . 79
xii
CHAPTER 1
INTRODUCTION
Brown dwarfs are sub-stellar objects with masses less than what is required to maintain
hydrogen nuclear fusion in their cores. There is still some debate on how to differentiate
Brown dwarfs from massive planets.The most commonly used mass range for brown dwarfs
is between 13MJ − 78MJ , which represents the deuterium burning limit and the hydrogen
burning limit, respectively.
Brown dwarfs’ mere existence from theory (Kumar (1962)) to their discovery took more
than three decades. In the last fifteen years a lot of progress has been made in our under-
standing of these objects. Ultra cool dwarfs (later than spectral type M7) search criteria
could vary on parameters such as age, distance, metallicity etc., but most of the common
methods involve multi-color imaging surveys. Once a UD candidate is identified based on
multi-color imaging (color-color and/or color-magnitude diagrams), its temperature or mass
needs to be estimated. It turns out that in most cases the temperature estimation is rela-
tively easy to make based on the spectral types. Two new spectral types, the L type and the
T type, were proposed (Martın et al. (1999) & Kirkpatrick et al. (1999)) after the discovery
of the first brown dwarfs. The L type is defined in the red optical region by strengthening
metal hydride bands (FeH, MgH, CrH, CaH), prominent alkali lines (K I - 7655/7699A, Rb
1
I - 7800/7948A, Cs I - 8521/8943A), and weakening metal-oxide bands (VO, TiO). The T
dwarfs are characterized by the presence of methane bands, weakening metal hydrides, and
broad alkali metal absorption features. Based on the statistical properties of brown dwarfs
such as mass function, clustering properties, accretion rates, and binary statistics, one can
argue that brown dwarf formation mechanisms are not much different from low-mass stars.
Although not exclusive, various mechanisms have been proposed for the formation of brown
dwarfs like (Whitworth et al. (2006)) :
• turbulent fragmentation of molecular clouds, forming very low-mass prestellar cores by
shock compression;
• collapse and fragmentation of more massive pre-stellar cores;
• disk fragmentation;
• premature ejection of protostellar embryos from their cores;
• photo-erosion of pre-exisiting cores overrun by HII regions;
The role of these mechanisms probably depends on the environment where the brown
dwarfs form, but their contribution can be judged by their ability to produce observational
parameters like initial mass function (IMF)1 , distribution and kinematics, binary statistics,
and disk retention.1The initial mass function (IMF) is an empirical function that describes the mass distribution (number
of objects per unit of mass) of a population of stars in terms of their theoretical initial mass (the mass theywere formed with).
2
After ground breaking studies by Salpeter (1955), Miller and Scalo (1979), and Scalo
(1986) a lot of effort has been put in studying the stellar IMF in different regions to inves-
tigate the universality of the IMF. Young open clusters and star-forming regions were the
most important targets as they have coeval populations of stars at a known distance and
environment. Some of these studies of open clusters (Kroupa and Boily (2002) , Chabrier
(2003))suggest no clear variation in the IMF and the IMF is consistent with the field IMF.
However, some other studies (Luhman et al. (2003). Luhman (2004) find evidence of a clear
variation in the IMF ( for example Taurus IMF peaks at 0/8M⊙ and IC 348 peaks at 0.1-0.2
M⊙)
Multiplicity of brown dwarfs fall under two main categories: brown dwarf secondaries
around a solar-type primary and around another brown dwarf primary. Brown dwarfs around
solar-type stars are quite rare. At close separations (≤ 5AU), the frequency of companions
with masses in the range 0.01 to 0.1M⊙ is ∼ 0.5% (Marcy and Butler (2000)). This deficit
in brown dwarfs at small separations to solar-type stars relative to frequency of planetary
companions and hydrogen burning stars is commonly known as brown dwarf desert. Various
mechanisms are suggested for the existence of this desert, such as differential fragmentation
in the proto-stellar disk or bias in the inward migration. At larger separations (≥ 100AU)
the binary frequency is < 2% (McCarthy and Zuckerman (2004) & Gizis et al. (2001)). BD-
BD binaries are relatively more numerous. For separations ∼2 AU, BD-BD binaries have an
observed multiplicity of ∼ 10 − 20% (Bouy et al. (2003); Burgasser et al. (2003); Gizis et al.
3
(2003)). Some authors speculate that this multiplicity could be as high as 50% (Pinfield
et al. (2003)).
The atmospheres of ultra cool dwarfs are rich in molecular gas species and metal ele-
ments. These are responsible for the bulk properties such as effective temperature (Teff),
metallicity and the surface gravity. The effective temperature of L dwarfs (1400K . Teff .
2300K) covers a range of temperature where gas-liquid and gas-solid phase transitions for a
number of refractory species occur. These species include iron, silicates, titanates and metal
oxides (Lodders (2002)). The existence of these molecular species has been confirmed in the
spectrum of L dwarfs (depleted TiO and VO which are gaseous precursors to condensates
Burrows and Sharp (1999); Allard et al. (2001)) and their red near infrared colors (thermal
emission from hot dust). Cushing et al. (2006) has detected silicate grain absorption in the
mid infrared. T dwarfs with 600K . Teff . 1400K have relatively condensate free photo-
spheres. This can be inferred observationally from their blue near-infrared colors and strong
molecular gas bands and also by pressure-broadened strong neutral alkali metal absorption
features in the spectrum.
A number of models have been proposed to explain the atmospheric physics of these
objects (Burrows and Sharp (1999); Burrows et al. (2001); Ackerman and Marley (2001a);
Allard et al. (2001); Tsuji et al. (2004); Cooper et al. (2003) ; Helling et al. (2008) ). The
current state of our understanding and observational evidence can be summarized in the
following paragraphs (more details can be found in the review article Burgasser (2009)).
4
One dimensional condensate models have been proposed by Ackerman and Marley (2001a)
to explain BD atmospheres. These models suggest that the condensate species are contrained
between a phase transition layer and a cloud top by a balance between gravitational set-
tling and convective mixing. The cloud top is characterized by a “settling efficiency” or
“cloud top temperature”. In L dwarf atmospheres the cloud layers reside at the photosphere
which makes the effect of condensate species prominent on the spectral energy distribution
of these dwarfs. In T dwarfs the cloud layer sinks below the photosphere which makes the
spectral energy distribution condensate-free. Even though the condensate species in these
dwarfs are constrained between the phase transition layer and the cloud top, the different
species are well mixed by convection and turbulent gas motion. This raises the possibility
of fairly complex grain chemistry including nonequilibrium hybrid grain formation (Helling
et al. (2008)).
Observational evidence supporting the presence of dust clouds comes in a number of
ways. Large near infrared color variations at wavelengths where cloud opacity is prominent
was found in Knapp et al. (2004). Burgasser et al. (2008) shows evidence for a correlation
of the near infrared color of these objects with the 9 µm silicate grain feature. Several ultra
cool dwarfs with atmospheres too neutral to couple with magnetic fields to form spots show
both spectroscopic and photometric temporal variability. This variability probably arises
from rotational modulation of surface asymmetries in global cloud coverage since light curve
periods are generally consistent with rotational line broadening (Reiners and Basri (2008)).
5
The condensate species grain size is also an important parameter which could effect how well
they mix and flow. Ackerman and Marley (2001b) has suggested sub micron size particles.
6
CHAPTER 2
A SEARCH FOR BROWN DWARF COMPANIONS OF
SOLAR-TYPE STARS IN UPPER SCO
2.1 BACKGROUND
The IMF for M ≥ 0.2M⊙ , i.e. very low mass stars and BDs, has been studied in the
solar neighborhood, a few young clusters, and star forming regions. It has been suggested
that loose star forming regions, such as Taurus-Auriga, could have different low-mass IMF
than crowded clusters like the Trapezium (Guieu et al. (2005)). With most of the focus
on open clusters and star-forming regions to derive the low-mass IMF, OB associations
have not been deeply studied for the low-mass IMF because of their large area. Upper Sco
is the youngest (5Myr Preibisch et al. (2001)) of the three subgroups of the nearest OB
association, Scorpius-Centaurus (145 pc de Zeeuw et al. (1999)). Because of its youth and
relative proximity and low reddening (AV ≤ 2), Upper Sco provides ideal search place for
BDs and pre-main sequence objects.
Lately Upper Sco has been studied, and more than 50 objects have estimated masses
that range from 0.2 to 0.02M⊙ in about a 75 sq degree area(Ardila et al. (2000); Martın
7
et al. (2004)). Also follow-up studies have shown evidence of mass accretion and mid-infrared
emission due to dusty disks.
We carried out a deep survey to identify BD companions orbiting at large distances (wide
binaries) from solar-type stars in the Upper Sco OB association. Our search is sensitive to
companions as faint as I=20 and 2 arcsec , which according to current evolutionary models
for brown dwarfs, correspond to masses down to 10MJ and 300 AU from the primary stars.
The frequency of such wide companions to solar-type stars is still uncertain, and the most
optimistic estimate is for 15% (Gizis et al. (2001)). Recently it has been proposed that this
frequency could be no larger than 3% (McCarthy and Zuckerman (2004)). Our study aims
at shedding more light on the apparent controversy.
2.2 METHOD
We have observed 65 G and K type stars in the Upper Sco OB association. Our objects were
selected from Preibisch et al. (2001) supplemented by H.Bouy proper motion study (private
communication) made available to us. We have obtained H band imaging with the 1.5m
Carlos Sanchez Telescope and I band Imaging with the IAC-80 Telescope for these objects
in early June 2006. Both of these telescopes are located in the Teide Observatory Complex
on the island of Tenerife. The individual images (dither pattern for each source) were bias
subtracted and flat fielded before stacking. The effective exposure time for the I band and
H band images were 1200 seconds and 800 seconds, respectively. Aperture photometry was
8
performed using the APPHOT package in IRAF. The calibration was done using stars in the
field of view of each image and comparison with DENIS and 2MASS photometry. Figure 2.1
and 2.2 show the final reduced images of one of the target star in I and H band respectively.
The color magnitude criteria used to select the candidates (shown in circles) are discussed
later in this section. Figure 2.3 and 2.4 show the photometric calibration for sources around
one of the target stars. The linear fit parameters were used to calibrate all the sources
around that target i.e. every a different calibration plot and fit were obtained to calibrate
sources around each of the target star. Figure 2.5, 2.6 and 2.7 show the distribution of the
I band magnitude, the H band magnitude and the projected separation from the target star
respectively for all of selected sources in the area of interest. This shows that our study is
sensitive for companions for up to 18th magnitude in H band and 20th magnitude in the I
band and companions at a minimum angular separation of 5 arcseconds.
Once the photometry was done for the I and the H band for the sources within a 1.5
arcmin box around the primary star, we plotted the I vs I-H plot for all the sources shown in
figure 2.8. Theoretical model isochrones were over plotted for a range of ages and distances.
BD candidates were selected based on their position in the plot, within the region surrounded
by 1Myr - 10 Myr isochrones in 2σ error bars of the photometric measurements. A list of
41/60 of the selected candidates is given in Table 2.1. The other 19 objects were selected
based on the 2MASS colors. These 19 candidates were detected only in our H band images
but had 2MASS J and K detections. In Table 2.1, column 1 gives the name of the target
9
Figure 2.1: Reduced and stacked I band image of one of the sources V1121Oph. The effective
exposure time was ∼1200s. The field of view shown here ∼ 100x100 arcsec. The
photometric UD candidates are shown in the circles.
10
Figure 2.2: Reduced and stacked H band image of one of the sources. The effective exposure time
was ∼800s. The field of view shown here ∼ 100x100 arcsec. The photometric UD
candidates are shown in the circles.
11
Figure 2.3: DENIS catalog magnitude vs I band instrumental magnitude for sources in the field of
view of one of the targets. The linear fit was then used to calibrate the instrumental
magnitudes of all the sources and estimate the calibration error bars.
12
Figure 2.4: 2MASS catalog magnitude vs H band instrumental magnitude for sources in the field
of view of one of the targets. The linear fit was then used to calibrate the instrumental
magnitudes of all the sources and estimate the calibration error bars.
13
Figure 2.5: I band magnitude distribution of all the sources in the area of interest (1.5 arcmin
box) around all of the targets.
star and column 2 gives the projected separation (assuming all the target stars are at 145
pc).
14
Figure 2.6: H band magnitude distribution of all the sources in the area of interest (1.5 arcmin
box) around all of the targets.
15
Figure 2.7: The black curve shows the separation distribution of all the sources in the area of
interest (1.5 arcmin box) around all of the targets. The blue curve shows the same
distribution divided by the separation. The projected separation was calculated as-
suming the distance of all the targets at 145 pc.
16
Figure 2.8: Initial selection Color-Magnitude Diagram of our candidates. The red lines are iso-
chrones for different ages using dusty brown dwarf evolutionary models for 1, 5, and
10Myr from right to left, respectively. The blue isochrone is the CONDENSATE
model, and green is the NEXGEN model for 5Myr.
17
Table 2.1. Ultracool dwarf candidates
Target name Separation (AU) I meas I err H meas Herr I-H
(1) (2) (3) (4) (5) (6) (7)
RX1625.2-2455ab 3356.43 17.34 0.02 14.33 0.04 3.01
RX1625.2-2455ab 4722.62 19.29 0.07 16.14 0.07 3.15
RX1625.2-2455ab 5046.56 20.61 0.21 16.97 0.11 3.64
RX1625.2-2455ab 5099.43 19.93 0.11 16.74 0.1 3.19
v1121oph 6381.04 18.47 0.05 14.93 0.05 3.54
v1121oph 5332.3 17.48 0.04 14.76 0.05 2.72
v1121oph 7141.76 18.43 0.05 15.19 0.06 3.24
v1121oph 6329.28 20.05 0.13 15.97 0.1 4.08
pds145 1158.32 16.09 0.02 13.88 0.04 2.2
pds145 8190.5 14.51 0.01 12.42 0.03 2.09
woph5 4617.3 17.67 0.03 14.85 0.04 2.83
woph5 3833.62 17.25 0.03 13.95 0.03 3.3
woph5 873.57 16.7 0.03 13.97 0.06 2.73
woph5 5961.86 19.05 0.06 16.02 0.1 3.04
gsc997 4483.35 15.32 0.03 13.15 0.03 2.17
gsc997 4337.93 16.71 0.04 13.7 0.03 3.01
gsc1406 1280.45 17.17 0.07 14.85 0.04 2.32
gsc1406 5089.11 18.7 0.12 15.84 0.06 2.86
18
Table 2.1 (cont’d)
Target name Separation (AU) I meas I err H meas Herr I-H
(1) (2) (3) (4) (5) (6) (7)
rox20-1 4095.03 19.08 0.05 15.91 0.07 3.17
rox20-1 3158.27 20.17 0.15 16.56 0.11 3.61
haro8jan 4988.57 18.79 0.04 15.51 0.05 3.28
haro8jan 1472.52 18.64 0.06 15.17 0.05 3.47
haro8jan 1959.4 18.18 0.03 15.5 0.06 2.68
haro8jan 5837.1 17.15 0.02 13.98 0.03 3.17
1rs48 3701.21 19.66 0.1 15.1 0.04 4.56
6228-13 5229.53 14.42 0.03 12.27 0.03 2.15
v866-sco 3287.73 18.06 0.03 15.26 0.05 2.8
B62-HA4 5867.5 15.4 0.01 12.69 0.03 2.71
B62-HA4 5040.51 18.27 0.09 15.27 0.05 4.49
B62-HA4 2535.94 18 0.03 15.29 0.05 2.7
B62-HA4 4917.42 15.6 0.01 13.56 0.04 2.05
B62-HA4 2837.57 15.9 0.01 13.42 0.03 2.48
HARO4JAN 6118.65 17.05 0.02 14.66 0.04 2.4
HARO4JAN 4379.82 16.41 0.01 14.09 0.04 2.31
V1725 4303.73 19.46 0.07 15.2 0.04 4.26
V1725 2569.38 18.83 0.04 14.74 0.04 4.09
19
Table 2.1 (cont’d)
Target name Separation (AU) I meas I err H meas Herr I-H
(1) (2) (3) (4) (5) (6) (7)
gsc6794-537 5884.24 19.06 0.06 16 0.09 3.06
gsc6794-537 5155.27 16.3 0.04 13.8 0.03 2.5
gsc6794-537 5082.41 19.15 0.07 15.47 0.06 3.68
WOph1 3554.58 19.16 0.15 16.39 0.1 2.77
pds91 1311.13 18.57 0.05 15.83 0.09 2.74
The separation was calculated assuming the distance of Upper Sco as 150pc
Follow-up spectroscopic observations were performed for a small fraction (< 15%) of the
candidates in the optical and/or near infrared wavelengths. These observations were made
using the DOLORES/3.6m Galileo telescope, the ISIS/4m William Herschel Telescope, and
the RCSpec/4m Blanco CTIO telescope in May 2007, June 2007, and July 2007, respectively.
The data was reduced from all the instruments was reduced using the APALL task in the
IRAF software.
2.3 RESULTS
We were able to get spectra for 8 out of 60 BD candidates. A quick analysis ruled out 7
of these candidates as background objects. One of the object had an M-dwarf-like spectra,
20
but it was also ruled out as it did not have any H-α emission which is typical for substellar
objects in Upper Sco. The normalized spectra are shown in figure 2.9
21
Fig
ure
2.9:
Opt
ical
spec
tra
ofU
pper
Sco
BD
cand
idat
es.
Non
eof
thes
esp
ectr
asu
gges
ta
youn
gsu
bste
llar
obje
ct.
22
It will be very premature to speculate about wide binary BDs in Upper Sco with the
current status of our study. We can still place an upper limit on the frequency of Sun-
like stars with with BD companions in our sample. Out of 65 stars that were observed,
even though most of them had sources close enough to be in our search box, we only found
candidates around 16 of these stars. 7 stars out of the 16 stars can be ruled out as not having
any BD companions based on follow-up spectroscopy. Hence, we place an upper limit of 14%
for Sun-like stars having BD companions. Figure 2.10 shows that the candidates with follow-
up spectra fall uniformly around the isochrone, and any contaminents are probably well
intertwined with the BDs in this plot. To differentiate contaminants from actual dwarfs, we
also plotted the 2MASS color-magnitude diagrams for our candidates with 2MASS detections
(see Figures 2.11 and 2.12). We do not see any trend that we can use to filter out the
contaminants.
23
I vs I-H14
15
16
17
18
19
20
21
1.5 2.5 3.5 4.5
I-H
I
Candidates
Observed
5Myr DustyModel
No 2MASSdetection
Red Sourcesnon-candidates
M4
6''
8''
9''
9''10''
14''
Figure 2.10: Color-Magnitude Diagram of our candidates and current status of the project
24
The current estimation of the upper limit keeps us optimistic on finding some BDs among
our candidates. Based on what we find in future, we will be able to answer the questions:
whether star formation efficiency is mass dependent or not and whether BD formation by
fragmentation of a circumstellar disc is a major contributing mechanism. If we indeed see
a BD desert at wide separations, it becomes quite challenging to explain BD formation
within the context of star formation. McCarthy and Zuckerman (2004) had even suggested
that perhaps star formation efficiency contains a mass dependent term like [1 − e−M/Mcrit ],
where Mcrit ∼ 0.1 − 0.2M⊙. Burgasser et al. (2005), Bate et al. (2002), Bate et al. (2003)
have proposed a mechanism where most of the BDs are ejected from stellar systems before
accreting enough mass to become stars. However, this mechanism can not explain the BD
desert at small separations, as they would be more tightly bound.
Burgasser et al. (2005) has also proposed the possibility of some BD-BD binaries may
survive the ejection after forming in the disc of a Sun-like star due to higher binding energy
of the BD-BD system with the sun. This makes the four red sources (Figure 2.3 sources to
the right of the isochrones) not selected as candidates in the current study viable candidates
for follow-up spectroscopy.
25
14
15
16
17
18
19
20
2 2.5 3 3.5 4 4.5 5 5.5 6I-K
I
Candidates
Dusty Model
Observed
Figure 2.11: Color-Magnitude diagrams showing all the candidates (with 2MASS detections -
open circles) and candidates observed for optical spectra (filled circles). A Dusty
model isochrone for 5Myr is overplotted (dashed line).The J, K photometry is derived
from 2MASS.
26
12
13
14
15
16
17
18
19
0 1 2 3 4 5 6 7 8J-K
J
Candidates
Dusty Model
Observed
Figure 2.12: Color-Magnitude diagrams showing all the candidates (with 2MASS detections -
open circles) and candidates observed for optical spectra (filled circles). A Dusty
model isochrone for 5Myr is overplotted (dashed line).The J, K photometry is derived
from 2MASS.
27
CHAPTER 3
DISCOVERY OF NEW NEARBY L AND LATE-M DWARFS
AT LOW GALACTIC LATITUDE FROM THE DENIS
DATABASE
3.1 BACKGROUND
Nearby stars are the brightest representatives of their class and therefore provide observa-
tional benchmarks for stellar physics. This is particularly true for intrinsically faint objects,
such as stars at the bottom of the main sequence and brown dwarfs. In the last decade,
many nearby ultracool dwarfs have been discovered by using the DENIS (Epchtein et al.
(1997)), 2MASS (Skrutskie (1997)), SDSS (York et al. (2000)), UKIDSS (Lawrence et al.
(2007)) surveys (Delfosse et al. (1999); Martın et al. (1999);Phan-Bao et al. (2001);Phan-
Bao et al. (2003); Reyle et al. (2002); Kirkpatrick et al. (1999); Reid and Cruz (2002); Cruz
et al. (2003); Burgasser et al. (2004); Deacon et al. (2005); Lodieu et al. (2005) Lodieu et al.
(2007); Knapp et al. (2004); Chiu et al. (2006)) or the proper motion measurement (Lepine
et al. (2002)). However, most surveys for nearby UDs have tended to avoid the crowded
regions of the galactic plane. The first systematic search for UDs (|b| < 10◦) was carried out
28
by Reid (2003), where he discovered one M8 and reidentified one L1.5 previously discovered
by Salim et al. (2003). Hambaryan et al. (2004) reported an M9.0 dwarf 1RXS J115928.5-
524717 (b = 9.3◦) showing strong X-ray flaring emission. Recently, Folkes et al. (2007)
discovered an L-T transition object close in the galactic plane 2MASS J11263991-5003550
(L9, b = 10.6◦). Four other UDs at low galactic latitudes (|b| < 15◦) were also discovered:
one M8 (Cruz et al. (2003)); one L2 (Scholz and Meusinger (2002)); one L0 and one L4.5
(Kendall et al. (2007)), the latter L4.5 dwarf is actually an L1+L3.5 (Kendall et al. (2007))
or L1.5+L4.5 (Burgasser et al. (2007)) binary. These discoveries of UDs in the galactic plane
provide additional targets with which basic physical properties of stars at the bottom of the
main sequence and brown dwarfs can be studied; however, many remain to be discovered.
The results of this study are already published in Phan-Bao et al. (2008).
3.2 METHOD
We identified 47 very cool and nearby (d < 32 pc) dwarfs in the literature with trigonometric
parallaxes and good photometry (e.g. errors smaller than 0.2 mag) . For our limited present
purpose the DENIS system is close enough to both the Cousins-CIT (≤ 0.05 mag, Delfosse
et al. (1997)) and 2MASS systems (≤ 0.02 mag, Carpenter (2001)). We therefore ignored the
small corrections needed to transform between these photometric systems. We did, on the
other hand, transform the I814 magnitudes of (Reid et al. (2001); Burgasser et al. (2003))
to IC , using the linear correction of Reid et al. (2001). Figure 3.1 shows the resulting (I - J,
29
MJ) plot, and the corresponding 4th order polynomial fit:
MJ = a0 + a1(I − J) + a2(I − J)2 + a3(I − J)3 + a4(I − J)4 (3.1)
where a0 = 130.164, a1 = −124.899, a2 = 46.948, a3 = −7.4842, a4 = 0.4341, valid for 3.0
≤ I − J ≤ 6.0. The rms dispersion around that fit is 0.3mag, corresponding to a 14% error
on distances.
We systematically searched the DENIS database (available at the Paris Data Analysis
Center, PDAC) for potential members of the solar neighbourhood, with simple and well
defined criteria. We first selected all DENIS objects which matched |b| < 15◦ (low galactic
latitude) and I − J > 3.0 (spectral type later than nominally M8.0, Leggett (1992)). We
then required that the position of the candidates in the (I - J, J - K) color-color diagram be
within J − K = ±0.5 mag of our linear fit to the locus of ultracool dwarfs (Figure 3.2). We
then computed photometric distances using the (MJ , I - J) relation established in equation
3.1 and retained the candidates with dphot ≤30 pc.
Before setting out to measure the labour-intensive proper motions needed to use reduced
proper motions as a discriminant between nearby dwarfs and distant giants, we eliminated
the bright candidates (I < 13.0, dphot ø 30pc) within the boundaries of known low galactic
latitude molecular clouds in SIMBAD. This left 29 ultracool dwarf candidates (Table 3.1)
for which we needed to measure proper motions. Their color inferred spectral types range
from M8.0 to ∼ L8.0 (Figure 3.2) .We also searched for T dwarf candidates by using the (I
- J, J - K) color-color diagram along with a linear fit to the colors of T dwarfs, however no
30
Figure 3.1: (MJ , I - J) color-magnitude diagram for late-M (> M8.0) and L dwarfs with known
trigonometric parallaxes. 2MASS 0559-1404 (T5.0, Burgasser et al. 2002) is overlu-
minous, by more than the 0.75 magnitude offset expected for an equal mass binary
system.
31
Figure 3.2: (I - J, J - K) color-color diagram for our 29 late-M and L dwarf candidates (Table
3.1), plotted as solid circles, as well as known late-M and L dwarfs; plotted as open
triangles). Representative error bars for one object are plotted.The line represents
our linear fit to the colors of the literature dwarfs:J-K=-0.462+0.532(I-J), σ = 0.15
32
T dwarf candidates were found. This reflects the fact that the DENIS survey with limiting
magnitudes of I = 18.5, J = 16.5, K = 14.0 is not very sensitive for T dwarf detections. For
example, a T0.5 dwarf (MI ∼ 19.3, MJ ∼ 14.8 and MJ ∼ 13.2, Vrba et al. (2004)) could
be detected by DENIS in all three I, J and K bands were its distance within ∼ 7 pc, e.g.
Epsilon Indi B (Scholz et al. (2003)). A relaxation in the request for detections in all three
DENIS bands, for example requiring detections in the J and K bands but not the I band
may allow us to detect nearby T dwarfs. Moreover, finding T dwarfs was not the primary
objective of our study.
We queried ALADIN and the Digital Sky Survey (DSS) server for publically available
scanned plates containing these candidates. Both ALADIN and the DSS archive provide
access to the plates of the POSS-I (R-band), SERCR (R-band), SERC-I (I-band), POSS-II F
(R-band), and POSS-II N (I-band) surveys. ALADIN additionally contains digitized images
of the ESO-R survey (R-band), which often extend the time baseline enough to significantly
reduce the standard error of our proper motion measurements.We used SExtractor (Bertin
& Arnouts 1996) to extract coordinates of our targets from all available digitized images.
We then determined proper motions through a least-square fit to the positions at the 3 to
5 available epochs with time baselines spanning from 10 to 21 years, including DENIS and
2MASS. The uncertainty on the proper motion measurement is the rms dispersion around
that fit. Table 3.2 present our proper motions and the associated standard errors for 26
ultracool dwarfs.
33
Table 3.1. DENIS ultracool dwarf candidates
Target name RA(J2000) DEC (J2000) epoch I I-J J-K
(1) (2) (3) (4) (5) (6) (7)
J0615493-010041 06 15 49.32 -01 00 41.2 1996.06 17 3.24 1.45
J0630014-184014 06 30 01.43 -18 40 14.7 1999.863 15.88 3.17 1.33
J0644143-284141 06 44 14.34 -28 41 41.9 1996.954 16.96 3.17 1.14
J0649299-154104 06 49 29.98 -15 41 04.0 1998.123 18.36 3.96 1.74
J0652197-253450 06 52 19.73 -25 34 50.5 2001.066 15.95 3.26 1.17
J0716478-063037 07 16 47.88 -06 30 37.5 1997.107 17.46 3.55 1.3
J0719234-173858 07 19 23.44 -17 38 58.8 1996.948 18.1 3.76 2
J0719358-174910 07 19 35.86 -17 49 10.4 1998.978 18.06 3.52 1.87
J0751164-253043 07 51 16.44 -25 30 43.3 1999.189 16.53 3.31 1.15
J0805110-315811 08 05 11.03 -31 58 11.8 2000.06 16.5 3.29 0.84
J0812316-244442 08 12 31.69 -24 44 42.1 1996.301 17.27 3.38 1.39
J0823031-491201 08 23 03.17 -49 12 01.3 1999.973 17.14 3.56 1.72
J0828343-130919 08 28 34.34 -13 09 19.9 1996.109 16.09 3.37 1.4
J1048278-525418 10 48 27.82 -52 54 18.4 2001.359 17.25 3.26 1.45
J1126399-500355 11 26 39.93 -50 03 55.3 1999.263 17.8 3.85 1.21
J1157480-484442 11 57 48.08 -48 44 42.6 1996.246 17.41 3.38 1.41
J1159274-524718 11 59 27.42 -52 47 18.7 1999.389 14.49 3.12 1.08
J1232178-685600 12 32 17.80 -68 56 00.7 1999.167 15.43 3.05 1.08
34
Table 3.1 (cont’d)
Target name RA(J2000) DEC (J2000) epoch I I-J J-K
(1) (2) (3) (4) (5) (6) (7)
J1253108-570924 12 53 10.87 -57 09 24.9 2000.216 16.74 3.29 1.42
J1347590-761005 13 47 59.07 -76 10 05.5 1999.189 17.05 3.24 1.38
J1454078-660447 14 54 07.84 -66 04 47.1 1998.375 16.9 3.7 1.48
J1519016-741613 15 19 01.62 -74 16 13.9 1996.322 16.56 3.07 0.92
J1520022-442242 15 20 02.24 -44 22 42.2 1999.301 16.69 3.45 1.42
J1705474-544151 17 05 47.41 -54 41 51.2 2000.301 16.66 3.2 1.27
J1733423-165449 17 33 42.31 -16 54 49.8 1996.369 16.82 3.21 1.17
J1745346-164053 17 45 34.67 -16 40 53.7 2000.552 17.11 3.37 1.36
J1756296-451822 17 56 29.64 -45 18 22.4 1999.321 15.46 3.14 1.13
J1756561-480509 17 56 56.19 -48 05 09.5 2000.533 16.74 3.41 1.19
J1909081-193748 19 09 08.17 -19 37 48.2 2001.466 17.92 3.58 1.72
35
Table 3.2. DENIS ultracool dwarf candidates measured proper motions
DENIS name Time No µRA errµRA µDEC errµDEC HI errHI HmaxI d errd
(b.l. years) of Obs. 00yr−1 00yr−1 mudec errdmue mag mag mag pc pc
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) (12)
D* J0615493-010041 15.616 4 0.226 0.012 -0.075 0.014 18.9 0.2 7.9 26.8 4.7
D* J0630014-184014 11.792 4 0.35 0.008 -0.503 0.004 19.8 0.1 7.9 18 3.1
D* J0644143-284141 10.912 4 0.329 0.019 -0.105 0.039 19.7 0.3 7.9 29.6 5.3
D* J0652197-253450 21.151 3 -0.233 0.005 0.086 0.002 17.9 0.1 7.9 16 2.8
D* J0716478-063037 18.12 3 -0.016 0.013 0.121 0.005 17.9 0.2 8.1 18.7 3.4
D* J0751164-253043 15.123 3 -0.885 0.003 0.142 0.004 21.3 0.1 7.9 19.1 3.3
D* J0805110-315811 21.024 3 -0.231 0.004 0.103 0.01 18.5 0.2 7.9 19.6 3.4
D* J0812316-244442 18.899 3 0.096 0.001 -0.165 0.007 18.7 0.2 8 23.7 4.4
D* J0823031-491201 20.994 3 -0.137 0.005 0.017 0.001 17.8 0.2 8.1 15.9 2.9
D* J0828343-130919 14.293 4 -0.547 0.018 0.078 0.009 19.8 0.1 8 14 2.4
D* J1048278-525418 21.053 3 -0.179 0.009 0.033 0.017 18.6 0.3 7.9 29.1 5.5
D* J1126399-500355 14.153 3 -1.57 0 0.438 0.011 23.9 0.2 8.4 12.5 2.2
D* J1157480-484442 16.178 3 -0.052 0.016 0.001 0.001 16 0.8 8 25.3 4.5
D* J1159274-524718 19.154 5 -1.067 0.003 -0.119 0.026 19.6 0 7.8 10.2 1.7
D* J1232178-685600 15.89 3 -0.196 0.01 -0.068 0.008 17 0.2 7.8 17.4 2.9
D* J1253108-570924 20.033 4 -1.575 0.001 -0.435 0.016 22.8 0.1 7.9 21.8 3.8
D* J1347590-761005 19.836 4 0.193 0.009 0.049 0.02 18.5 0.3 7.9 27.5 4.9
36
Table 3.2 (cont’d)
DENIS name Time No µRA errµRA µDEC errµDEC HI errHI HmaxI d errd
(b.l. years) of Obs. 00yr−1 00yr−1 mudec errdmue mag mag mag pc pc
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) (12)
D* J1454078-660447 19.054 3 0.525 0.002 -0.376 0.006 21 0.1 8.3 10.9 1.9
D* J1519016-741613 15.781 3 0.317 0.01 0.097 0.009 19.2 0.1 7.8 28.5 5
D* J1520022-442242 20.978 4 -0.623 0.008 -0.39 0.016 21 0.1 8 15.9 2.7
D* J1705474-544151 20.879 4 -0.072 0.008 0.03 0.001 16.1 0.3 7.9 24.5 4.4
D* J1733423-165449 17.905 3 0.081 0.015 -0.048 0.015 16.7 0.6 7.9 26 4.7
D* J1745346-164053 17.914 4 0.116 0.005 -0.111 0.019 18.1 0.3 8 22.4 4.2
D* J1756296-451822 19.916 5 0.064 0.005 -0.183 0.006 16.9 0.1 7.8 15.5 2.6
D* J1756561-480509 13.045 4 0.078 0.004 0.05 0.011 16.6 0.3 8 17.6 3.2
D* J1909081-193748 14.776 5 -0.064 0.021 -0.145 0.034 18.9 0.7 8.1 21.9 4.2
We use the Maximum Reduced Proper Motion (MRPM Phan-Bao et al. (2003)) method
to reject giants from our sample. This method uses the reduced proper motion (H = M + 5
log(Vt/4.74) = m+5+log(µ)) versus color diagram, where nearby ultracool dwarfs and distant
giants segregate very clearly. We use the color-magnitude relation of giants to compute the
maximum reduced proper motion that a red giant can have at a given color setting its
tangential velocity to the galactic escape velocity, for which we adopt a conservative 800 km
s−1 (see Phan-Bao et al. (2003) and references therein) , and then declare any object with a
reduced proper motion above that value for its color to be a dwarf. To use this method here
37
for the cooler M8.0-L8.0 candidates, we had to extend the maximum reduced proper motion
curve of giants towards redder colors (I - J ≥3.0). The following cubic least-square fit (see
Fig. 3.3) is valid for 3.0 ≤ I − J ≤ 4.5, or dwarf spectral types between M8.0 and L8.0.
HmaxI = 6.79 + 1.25(I − J) − 0.630(I − J)2 + 0.108(I − J)3 (3.2)
The rms dispersion around the fit is 0.1 mag. Figure 3.3 shows the position of the resulting
HmaxI vs (I-J) curve relative to the 29 candidates. The curve well classifies 26 candidates as
late-M and L dwarfs. The likely giants are DENIS 0649-1541 and DENIS 0719-1749 that are
in the ultracool dwarf part of the diagram but within the 2 σ limit, and the remaining one
DENIS 0719-1738 that is within 1σ.
We observed the candidates on December 17-18 2006 and on May 20-25 2007 with the
Double Beam Spectrograph on the ANU (Australian National University) 2.3m telescope
at Siding Spring Observatory. Both gratings 158g/mm and 316g/mm were used and they
provided a wavelength coverage of 6100-10200 A at 5 and 2A resolution for gratings 158g/mm
and 316g/mm, respectively. Exposures of 600-1800 seconds were taken, depending on the
target magnitude. The seeing was around 2.0 arcsec . The data were reduced using the
FIGARO5 data reduction system. Smooth spectrum stars were observed at a range of
airmass to remove the telluric lines using the technique of Bessell (1999). The EG131 (Bessell
(1999)) spectrophotometric standard was used for relative flux calibration, and a NeAr arc
provided the wavelength calibration. All spectra were normalized over the 7540-7580 A
interval, the denominator of the PC3 index ( Martın et al. (1999) ) and a region with a
38
Figure 3.3: I-band reduced proper motions versus I - J. The dashed curve represents the maximum
possible reduced proper motion for a giant, HmaxI . Objects above this curve must be
dwarfs. Three red, distant dwarfs are indicated: DENIS 0649-1541, DENIS 0719-1738
and DENIS 0719-1749
39
good flat pseudo-continuum. We also observed VB 8 (M7, for spectral classification, see
Martın et al. (1999)), VB 10 (M8, Martın et al. (1999)), LHS 2065 (M9, Martın et al.
(1999)), 2MASS J22431696-5932206 (L0, Kendall et al. (2007)), 2MASS J01282664-5545343
(L1, Kendall et al. (2007)) and 2MASS J09211410-2104446 (L2, Reid et al. (2006)) with
the same instrument setup and used them as templates. At the resolution of the spectra,
M and L dwarfs are immediately distinguished from M giants by the presence of the NaI
and KI doublets, the presence of FeH bands, the appearance of strong CaH cutting into the
continuum shortward of 7000 A, and by the absence of the CaII triplet (e.g. Bessell 1991).
To estimate the spectral types of L and late-M dwarfs, we used the PC3 index defined
in Martın et al. (1999) and TiO5 defined in Reid et al. (1995). The VOa index, defined
in Kirkpatrick et al. (1999), saturates in the M8-L0.5 spectral type range which covers a
significant number of our ultracool dwarfs, we therefore only used it as a reference when
there was a large difference (e.g. greater than 2 subclasses) in the spectral type estimate
between PC3 and TiO5. We used spectral type versus spectral index relations to estimate
spectral types, the Cruz and Reid (2002) relation for TiO5 and the Martın et al. (1999)
relation for PC3. The uncertainty in the spectral indices was calculated based on the rms
dispersion of the flux around the mean in the wavelength regions that were used to calculate
the spectral indices.
40
3.3 RESULTS
Our low-resolution optical spectra confirmed that 26 were indeed UDs, with spectral types
from M8.0 to L5.5.Two contaminants and one rejected by the maximum reduced proper
motion cut-off were all reddened F-K main sequence stars. Twenty of these 26 UDs are new
nearby UD members, three L dwarfs within 15 pc with one L3.5 at only ∼ 10 pc. Figure 3.4
shows the optical spectra of the 26 dwarfs. Table 3.3 shows the determined spectral types
and other measured parameters for them.
41
Figure 3.4: Spectra of the 26 M8-L5.5 ultracool dwarfs are plotted by increasing spectral type
42
Table 3.3. 26 nearby L, and late-M dwarfs
DENIS name SpT MJ dsp err Vt errVt EW
pc pc km/s km/s HαA
(1) (2) (3) (4) (5) (6) (7) (8)
DENIS 0615-0100 L2.5 12.74 16 3.3 18 5 <7.0
DENIS 0630-1840 M8.5 11.28 19.3 3.9 56 12 <18.0
DENIS 0644-2841 M9.5 11.65 26.8 5.6 44 13 <10.0
DENIS 0652-2534 L0.0 11.83 14.9 3 17 4 <1.0
DENIS 0716-0630 L1.0 12.19 22 4.6 13 3 <4.0
DENIS 0751-2530 L1.5 12.38 14.8 2.9 63 13 <5.0
DENIS 0805-3158 M8.0 11.1 26.4 5.4 32 7 <3.0
DENIS 0812-2444 L1.5 12.38 20.1 4.3 18 5 <5.0
DENIS 0823-4912 L1.5 12.38 17.4 3.6 11 3 <5.0
DENIS 0828-1309 L1.0 12.19 12.7 2.5 33 8 <4.0
DENIS 1048-5254 L1.5 12.38 21 4.6 18 5 <3.0
DENIS 1126-5003 L5.5 13.83 10.6 2.2 82 17 <12.0
DENIS 1157-4844 L1.0 12.19 23.3 4.9 6 3 <2.0
DENIS 1159-5247 M9.0 11.47 9.6 1.9 49 10 15
DENIS 1232-6856 M8.0 11.1 18 3.5 18 4 <6.0
DENIS 1253-5709 L0.5 12.01 19.4 4 150 31 <7.0
DENIS 1347-7610 L0.0 11.83 24.9 5.2 23 7 <10.0
43
In our study only the M9 dwarf DENIS 1159-5247 (or 1RXS J115928.5-524717) shows
strong Hα emission. This M9 dwarf has also shown strong X-ray flaring emission (Fuhrmeis-
ter and Schmitt (2003); Hambaryan et al. (2004)) demonstrating that the dwarf has chromo-
spheric and coronal activity. We measured an upper limit for the remaining dwarfs, which
have weak or no Hα emission or low signal-to- noise spectra. As expected, the dissipation
of chromospheric activity in almost our ultracool dwarfs is due to their predominantly cool,
dense, and highly neutral atmospheres (Meyer & Meyer-Hofmeister (1999); Mohanty et al.
(2002)).
The differential spectrophotometric distance distribution of that sample (Figure 3.5) is
well fitted by a d2 distribution, as expected for a constant density population, out to ∼ 22 pc.
The difference from the initial 30 pc selection cutoff reflects the detection limiting distance of
DENIS for ultracool dwarfs and the different color-magnitude relations used in the selection
and in the final spectrophotometric distance estimate. We conservatively adopted 22 pc
as the completeness limit of our sample, and use the 18 objects within that distance to
determine the local density of ultracool dwarfs. A sample limited by spectrophotometric
distance is effectively a magnitude-limited sample with a spectral type (or color)-dependent
magnitude limit. As such, the dispersion in the spectral type-absolute magnitude relation
has two distinct effects on the statistics of the stellar population (Stobie et al. (1989)). First,
the average luminosity of stars at a given spectral type (or color) is increased; this is the
classical Malmquist bias (Malmquist (1936)). Second, due to the dispersion of the distance
estimator a magnitude-limited sample includes more stars at a given spectral type. Here we
44
Table 3.3 (cont’d)
DENIS name SpT MJ dsp err Vt errVt EW
pc pc km/s km/s HαA
(1) (2) (3) (4) (5) (6) (7) (8)
DENIS 1454-6604 L3.5 13.1 10.5 2.1 32 7 <5.0
DENIS 1519-7416 M9.0 11.47 25.4 5.2 40 10 <2.0
DENIS 1520-4422 L1.0 12.19 16.2 3.2 56 12 <4.0
DENIS 1705-5441 M8.5 11.28 27.2 5.7 10 3 <6.0
DENIS 1733-1654 L1.0 12.19 19.2 4 9 4 <2.0
DENIS 1745-1640 L1.5 12.38 18.7 4.1 14 5 <9.0
DENIS 1756-4518 M9.0 11.47 14.8 3 14 3 <2.0
DENIS 1756-4805 L0.0 11.83 20 4.3 9 3 <13.0
DENIS 1909-1937 L1.0 12.19 26.9 6 20 10 <5.0
dsp - our computed spectrophotometric distance
Vt - tangential velocity computed using dsp
EWHα was measure from the spectra using IRAF SPLOT task
45
Figure 3.5: Number of ultracool dwarfs per 2.5 pc spectrophotometric distance bin over 4,800
square degrees. The errorbars are poissonian 1σ errors and the curve is the expected
d2 distribution, normalized at 16 pc.
46
study the sample of ultracool dwarfs M8-L3.5 over 11.1≤ MJ ≤13.1. The first component
of the Malmquist bias is therefore irrelevant, since firstly we donot look for any significant
luminosity resolution, and secondly the luminosity function is sufficiently flat over theM8-
L3.5 spectral range (Cruz et al. (2007)) that a small shift in the average luminosity will not
measurably affect the resulting density. The second component of the bias, on the other
hand, is significant. For a gaussian dispersion of the color-luminosity relation it can be
computed analytically(Stobie et al. (1989)):
∆Φ
Φ=
1
2σ2(0.6ln10)2 (3.3)
whereΦ is the luminosity function andσ is the intrinsic rms scatter in the spectral type-
luminosity relation.
The mean surface density of our sample, 0.38 ± 0.10 objects per 100 square degrees
out to 22 pc, corresponds to an uncorrected luminosity function of ΦJ = (1.76 ± 0.46).10−3
dwarfs/mag/pc3. After correcting for the Malmquist bias, this becomes ΦJcor = (1.64 ±
0.46).10−3 dwarfs/mag/pc3, averaged over 11.1≤ MJ ≤13.1. Our value is in good agreement
within error bars with the Cruz et al. (2007) measurement of ΦJ = (0.95 ± 0.30).10−3
dwarfs/mag/pc3, averaged over 11.25 ≤ MJ ≤13.25 for M8-L4 ultracool dwarfs at high
galactic latitude (|b| > 10◦). This clearly demonstrates that the space density of ultracool
dwarfs (M8-L3.5) in the solar neighborhood does not depend on galactic latitude.
47
CHAPTER 4
USING VLTI OPTICAL INTERFEROMETRY TO VERIFY
HST-ASTROMETRY DERIVED MASS AND ORBITAL
PARAMETERS OF THE COMPANION OF HD33636
4.1 BACKGROUND
In the past decade, radial velocity surveys have detected over 200 extrasolar planets with
masses Msin(i) < 13MJ (where MJ is the Jupiter mass). Even though the radial velocity
method provides initial detection, important orbital information and an estimate of the
minimum mass, the true mass of the companions depends on the inclination angle. For
this reason, high precision astrometry is required to further constrain the orbit, obtain the
inclination angle, and determine the true mass of the companions. High precision and long
time-baselines are two challenging and time consuming factors for using astrometry. Follow-
up astrometry to detect the reflex motion due to very low mass companions requires an
outstanding astrometric precision (∼ 1mas) for a time span of the order of the orbital
period.
48
The advantage of using optical interferometry over high precision astrometry is that we
can constrain the inclination angle of the binary and measure its true mass by observing it
in less than half a night as opposed to 1 year. The documented precision of 9% in visibility
measurements and 1% in closure phase measurements is ideal for such direct determination
of the flux-ratio and angular separation of binary systems. Using ASPRO (JMMC tool) and
VisCalc (ESO tool), we have modeled visibilities and closure phase signals for different incli-
nations for our target during the observing epoch and have estimated an absolute visibility
of about 0.9 and a closure phase amplitude of about 3◦.
The instrument sensitivity and precision place limits on the flux-ratios (and angular
separations) that could be detected. In all cases the highest detectable true-mass of these
companions is > 13MJ and hence are brown dwarfs or low-mass stars. It is important to note
that even a non-detection would give us an upper mass limit for these companions. Detection
of brown dwarfs at these small separations becomes a very significant discovery as these
companions would become some of the very few brown dwarfs in the brown dwarf desert.The
deficit of close (a < 4AU , a being the separation of the binary) BD companions to solar type
stars relative to stellar companions or planetary mass companions is commonly referred to
as the brown dwarf desert. The existence and extent of this brown dwarf desert can provide
important constraints on star and planet formation mechanisms since the BDs represent
extreme mass limits of both of the processes (Kraus et al. (2008)). For example, the BD desert
suggests a different spectrum of gravitational fragmentation in the formation environment.
The BD desert could also be explained by the disinclination of post formation migratory
49
mechanisms to leave BDs in close orbits (Grether and Lineweaver (2006), Ribas and Miralda-
Escude (2007)). Hence it is clear that discovery of these rare brown dwarfs is in itself a high
profile result.
Recently, speckle interferometry was used to detect close companions (6AU − 435AU)
to young stars in upper scorpius (Kraus et al. (2008)).Speckle interferometry has proven to
probe the high contrast close substellar companions deeper than conventional imaging meth-
ods (∼5 times smaller angular separations). With VLTI/AMBER + FINITO, we present
a case where we can detect companions with even higher contrast (flux ratio ≤ 10−4) and
smaller angular separations (≤ 100mas).
4.2 METHOD
Conventional imaging can be described mathematically as
I(l, m) =
Z ZP (l − l0, m − m0)O(l0, m0)dl0dm0, (4.1)
which implies that the observed brightness distribution I is the true source brightness distri-
bution O(l0, m0) convolved with a point-spread function, P(l,m) where l and m are angular
coordinates on the sky measured in radians. The Fourier transform of the above equation
looks like
I(u, v) = T (u, v)XO(u, v), (4.2)
50
where I(u,v), O(u,v) are the Fourier transforms of I(l,m) and O(l,m) respectively and T(u,v)
is the Fourier transform of the PSF. u and v are now the spatial frequencies measured in
radians−1
This shows the decomposition of an image into a series of spatially separated compact
PSFs. That also implies that if one can measure the Fourier components of the source, we
can get a lot of information on the source itself.
Figure 4.1 shows a schematic of a 2-element interferometer. The telescopes are located
at x1 and x2. The baseline B=x1-x2. This governs the sensitivity to different angular scales.
S is the pointing direction, and the geometric delay is s.B where s = S/|S|. d1 and d2
represent the the optical paths along the two arms.
The electromagnetic fields from the two collectors can be written as
™1 = Aeik(s.B+d1)e−iωt (4.3)
™2 = Aeikd2e−iωt (4.4)
∴ ™ = ™1 + ™1 = A(eik(s.B+d1) + eikd2)e−iωt (4.5)
I = h™™∗i ∝ 2 + 2cos(k[s.B + d1 − d1]) ∝ 2 × (1 + cos(kD)) (4.6)
where D = s.B + d1 − d2 This implies that the intensity of the combined light varies
cosinusoidally with kD with k = 2π/∏. In interferometry the key observables are the contrast
(VISIBILITY) and the location of these modulations (PHASE) in intensity.
V isibility, V =Imax − Imin
Imax − Imin(4.7)
51
Figure 4.1: A schematic representation of a 2-element interferometer.
52
In other words, the fringe amplitude and phase are a measure of the amplitude and phase
of the Fourier transform of the source at one spatial frequency dependent on the baseline
vector B. If the baseline is large the interferometer probes small scale structures and if the
baseline is short, it probes large structures. Further details can be found in Haniff (2007)
and references therein.
Astronomical Multi-BEam combineR (AMBER), the near-infrared/red focal instrument
of the VLTI, operates in the bands J, H, and, K (ie 1.0 to 2.4 µm). The instrument has been
designed to be used with two or three beams, allowing to measure the closure phase. Its
angular resolution is set by the maximum separation of the telescopes and the field of view of
the instrument. In short, AMBER is able to resolve features between 2mas (milli-arcsecond)
and 50mas with the 8.2m unit telescopes(UTs), and between 2mas and 140mas with the
1.0m auxiliary telescopes (ATs). AMBER is used with an external fringe tracker FINITO.
On the UTs it is possible to reach H=7. On the AT’s, the limiting magnitude is H=5.
We selected HD33636 as it is a bright star with H < 7. A known RV companion with
msini = 9.3MJ , Period P = 2128 days , eccentricity e = 0.48 and velocity semiampliude
K = 164ms−1(Perrier et al. (2003), Butler et al. (2006) ). Using the radial velocity measured
parameters, Bean et al. (2007) further constrained the orbit using Hubble Space Telescope
(HST) astrometry. The inclination was estimated to be 4.1 ± 0.1, degrees and the corre-
sponding mass was estimated to be 142 ± 11MJ . This relatively high mass companion made
HD33636 ideal for our pilot study as this would mean a smaller contrast ratio and detectable
visibility, phase, and closure phase signals.
53
Figure 4.2: AMBER data reduction work flow
HD33636 was observed with AMBER and FINITO combining light from three of the four
8.2 meter unit telescopes (UT1-UT2-UT4) in October 2008. AMBER data reduction was
performed using AMDLIB and MYAMBERGUI packages. An overview of AMBER data
reduction is given in the Figure 4.3. The data was reduced after choosing 20% of the frames
with the highest fringe SNR. The calibrator used was HD34137 with a diameter of 0.8 mas.
54
AMBER.2008−10−14T07:13:47.387_OIDATA_AVGTarget : HD33636Exposure time (s): 0.1000Observing date: 2008−10−14T07:13:47.3873Observations category: SCIENCEObservation type: OBJECTAir mass: 1.196Seeing ("): 1.04Coherence time (s): 0.001747Central wavelength (µm): −0.001Grating order: 0Instrumental mode: 3Tstd_Low_JHKSpectral shift 1: 4.218750Spectral shift 2: 3.781250Spectral shift 3: −0.781250BCD: OUTFINITO: ONTelescope 1: U1Telescope 2: U2Telescope 3: U4
0.2.4.6.8.10+6
0.2.4.6.8.10+6
1.5 2.0 2.50.2.4.6.8.10+6
0.000.050.100.15
0.000.050.100.15
1.5 2.0 2.50.000.050.100.15
A0A1
B0B1B2B3B4B5
C0C1C2C3
D0
D1D2
E0 G0
G1
G2H0
I1
J1
J2
J3J4
J5
J6
K0L0M0U1
U2
U3
U4
1 = U1−U2 : 47.43m, 22.05° 2 = U2−U4 : 83.78m, 83.23° 3 = U4−U1 : 114.4m, 61.94°
−100−50 0 50 100
−100
−50
0
50
100
−2−1 0 1 2
−2−1 0 1 2
−2−1 0 1 2
1.5 2.0 2.5−2
0
2
V2 1V2 2
V2 3
Wavelength (µm)
φ 1 (r
ad)
φ 2 (r
ad)
φ 3 (r
ad)
Wavelength (µm)
ψ (r
ad)
E <−−−−− U (m)
V (m
) −−−−−
> N
E <−−−−− U (m)
V (m
) −−−−−
> N
E <−−−−− U (m)
V (m
) −−−−−
> N
Flux
(e− )
Flux
(e− )
Flux
(e− )
Wavelength (µm)
Figure 4.3: One of the averaged frame of HD33636. This is the output of the program YORICK
which is a part of AMDLIB and was used for data reduction.
55
4.3 RESULTS
Interferometric observables visibility, phase, and closure phase were obtained (Figures 4.3,
4.4, 4.5, 4.6). The flux was derived by performing the inverse Fourier transform of the
visibility. To model our measurments we adopt a binary system with two point sources
as outlined by Berger and Segransan (2007). For clarity, we will derive the point source
binary model from the first principles here. Consider now an astrophysical object that
can be described by the addition of n components of known morphologies. Denote the
brightness distributions of such objects Ij(α, β) their position in the plane of sky being
(αj,βj) respectively and the corresponding normalized visibilities V (u,v) with j = 1 . . . n.
To compute the normalized visibility of such an object one should take into account their
respective contributions to the total brightness which we will name Fj. The total brightness
distribution can therefore be written:
I(α, β) =X
j=1..n
Ij(α, β)δ(α − αj, β − βj) (4.8)
The addition property of the Fourier transform allows us to write the visibility function
as:
∫(u, v) =X
j=1..n
FjV (u, v)exp(2πi(uαj + vβj)) (4.9)
Normalization gives the final visibility:
56
V (u, v) =
Pj=1..n FjV (u, v)exp(2πi(uαj + vβj))P
j=1..n Fj(4.10)
For a binary system with two point sources S1 and S2, with separation ρ, position angle
θ and respective fluxes F1 and F2,
I(α, β) = F1δ(α − α1, β − β1) + F2δ(α − α2, β − β2) (4.11)
where, (α1,β1) and (α2, β2) are the angular coordinates for sources S1 and S2. The
corresponding Fourier transform gives the unnormalized complex visibility function:
∫(u, v) = F1exp(2πi(uα1 + vβ1)) + F2exp(2πi(uα2 + vβ2)) (4.12)
The normalized squared visibility amplitude is then:
|V (u, v)|2 =∫(u, v)∫(u, v)∗
|∫(0, 0)|2 =F 2
1 + F 22 + 2F1F2cos(2π(u(α1 − α2) + v(β1 − β2)))
(F1 + F2)2(4.13)
where the normalization factor is the total flux squared (∫2(0, 0)). If we introduce the
flux ratio f=F 2F 1 , the baseline vector ~B(|B| = ∏
√u2 + v2), and the separation vector ~ρ(|ρ| =
p(α1 − α2)2 + (β1 − β2)2)), the above equation becomes
|V (u, v)|2 =1 + f2 + 2cos(2π/∏ ~B~ρ)
(1 + f)2(4.14)
We use the above equation to model our observed visibilities as shown in Figure 4.7.
Even though measurements of visibility at intermediate baselines to our observations are
57
required to break some model fit degeneracies, a value for the flux ratio could be estimated
from this model fit. We estimated a flux ratio of 0.05, which seems consistent with the
a mass of 0.1M⊙ − 0.2M⊙ for the companion. This flux ratio was estimated using the
radius from the mass-radius relation for stars, and the luminosity was estimated using the
theoretical evolutionary models . Some inconsistencies remain in the orbital parameters from
our observations when compared to the oribt as derived by Bean et al. (2007). Figure 4.8
(solid red) shows the projected (apparent) orbit of the companion of HD33636 as derived by
Bean et al. (2007). Our VLTI AMBER observations were performed in October 2009. At
that time the angular separation between the primary and the secondary was ∼ 100mas.
The AMBER field of view with UTs and FINITO ∼ 60mas. This implies that if Bean et al.
(2007) orbit is correct then we did not have the same companion in the field of view (see
Figure 4.9), and hence we are seeing some other source which could be one of the following:
• a background source very close the primary : Future interferometric observations can
rule-out this possibility as the background source will have a different parallax
• another companion much closer (at the time of observations) to the primary : This
companion has to be almost face-on i < 0.5degrees as no additional signal is detected
in the radial velocity of the star.
The other possible explanations for this inconsistency are that either Bean et al. (2007)’s
orbit is not accurate or there are some missing parameters in how we plot the orbit.
58
Base 2−3
1.6 1.8 2.0 2.2 2.40.0
0.5
1.0
Base 1−2
1.6 1.8 2.0 2.2 2.40.0
0.5
1.0
Base 1−3
1.6 1.8 2.0 2.2 2.40.0
0.5
1.0
V2V2
Wavelength (µm)
V2
Figure 4.4: Squared visibility vs wavelength plot. Base 1-3 represent the UT1-UT4; Base 1-2
represent UT1-UT2; and Base 2-3
59
1.5 2.0 2.50.2.4.6.8.
10+6
1.5 2.0 2.50.2.4.6.8.
10+6
1.5 2.0 2.50.2.4.6.8.
10+6
Flux
(e− )
Flux
(e− )
Flux
(e− )
Wavelength (µm)
Figure 4.5: Flux vs wavelength plot. The order is same as figure 4.4.
60
1.6 1.8 2.0 2.2 2.4−1.0
−0.5
0.0
0.5
1.0
Wavelength (µm)
Clos
ure
Phas
e (ra
d)
Figure 4.6: Closure phase vs wavelength plot.
61
Visibilities.pdf Visibilities.bb
0
0.1
0.2
0.3
0.4
0.5
0.6
0.7
0.8
0.9
1
0 20 40 60 80 100 120 140
Projected Baseline
V2
Figure 4.7: Model fit for the visibility measurements. The black squares represent our data and
the magenta diamonds represent model parameters for flux ratio=0.05 and angular
separation=30mas.
62
-80
-60
-40
-20
0
20
40
60
80
100
120
140
-120 -100 -80 -60 -40 -20 0 20 40 60
!" (
ma
s)
!# (mas)
Oct 2008
Aug 2009
Oct 2009
Jan 2010
N
E
MB = 143 Mjup
Figure 4.8: Apparent orbit of HD33636 and its companion.
63
20
40
60
80
100
120
140
160
0 20 40 60 80 100 120 140 160 180
An
gu
lar
se
pa
ratio
n (
ma
s)
Orbital inclination (deg)
Oct 2008
Aug 2009
Oct 2009
Jan 2010
Figure 4.9: Angular separation vs Orbital inclination of the companion using the RV derived
parameters only for different epochs. The solid horizontal black line represents the
AMBER+UT field of view. The colored solid curves represent the angular separation
for all possible values of i for different epochs.
64
CHAPTER 5
OPTICAL LINEAR POLARIZATION IN ULTRA COOL
DWARFS: A TOOL TO PROBE DUST IN ULTRA COOL
DWARF ATMOSPHERES
5.1 BACKGROUND
A large number of UDs have been detected in the last decade, and our understanding of
these faint objects has kept improving. One of the challenging and fundamental aspects in
the study of these objects is to understand the properties and distribution of condensate
dust in the atmosphere. Observations of L dwarfs with effective temperatures of 1400-2200
K have led to the investigation of dust condensates in their atmospheres (Kirkpatrick et al.
(1999); Tsuji et al. (1996)). Because of complete gravitational settling, grains are expected
to condense beyond the visible atmosphere for objects with effective temperatures below
1400 K (T-Dwarfs - Allard et al. (2001); Chabrier et al. (2000)). At higher effective temper-
atures (1400-2200 K), grains can be present in the visible atmosphere because of incomplete
gravitational settling (Burrows and Sharp (1999); Burrows et al. (2001); Ackerman and
Marley (2001a); Allard et al. (2001); Tsuji et al. (2004); Cooper et al. (2003); Helling et al.
65
(2008) ). Recent discoveries of blue L dwarfs and L-T transition type dwarfs (as identified
in Knapp et al. (2004); Chiu et al. (2006); Tsuji and Nakajima (2003)) have brought forth
models which could explain this phenomenon (eg. Burrows et al. (2006), Knapp et al. (2004))
by mechanisms which involve dust settling. It would be very important to validate these
mechanisms.
Linear polarization could be a very useful tool in understanding the observationally poorly
constrained dust properties in the atmospheres of L dwarfs. The possibility of detecting
polarization at optical wavelengths from grains in the atmospheres of L dwarfs was first raised
by Sengupta and Krishan (2001). Fast rotation of L dwarfs will induce the shape of their
photosphere into the form of an oblate ellipsoid,(Basri et al. (2000)) and this nonsphericity
will lead to the incomplete cancellation of the polarization from different areas of the stellar
surface (Sengupta and Krishan (2001)). This prediction was first confirmed by the detection
of linear polarization at 768 nm from a few L dwarfs by Mnard et al. (2002). Recently,
Zapatero Osorio et al. (2005) have reported R and I band detection of linear polarization from
several L dwarfs. Since polarization in the optical is unlikely to be due to Zeeman splitting
of atomic or molecular lines or by synchrotron radiation, the observed polarization can be
explained by single dust scattering in a rotationally induced oblate atmosphere (Sengupta
(2003); Sengupta and Kwok (2005)) or it could be due to large and randomly distributed
dust clouds (Mnard et al. (2002)).
66
5.2 METHOD
Four very nearby (7pc -15pc) UDs were selected (SpT M9.5 - L5) as they are among the
brightest and nearest UDs with no known infrared excess and no evidence of multiplicity(see
Table 5.1). This selection criteria ensures that the targets are bright enough sources in R
band to get high S/N and to avoid other than intrinsic sources of polarization such as cir-
cumstellar disks or multiplicity. For calibration, one polarized (Cyg OB2 A Whittet et.al.
1992) and one unpolarized standard star were observed at two different times during the
night. All the objects were observed so that they were acquired at the same position on
the detector (5 pixel box). This procedure minimized contamination caused by instrumental
polarization within the detector and variations in the optical path.
The polarimetric observations were obtained using The Intermediate dispersion Spectro-
graph and Imaging System (ISIS) which is mounted at the Cassegrain focus of the 4.2m
William Herschel Telescope (located in La Palma, Canary Islands, Spain). ISIS in polariza-
tion mode is a modulation polarimeter with a double-beam analyzer (the calcite plate) and
a rotating halfwave plate modulator. ISIS is equipped with two detectors: a blue-sensitive
EEV12 (4096x2048 pixels) and a red-sensitive RED+ (4096x2048) detector. In our program,
we have used the RED+ detector.
67
Images were obtained using Bessel R- and I- filters centered on 641 and 812 nm, respec-
tively, on June 18, 2006 (UT Date). The night was photometric with stable average seeing
of 1.0 arcsec .
The raw images were bias-subtracted and flat-fielded before performing aperture photom-
etry. The flat-field images were obtained with the polarimeter optics. One of the reduced
frames is shown in Figure 5.1.
Fluxes were obtained for 0.8, 1.0, 1.2, 1.5, 2.0 times the average FWHM for each object.
The best aperture was chosen to be 1.5 times FWHM based on minimum photon contri-
bution of nearby sources, variable sky contribution, and maximum signal-to-noise ratio of
the measurements. The average FWHM of all images was 4.0 pixels which correspond to
1.0 arcsec. We have only one set of measurements for each object. Therefore, we have esti-
mated the uncertainty in degree of polarization from various apertures. (a similar method
was used by Zapatero Osorio et al. (2005) for some of the objects) There was no signifi-
cant instrumental polarization found as the unpolarized standard measured D(p)=0.086% ±
0.002 . Polarization is a measure of anisotropy in the radiation field and is caused by either
scattering or is due to the presence of magnetic field. The state of the polarization of light
is described by the Stokes parameters, I, Q, U and V. The parameter I is the total scalar
specific intensity of radiation. It is the complete flux of radiant energy inside the unit inter-
vals of frequency, time, solid angle, and area perpendicular to the flux. This flux includes
all radiation independently of polarization. Polarization is described by the parameters Q,
U, V. These parameters are proportional to the scalar specific intensity and have the same
68
Figure 5.1: Reduced data frame. The black dots in the center of the figure represent the ordinary
and the extraordinary component of the source. The vertical lines are from a mask
which was used.
69
dimension. Q and U represent the linearly polarized component, and V represents the cir-
cularly polarized component. For linear polarization, V=0 and the degree of polarization is
given as p =p
Q2 + U2/I. If we consider axial symmetry, then U = 0 and in that case we
define the degree of polarization p = −Q/I. The sign convention is such that if p > 0, the
light is polarized perpendicular to the scattering plane, and if p < 0, the light is polarized
parallel to the scattering plane. For an unresolved stellar object, the Stokes parameters are
integrated over the stellar disk.
From our obtained images, the degree of polarization and the polarization angles are
calculated using the following equations:
R2Q =
o(0◦)/e(0◦)
o(45◦)/e(45◦); Q/I =
RQ − 1
RQ + 1(5.1)
R2U =
o(22.5◦)/e(22.5◦)
o(67.5◦)/e(67.5◦); U/I =
RU − 1
RU + 1(5.2)
p =p
(Q/I)2 + (U/I)2 (5.3)
θ = 0.5tan−1 U/I
Q/I(5.4)
where
o, e ordinary and extraordinary fluxes respectively of the dual images of the
program source on a single frame
p degree of polarization
θ polarization angle
I the total scalar specific intensity of radiation
Q/I, U/I normalized Stokes parameters
70
As pointed out by Mnard et al. (2002), the observed linear polarization in the optical can-
not be due to magnetic field, and scattering remains the most viable mechanism for yielding
the detected linear polarization. Polarimetric observation at the R and I bands by Zapatero
Osorio et al. (2005) shows that polarization decreases significantly with the increase in wave-
length, which strongly supports this argument (Sengupta and Kwok (2005)). In the present
investigation, we report detection of polarization at both R and I bands which shows the
same wavelength dependency and hence strengthens the case for scattering polarization. If
the dust density is low, then the single scattering approximation is reasonable for the region
where the dust optical depth τd < 1 because scattering by atoms and molecules does not sig-
nificantly contribute to polarization. However, in the optical region τd is much higher than 1,
and therefore multiple scattering is the appropriate process. Multiple scattering reduces the
amount of polarization significantly because the planes of scattering are oriented randomly
and cancel each others’ contributions. In this work, however , we adopt a simple single scat-
tering polarization model because the main aim is to explain the wavelength dependency
of the observed polarization. Nevertheless, we find that a single dust scattering model in
rotation-induced oblate L dwarfs can fit the observed data well within the permissible range
of parameter space of dust composition, size, geometrical distribution and rotation period.
A detailed modeling with multiple scattering is in progress.
The simple theoretical model adopted here to explain the observed polarization is de-
scribed in details in Sengupta and Kwok (2005). At an edge-on view, the degree of polar-
ization integrated over the stellar disk is given by :
71
p(∏) =∏2
g
Z Pbot
Ptop
n(P )dP
ρ(P )
1X
l=2
Ωα2(l, m)P m
l (0)Fl2
Z 1
−1
Pl(µ)
[1 + (A2 − 1)µ2]1/2dµ
æ(5.5)
In the above expression, P is the total pressure (gas plus dust), Ptop and Pbot are the
pressures at the top (deck) and the bottom (base) of the dust layer, n(P ) is the grain
number density at different pressure heights, ρ(P ) is the total density at any pressure level,
and A is the ratio of the equatorial radius to the polar radius. Since the dust density is small
compared to the gas density, ρ(P ) is the effective gas density and is equal to nRT/P where
n is the mean molecular weight and R is the gas constant. In the above equation, Pl(µ) is
the Lengendre polynomial where µ = cos θ, θ being the scattering angle and Fl2 is given by
Flm = α(l, m)
Z 1
−1
i1 − i2
2P m
l (cos θ)d(cos θ). (5.6)
where
α(l, m) =
∑(2l + 1)(l − m)
4π(l + m)
∏1/2
, (5.7)
and P ml is the associated Legendre function of the first kind. i1 and i2 are the scattering
functions given by van de Hulst (1981).
The vertical dust distribution and the location of the cloud base and deck in the atmo-
sphere are calculated based on the one dimensional heterogeneous cloud model of Cooper
et al. (2003). This model assumes chemical equilibrium throughout the atmosphere and
uniform density distribution across the surface of an object at each given pressure and tem-
perature. The number density of cloud particles in this model is given by
72
n(P ) = qc
µρ
ρd
∂ µµd
µ
∂ µ3
4πa3
∂, (5.8)
where ρ is the mass density of the surrounding gas, a is the cloud particle radius, ρd is the
mass density of the dust condensates, µ and µd are the mean molecular weight of atmospheric
gas and condensates respectively. The condensate mixing number ratio (qc) is given as
qc = qbelowPc,l
P(5.9)
for heterogeneously condensing clouds. In the above equation, qbelow is the fraction of con-
densible vapor just below the cloud base, Pc,l is the pressure at the condensation point.
A log-normal size distribution is adopted for the spherical grains (Ackerman and Marley
(2001b)). The formation of dust makes it a prohibitive task to develop a fully consistent
atmospheric model for ultra-cool dwarfs. This is mainly because of the fact that the pres-
ence of dust clouds affects the radiative equilibrium of the upper atmosphere and hence
alters the T-P profile from that of a cloud-free atmosphere. On the other hand, the T-P
profile dictates the position and the chemical equilibrium of condensates. Allard et al. (2001)
presented atmospheric models for two of the limiting cases, e. g., one with inefficient gravita-
tional settling wherein the dust is distributed according to chemical equilibrium predictions
(AMES-dusty) and another with efficient gravitational settling in which situation dust has
no effect on the thermal structure (AMES–cond). Tsuji et al. (2004) have proposed a Uni-
fied Cloudy Model (UCM) in which the segregation of dust from the gaseous mixture takes
73
place in all the ultra-cool dwarfs and at about the same critical temperature. Ackerman
and Marley (2001b) treat the upward convective mixing of a gas, its condensation and the
sedimentation of the condensate through the atmosphere of the object while Woitke and
Helling (2004) consider an ensamble of dust grains falling downwards from the top of the
atmosphere. A detailed comparison of different atmospheric models of L dwarfs is presented
in Helling et al. (2008).
The oblateness of a rotating object has been discussed by Chandrasekhar (1933) in the
context of polytropic gas configuration under hydrostatic equilibrium. For a slow rotator, the
relationship for the oblateness f of a stable polytropic gas configuration under hydrostatic
equilibrium is given by
f =2
3C
≠2R3e
GM, (5.10)
where M is the total mass, Re is the equatorial radius, and ≠ is the angular velocity of
the object which is related to the linear velocity V = ≠ × Re. C is a constant whose value
depends on the polytropic index. For a polytropic index of n = 1.0, C = 1.1399, which is
appropriate for Jupiter (Hubbard (1984)). For non-relativistic completely degenerate gas,
n = 1.5 and C = 0.9669.
The effective temperature of the L dwarfs of different spectral type is determined by
adopting a sixth order polynomial fit given by Golimowski et al. (2004) which is based on
bolometric luminosities. The Teff calibration of Golimowski et al. (2004) agrees well in the
interval L3-L8, but there are significant differences in earlier types. In our calculations for
74
the degree of polarization, the effective temperature Teff is used and hence the degree of
polarization should be considered strictly as a function of Teff rather than of spectral type.
The mass and radius of the L dwarfs of different spectral types are estimated by adopting
the empirical relationship given by Marley et al. (1996).
75
Tab
le5.
1:Ta
rget
list
Ob
ject
RA
DE
CS
pTR
IJ
dH
αre
f
(mag
)(m
ag)
(mag
)(p
c)E
W(A
)
2MA
SSW
J143
8082
+64
0836
1438
08.2
+64
0836
M9.
5..
.16
.812
.92
7.2
±0.
5−
2.36
1
2MA
SSW
J150
7476
-162
738
1507
47.6
9-1
627
38.6
L5
18.9
16.5
12.8
7.33
±0.
03−
5.01
3,4,
5
2MA
SS
J173
1297
4+27
2123
317
3129
.74
+27
2123
.3L
0..
...
.12
.09
11.8
0±
0.70
−5.
981,
2
2MA
SSI
J180
7159
+50
1531
1807
15.9
3+
5015
31.6
L1.
5..
...
.12
.96
14.6
±1.
0−
1.38
1,2,
3
ref-
1.Sch
mid
tet
al.
(200
7),
2.Ja
mes
onet
al.
(200
8),
3.C
ruz
etal
.(2
003)
,4.
Rei
det
al.
(200
0),
5.K
nap
pet
al.
(200
4)H
-alp
ha
EW
from
Sch
mid
tet
al.
(200
7)
76
5.3 RESULTS
We find a trend (more data are required to confirm our theory) of higher polarization in the
R band when compared to the I band. This wavelength dependency strongly supports the
argument by Sengupta and Kwok (2005) that the polarization arises due to scattering and
not because of magnetic field. In dust scattering as described by Mie theory, the amount
of polarization depends on the ratio of the grain radius to the wavelength. For the same
kind of dust species, the polarization usually peaks when the ratio is one. As a consequence,
the increase in polarization with the decrease in wavelength implies the presence of sub-
micron size grains in the photosphere of the L dwarfs. We present our measurements in
Table 5.2, and our model fit in Figure 5.2. One of the L dwarfs from Zapatero Osorio et al.
(2005) (2MASSW J1507476-162738) shows a null polarization in our measurements in the
I band, whereas Zapatero Osorio et al. (2005) present a higher polarization (1.36% ±0.30)
for the same object. This suggests variability in linear polarization which in turn suggests
atmospheric activities like dynamical variations of the cloud cover.
We also find relatively high polarization in the L0 dwarf. Additionally, from Schmidt
et al. (2007), the Hα equivalent width is the highest for the L0 dwarf among the objects
from our sample. This could be an indirect evidence of a disk around this UD. We therefore
searched the Spitzer public archive for mid-IR data. 2MASS J17312974+2721233 has been
observed with IRAC and IRS in the course of program 3136 (P.I. Cruz), and we retrieved the
pipeline processed data. We extracted the IRAC photometry using standard PSF photom-
77
Table 5.2: Polarization measurements
Object SpT Filter Exposure Time D(p) θ
(seconds) (%) (degrees)
2MASSW J1438082+640836 M9.5 I 90 0.482 ±0.071 66.24
R 300 0.54 ±0.097 115
2MASSW J1507476-162738 L5 I 150 0.0 ±0.036 . . .
R 300 0.216 ±0.022 100.26
2MASS J17312974+2721233 L0 I 60 5.195 ± 0.854 52.57
R 240 0.666 ± 0.169 111.08
2MASSI J1807159+501531 L1.5 I 60 0.711 ± 0.142 46.67
R 120 1.669 ± 0.605 101.2
The uncertainty in θ due to calibration errors is about 1.2 degrees
etry procedures within the Interactive Data Language. Uncertainties were estimated from
the Poisson noise weighted by the coverage maps of the mosaics. Table 5.4 gives a summary
of the photometry. Figure 5.3 shows the spectral energy distribution (SED) of the source
and a L0 comparison object from the literature (2MASS J1204+3212, Patten et al., 2006).
2MASS J17312974+2721233 does not show any significant mid-IR excess up to 15 µm. The
presence of a young circumstellar disc can therefore be ruled out at a high level of confidence.
In the current state of the data, we cannot rule out the presence of a cold debris disc, as it
would produce an excess at longer wavelengths.
78
Figure 5.2: Best model fit of the observed data. The solid lines represent the model with the poly-
tropic index n=1.0 and dashed lines represent that with n=1.5. For other parameters
see Table 5.3.
79
1 10λ [µm]
10-15
10-14
10-13
Flux
[W m
-2]
2M1731+27212M1204+3212 (L0)
Figure 5.3: Spectral energy distribution of 2MASS J17312974+2721233 (dots). V-band photom-
etry from the LSPM-North proper-motion catalog nearby stars (Lepine et a., 2005).
J,H and Ks photometry from 2MASS. IRAC and IRS photometry from this work. A
comparison L0 dwarf (2MASS J1204+3212, Patten et al., 2006) scaled to the same
J-band flux is overplotted (squares). The target does not show any significant excess
compared to the field L0 dwarf.
80
Table 5.3: Model fit
Object SpT n d0 log(g) V
(micron) (cgs) (km s−1)
2MASSW J1438082+640836 M9.5 1.0 0.7 5.0 18.7
1.5 0.7 5.0 20.5
2MASSW J1507476-162738 L5 1.0 0.44 5.255 27.2
1.5 0.44 5.230 27.2
2MASSI J1807159+501531 L1.5 1.0 0.5 5.41 76.0
1.5 0.44 5.37 76.0
Bailer-Jones 2004 has measured the projected rotational velocity for 2MASSW J1507476-
162738 as 27.2 km s−1, while Reiner & Basri 2008 have reported the projected velocity for
2MASSW J1807159+501531 to be 76 km s−1.
Table 5.4: IRAC photometry of 2MASS J17312974+2721233
Wavelength Flux
(micron) (mJy)
3.6 21.82±0.02
4.5 14.16±0.02
5.8 11.30±0.06
8.0 6.91±0.03
81
The surface gravity of L dwarfs older than approximately a few hundred million years
varies from g = 105 to 3 × 105 cms−2. Evolutionary models by Chabrier et al. (2000) show
younger L dwarfs to have a surface gravity smaller than 105 cms−2. Our best fit model
parameters are presented in Table 5.3 (best fit obtained by visual inspection). For the L
dwarfs 2MASSW J1438082+640836, we find the best fit with log(g)=5.0, and rotational
velocity V=18.7 km s−1 when n = 1.0, and V=20.5 km s−1 when n = 1.5. For a fixed
rotational velocity, the oblateness is lower when the polytropic index n is higher. As a
result, higher rotational velocity is needed to fit the observed data if n = 1.5. The projected
rotational velocity of the L dwarf 2MASSW J1507476-162738 is measured by Bailer-Jones
(2004) while the projected rotational velocity of 2MASSW J1807159+501531 is estimated
by Reiners and Basri (2008). In the absence of any knowledge on the projection angle, we
consider the minimum rotational velocity of these objects as V=27.2 km s−1 and V=76 km
s−1. We find the best fit for the observed polarization of 2MASSWW J1507476-162738 with
Log(g)=5.255 for n = 1.0 and Log(g)=5.230 for n = 1.5 For 2MASSW J1807159+501531
the best fit is obtained with log(g)=5.41, when n = 1.0, and log(g)=5.37 when n = 1.5.
Note that the oblateness decreases with the increase in surface gravity. For all the cases,
the observed polarization profiles can be fitted with sub-micron size grains, and the mean
size of grains that are required to fit the observation is consistent with the recent theoretical
calculations of dust properties (Woitke and Helling (2004); Woitke and Helling (2003)).
Polarization measurements for one of the above three objects (2MASSW J1507476-162738)
were also recently published Goldman et al. (2009). Our results are consistent with the
82
Goldman et al. (2009) measurements within 1σ error bars for this object. This work differs
from Goldman et al. (2009) in the aspect that our sample could reproduce the predictions of
Sengupta and Kwok (2005) whereas the Goldman et al. (2009) sample does not. Both the
studies (this work and Goldman et al. (2009)) have sparse statistics, which discourages us
to make any speculations on this apparent discrepancy.
To summarize,
1. We report linear polarizaion measurements of 4 very nearby UDs in the R and I bands.
2. We find that there is a trend (3 out of 4) of a higher degree of polarization at shorter
wavelengths (R band) when compared to the I band as predicted by the theoretical
models of Sengupta and Kwok (2005).
3. The L0 dwarf 2MASS J17312974+2721233 is interesting because of its relatively high
polarization and requires follow-up studies.
4. We also fit theoretical models to predict the dust grain size and rotational velocities
of three of the UDs.
5. We find evidence for variability in the linear polarization for (2MASSW J1507476-
162738). This suggests atmospheric activities like dynamical variations of the cloud
cover in this object.
83
CHAPTER 6
DISCUSSION
In absence of observations, the field of study of sub-stellar objects was led by theorists alone.
The mere discovery of these objects, in a sense, vindicates the scientific process of research.
Taking it a step further one can say that if some phenomenon is permissible by nature (or
say physics), it is possible that it had occurred somewhere in this universe. It then becomes
a matter of time and technology that we either discover it or find why that phenomenon is
forbidden.
In this dissertation, I have presented a wide range of projects aimed to investigate a wide
range of physics involving sub-stellar objects. From their search in different regions and
environments (star-forming regions- Chapter 2; low-galactic latitude- Chapter 3; in the field-
Chapter 4) to their spectral characterization and their physical and atmospheric properties
(Chapter 5). We investigate the brown dwarf desert at wide and small separations (Chapter
2 and Chapter 4, respectively). Atmospheres of these objects have been well modeled, we
present a way to use observations to validate some of these models (Chapter 5).
To summarize our key results, the search for wide binaries around solar type stars in
Upper Sco OB association (Upper Sco) do indicate (the survey is not yet complete) a deficit
84
of BD binaries at these large separations (< 5AU). Twenty-six new UDs were discovered at
low galactic latitudes in our survey from archival data and a novel technique using reduced
proper motion. Six field UDs were discovered by spectroscopic follow-up of the candidates
selected from the deep survey (SURF). Optical interferometry was used to independently
determine the orbit of the companion of HD33636 which was initially determined using
Hubble Space Telescope(HST) astrometry and radial velocity found. Some inconsistency
in the HST-determined orbit and mass. Optical linear polarization in UDs was used to
investigate the dust propertied in their atmospheres. A trend in polarization as predicted by
theoretical models was validated, and atmospheric dust grain sizes and projected rotational
velocities for these objects were estimated using the model fitting with the data.
Even though in Chapter 2 our study is not complete (which will require follow-up spec-
troscopy of all of the candidates), if we do indeed see a deficit of BDs at wide separations,
that would have the following implications on BD formation theories (at least in Upper Sco).
• Star formation efficiency could be mass dependent which could also explain the dis-
continuity in the IMF close to the sub-stellar boundary.
• Ejection of brown dwarfs (single or binaries) after they form at wide separations (<
100AU) in the cooler part of the disc through fragmentation. The ejection could be
caused due to dynamical interactions with the surrounding objects. This scenario
could not explain the BD desert at close separations where the they would be more
tightly bound. Low-density formation environments like the TW Hydrae association
85
and MBM 12, where disruption due to interactions is less significant, appear deficient
in brown dwarf companions (McCarthy 2001).
• We have 4 candidates which are red and over luminous and could be BD-BD binaries.
If they are proven to be BD-BD binaries, that would be a good support of the frag-
mentation of the outer disc formation theory of the BDs. BDs and BD-BD binaries
being ejected from their initial formation environment would seem more viable.
In Chapter 3 we demonstrated a novel and efficient technique to find ultra cool dwarfs in
the solar neighborhood in the galactic plane. The implications and contribution to the field
of brown dwarf research of this study could be listed as follows.
• More than 20 new ultra cool dwarfs discovered in the solar neighborhood which provide
us with additional targets for detailed studies of the basic physical properties of stars
at the bottom of the main sequence.
• A novel technique is demonstrated successfully to discover UDs in the galactic plane in
the solar neighborhood. Similar techniques can be used with some of the other surveys
to discover more UDs.
• We found all but one UDs in our sample were not Hα emmitors. This suggests that
the dissipation of chromospheric activity in our ultracool dwarfs is due to their pre-
dominantly cool, dense, and highly neutral atmospheres.
86
• A space density of 1.64 ± 0.46x10−3dwarfs/pc3/mag averaged over 11.1 < MJ < 13.1
for low galactic latitude is derived. This value is in good agreement (within error bars)
with Cruz et al. (2007) measurement of 0.95 ± 0.3 for M8-L4 UDs at high galactic
latitude. This suggests that the space density of UDs in the solar neighborhood is not
dependent on the galactic latitude.
In Chapter 4 we demonstrated another novel technique in conjunction with the radial
velocity method to obtain orbital parameters, mass, and temperature of close companions to
bright solar-type stars. The contribution and implications of this study include the follwing.
• We detect a relatively faint source at ∼30mas from the star. This could be another
companion in addition to the RV detected companion. The flux signal from this com-
panion looks like a dusty late L type dwarf (high resolution dispersed interferometric
measurements are required to confirm the spectral type). If this is true then we present
one of the few direct observations of BDs in the BD desert. Moreover getting a direct
spectrum of a companion this close and faint in itself is a high profile result.
• We demonstrate the use of AMBER+FINITO+UTs for similar studies which may
require the measurement of very small visibility and closure phase signals. Once the
use of AMBER+FINITO+UTs/ATs is established to obtain/estimate the true mass,
this method can be used to effectively filter out rueexoplanet candidates for future
astrometric programs like PRIMA and SIM.
87
In Chapter 5 we used linear polarization in the optical in the imaging mode and were
able to demonstrate that these measurements can be used as observational tools for probing
dust and other physical properties of ultra cool dwarfs.
• We validate theoretical models by Sengupta and Kwok (2005) using our measurements.
• Our findings add to the evidence that L dwarf photospheres are dusty and the effective
dust grain is sub-micron size.
• We also find evidence for variability which suggest a dynamic cloud cover in L dwarfs.
• Brown dwarf atmosphere studies are also very important as they are the contributors
to the bulk properties which are observable. Moreover, the similar properties of cool
dwarf and hot exoplanet atmospheres (low temperatures, molecule-rich atmospheres,
presence of clouds and dynamics) but distinct approaches to their study (e.g, direct
observations versus transit detection and phase variation) indicate opportunities for
synergistic efforts in addressing common, outstanding problems. The advantage of UD
atmosphere study is the ease of detailed, direct observation.
Observational astronomy in itself is a very diverse field. This dissertation boasts the use of
most of the techniques of observational astronomy like photometry, spectroscopy, polarimetry
and interferometry. Each of these techniques required different telescopes, instruments, and
data reduction and analysis methods. Some of these techniques required theoretical modeling
for the interpretation of the data. In conclusion, this dissertation shows a synoptic study
using different observational techniques and theoretical modeling, of UDs.
88
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