11institutetext: Mark Marley 22institutetext: NASA Ames Research Center, Mail
Stop 245-3, Moffett Field CA 94035, 22email<EMAIL_ADDRESS>33institutetext: S. Leggett 44institutetext: Gemini Observatory, 670 North
A’ohoku Place Hilo, HI 96720 44email<EMAIL_ADDRESS>
# The Future of Ultracool Dwarf Science with JWST

Mark S. Marley and S.K. Leggett

###### Abstract

Ultracool dwarfs exhibit a remarkably varied set of characteristics which hint
at the complex physical processes acting in their atmospheres and interiors.
Spectra of these objects not only depend upon their mass and effective
temperature, but also their atmospheric chemistry, weather, and dynamics. As a
consequence divining their mass, metallicity and age solely from their spectra
has been a challenge. JWST, by illuminating spectral blind spots and observing
objects with constrained masses and ages should finally unearth a sufficient
number of ultracool dwarf Rosetta Stones to allow us to decipher the processes
underlying the complex brown dwarf cooling sequence. In addition the spectra
of objects invisible from the ground, including very low mass objects in
clusters and nearby cold dwarfs from the disk population, will be seen for the
first time. In combination with other ground- and space-based assets and
programs, JWST will usher in a new golden era of brown dwarf science and
discovery.

##  1 Introduction

The explosive growth of brown dwarf and ultracool dwarf discoveries over the
past dozen years has been so extraordinary that it is a rare paper in the
field that does not open by remarking upon it. The first undisputed brown
dwarf, Gl 229 B, was discovered as a companion to an M dwarf in 1995 (Nakajima
et al. 1995). The Two-Micron All Sky Survey (2MASS, Skrutskie et al. 2006))
and the Sloan Digital Sky Survey (SDSS, York et al. 2000) subsequently
revealed large numbers of ultracool low-mass field dwarfs. These surveys first
led to the discovery of the isolated field late-L dwarfs (Kirkpatrick et al.
1999), then the mid-T dwarfs (Burgasser et al. 1999, Strauss et al. 1999) and
finally the early-T dwarfs (Leggett et al. 2000). Today over 600 warm ($T_{\rm
eff}\sim 2400$ to 1400 K) L and cool (600 to 1400 K) T dwarfs are known111See
http://www.DwarfArchives.org and the quest for the elusive, even cooler “Y”
dwarfs is ongoing.

Note that collectively late M and later type dwarfs are often termed
‘Ultracool Dwarfs’ (or UCDs) to avoid having to distinguish whether particular
warm objects in this group are above or below the hydrogen burning minimum
mass, the requirement for bestowing the term ‘brown dwarf’. Brown dwarfs will
continuously cool over time. The more massive UCDs will eventually arrive on
the bottom of the hydrogen burning main sequence (Burrows et al. 1997).

Figure 1: The most prominent signatures of the ultra cool dwarf spectral
sequence are seen in these 0.65 to $14.5\,\rm\mu m$ spectra of a mid-M, L, and
T dwarfs as well as Jupiter (adapted from Cushing et al. (2006)). The spectra
have been normalized to unity at $1.3\,\rm\mu m$ and multiplied by constants.
Major absorption bands are marked. The collision-induced opacity of $\rm
H_{2}$ is indicated as a dashed line because it shows no distinct spectral
features but rather a broad, smooth absorption. Jupiter’s flux shortward of
$\sim 4\,\rm\mu m$ is predominantly scattered solar light; thermal emission
dominates at longer wavelengths (near- and mid-infrared Jovian spectra from
Rayner, Cushing & Vacca (in preparation) and Kunde et al. (2004),
respectively).

Ultracool dwarf science is exciting not only for the rapid pace of discovery,
but for a host of other reasons as well. First, since brown dwarfs lack an
internal energy source (beyond a brief period of deuterium burning), they cool
off over time; they thus reach effective temperatures below those found in
stars and enter the realm where chemical equilibrium favors such decidedly
‘unstellar’ atmospheric species as $\rm CH_{4}$ and $\rm NH_{3}$. Along with
these more typical ‘planetary’ gasses, silicate and iron clouds are found in
their atmospheres, leading to interesting, complex interactions between
atmospheric chemical, radiative-transfer, dynamical, and meteorological
processes. Ultracool dwarfs thus bridge the domain between the bottom of the
stellar main sequence and giant planets. They are a laboratory for
understanding processes that will also be important in the characterization of
extrasolar giant planets. Second, they occupy the low mass end of the stellar
initial mass function. Understanding the IMF requires that we understand the
masses of individual field objects, which ultimately requires an understanding
of their luminosity evolution as well as the dependence of their spectra on
mass, gravity, effective temperature and metallicity. Finally, as terra
incognita, brown dwarfs (some of our nearest stellar neighbors) have offered a
series of surprises that test our ability to understand the universe around
us.

Figure 2: Important chemical equilibrium boundaries for substellar objects
(modified from Lodders & Fegley 2006). Green, red, and blue lines denote
various condensation boundaries for a solar abundance mixture of gasses in a
substellar atmosphere. Light purple lines denote equilibrium boundaries
between important gaseous species. Grey dashed lines show model atmospheric
temperature-pressure profiles for M, L and T dwarfs (the latter specifically
for Gl 229 B) as well as for Jupiter. As one moves upwards in the diagram
along the model $(T,P)$ curves, the labeled species will condense at the
intersection with the condensation curves and would be expected to be absent
from the gas at lower temperatures further up along the model profiles. This
figure can be compared with Figure 1 to understand why spectral features for
various compounds are present or absent in each observed spectrum.

The most distinctive features of the UCD spectral sequence are highlighted in
Figure 1. At effective temperatures below those of late-M dwarfs, several
chemical changes (illustrated in Figure 2) occur that strongly impact the
spectral energy distribution. First, major diatomic metal species
(particularly TiO and FeH) become incorporated into grains, leading to the
gradual departure of hallmarks of the M spectral sequence (Kirkpatrick et al.
1999). Second, the formation of iron and silicate grains produces optically
thick clouds that veil gaseous absorption bands and redden the near-IR $JHK$
colors of L dwarfs. The atmospheric temperature domain where these clouds are
most important is $\sim 1500-2000\,\rm K$ (e.g. Ackerman & Marley 2001). At
lower $T_{\rm eff}$, the clouds lie near or below the base of the wavelength-
dependent photosphere, and only marginally affect the SEDs of T dwarfs.
Finally, CH4 supplants CO as the dominant carbon-bearing molecule. This
transition is first noted in the 3–4 $\mu$m spectra of mid-L dwarfs (Noll et
al. 2000) and appears in both the $H$ and $K$ bands of T0 dwarfs (Geballe et
al. 2002). Together, increasing CH4 absorption and sinking cloud decks cause
progressively bluer near-IR colors of T dwarfs. For types T5 and later,
significant collision-induced H2 opacity in the $K$ band enhances the trend
toward bluer near-infrared colors.

These changes in spectral features are used to assign spectral types to UCDs
as briefly explained in §2. Given assigned spectral types, measurement of the
bolometric luminosity of individual objects along with their parallaxes
connect the spectral sequence to effective temperature. Figure 3 illustrates
the effective temperature as a function of spectral type from late M through
late T. While the general correlation of increasing spectral type with falling
effective temperature is unmistakable, a remarkably rapid set of spectral
changes (as expressed in the variation in spectral type) happens over a
relatively small span of $T_{\rm eff}$ near 1400 K. As we will discuss,
understanding this variation in expressed spectral signatures, the ‘L to T
transition’, is a key subject of current brown dwarf research.

Figure 3: Effective temperature as a function of infrared spectral type for
ultracool dwarfs with known parallax (data from Golimowski et al. 2004, Vrba
et al. 2004, and Luhman et al. 2007). Note the roughly constant effective
temperature for dwarfs of spectral types from late L to early T. See
Kirkpatrick (2007) for further discussion.

Most of the scientific inquiry into these ultracool dwarfs has focused on
their formation and youth, on the resultant initial mass function, and on the
determination of their global properties, particularly mass, effective
temperature, metallicity, and cloudiness. This review will focus primarily on
the latter areas. Burrows et al. (2001) provide a much more in depth review
and background to brown dwarf science and is an excellent starting point for
those new to the subject. Kirkpatrick (2007) provides a more current look at
outstanding issues in the field from an observational perspective. Here we
first present a very brief review of the ultracool dwarf spectral types and
the nature of the current datasets. We then move on to discuss the role that
clouds and atmospheric mixing play in controlling the emitted spectra of these
objects and the enigmatic L- to T-type transition. Because clouds control the
spectral energy distribution of the L and early T dwarfs, and since clouds are
inherently difficult to model, constraining gravity solely by comparison of
observations to spectral data is particularly challenging. Finally we will
close with a look forward to some of the ultracool dwarf science opportunities
that will be enabled by JWST. Because of space limitations we neglect several
other important avenues of UCD research, including studies of the IMF, very
young objects, and of objects with unusual colors.

##  2 Spectral Type

The currently known ultracool dwarfs span spectral types from late M, L0
through L9, and T0 through T9. The TiO and VO bands, which dominate the
optical portions of late-M dwarf spectra, disappear in the L dwarfs
(Kirkpatrick et al. 1999), where metallic oxides are replaced by metallic
hydrides and where features due to neutral alkali metals are strong. To
systematize such objects Kirkpatrick et al. established spectral indices and
defined an optical classification scheme for L dwarfs, which is commonly used.
The indices measure the strengths of TiO, VO, CrH, Rb and Cs features as well
as a red color term, at wavelengths between 0.71 and 0.99 $\mu$m.

With the discovery of the T dwarfs, which have very little flux in the
optical, Geballe et al. (2002) defined a near-infrared classification scheme
that encompassed both the L and T dwarfs. The indices measure the strength of
the H2O and CH4 absorption features between 1.1 and $2\,\mu$m, and for the
early L dwarfs a red color term is also used which is slightly modified from
Kirkpatrick et al. (1999).

Burgasser et al. (2002a) introduced a near-infrared classification scheme for
T dwarfs that was very similar to that of Geballe et al. The two schemes were
unified in Burgasser et al. (2006) and this near-infrared scheme is the
commonly used scheme for typing T dwarfs. For L dwarfs, both the optical
Kirkpatrick et al. and the near-infrared Geballe et al. schemes are used;
these usually produce the same type (the Geballe et al. scheme was pinned to
the Kirkpatrick et al. types), but for L dwarfs with unusual colors they can
give significantly different types. Because of this, the classification scheme
used for L dwarfs should always be specified.

Leggett et al. (2007, and other work referenced therein) explore the possible
signatures of the next spectral type, for which the letter Y has been
suggested (Kirkpatrick 2005). It is likely that NH3 features will join the
familiar water and methane absorption seen in T dwarfs as the effective
temperatures approach 600 K. In actuality the situation is likely more
complex: it is already known that the atmospheric NH3 abundance (as seen in
Spitzer mid-infrared spectra of T dwarfs) is reduced by vertical mixing which
drags N2 up from deeper layers in the atmosphere (e.g., Saumon et al. (2006)).
If this mechanism continues to act at low effective temperatures (which is not
a certainty (Hubeny & Burrows 2007)) the near-infrared ammonia features may be
weaker than expected. Another outstanding problem with predicting the
signature of the proposed Y dwarfs, is that the linelist for NH3 is very
incomplete at 1.0-1.5 $\mu$m. Since the near-infrared flux of Y dwarfs is
expectedly to rapidly ‘collapse’ with falling $T_{\rm eff}$ (Burrows et al.
2003), it may even be appropriate to ultimately type these objects with mid-,
instead of near-infrared, spectra. In this case spectra obtained by Spitzer or
JWST would be required for spectral typing. Regardless, dwarfs with $T_{\rm
eff}\sim 650\,\rm K$ are now being found (e.g. Warren et al. 2007), and it is
likely that soon temperatures where significant spectral changes occur will be
reached.

##  3 Ultracool Dwarf Datasets

The L and T dwarfs were discovered primarily as a result of the far-red and
near-infrared Sloan Digital Sky Survey and 2 Micron All Sky Survey (e.g.
Kirkpatrick et al. 1999, Strauss et al. 1999). This continues with current
surveys - the Canada France Hawaii Brown Dwarf Survey and the UKIRT Infrared
Deep Sky Survey (e.g., Lodieu et al. 2007) are identifying extreme-T dwarfs by
their very red far-red and blue near-infrared colors. Spectral classification
is carried out in the far-red or near-infrared. Hence the existing data for L
and T dwarfs is primarily far-red and near-infrared imaging and spectroscopy.
The spectroscopy has been medium- or low-resolution, both because that is all
that is required for the spectral classification, but also because the dwarfs
are faint.

Some ground-based imaging and spectroscopy has been carried out at 3.0-5.0
$\mu$m (e.g. Noll et al. 2000, Golimowski et al. 2004). Such work is extremely
challenging due to the very high and rapidly variable sky background at these
wavelengths, and only the brightest dwarfs could be observed from the ground.

This situation changed with the launch of the Spitzer Space Telescope. Roellig
et al. (2004) and Patten et al. (2006) demonstrated the quality and quantity
of mid-infrared imaging and spectroscopy of L and T dwarfs that Spitzer could
produce, and such work continues through the current, final, cryogenic cycle.
IRS spectral data, which span 6 to 15 $\mu$m, show a strong 11 $\mu$m NH3
absorption feature in T dwarfs, as well as H2O and CH4 absorption features in
both L and T dwarfs (Figure 1). IRAC photometry covers the 3 to 8 $\mu$m
wavelength range, and the 3 to 5 $\mu$m range may continue to be available in
the warm-Spitzer era. These IRAC bandpasses include CH4, CO and H2O features,
and signatures of vertical transport have been recognized in the photometry
(Leggett et al. 2007).

Warren et al. (2007) further suggest that the $H-$ [4.49] color may be a very
good indicator of temperature for dwarfs cooler than 1000 K. For extreme-T
dwarfs the near-infrared CH4 and H2O bands are so strong that it will be
difficult to measure an increase in their strength, hence the mid-infrared may
prove to be vital to interpreting the cold objects.

##  4 Ultracool Dwarf Atmospheres

Ultracool dwarf emergent spectra are controlled by the variation in abundances
of important atomic, molecular, and grain absorbers both with height in the
atmosphere at a given age and over time as the objects cool. The major atomic
and molecular absorption features are imprinted on spectra that have no true
continuum222Sharp & Burrows (2007) and Freedman et al. (2008) discuss the
atmospheric opacity sources in detail.. Flux emerges from a many-scale-height-
thick range of depths in the atmosphere as a function of wavelength. For
example (Fig. 4), in an early L dwarf, brightness temperatures333Brightness
temperature (the temperature that a blackbody that emits radiation of the
observed intensity at a given wavelength) is commonly used in planetary
atmospheres studies to elucidate the temperature of the emitting level in an
atmosphere.  range from 1000 K in the depths of alkali absorption lines
(Burrows et al. 2000) in the far red, to well over 2000 K in the molecular
windows in between strong water absorption bands. By providing a continuum
opacity source, clouds can limit the flux emerging in some molecular window
regions, but not others. Furthermore the strength of some molecular absorption
features, particularly $\rm CH_{4}$, CO, and $\rm NH_{3}$, can depend on the
strength of mixing in the atmosphere. Thus a full description of an ultracool
dwarf atmosphere hinges on the dwarf’s gravity, effective temperature, cloud
properties, and mixing. In this section we summarize the important unsolved
problems related to these atmospheres.

Figure 4: Brightness temperature as a function of wavelength for atmosphere
models which include (solid) or exclude (dotted) silicate and iron clouds
(Ackerman & Marley 2001). Brightness temperature increases downward to suggest
increasing depth in the atmosphere from which the wavelength-dependent flux
emerges. The solid straight line indicates the base of the silicate cloud
while the long dashed line denotes the ‘top’ of the cloud (the level in the
atmosphere at which the cloud column extinction reaches 0.1). Shading suggests
the decrease in cloud extinction with altitude. Since cloud particle radii
exceed $10\,\rm\mu m$ in these models, the Mie extinction efficiency is not a
strong function of wavelength over the range shown. Shown are models
characteristic of (a) an early-type L dwarf with $T_{\rm eff}=1800\,\rm K$,
(b) a late-type L dwarf with $T_{\rm eff}=1400\,\rm K$, and (c) a T dwarf with
$T_{\rm eff}=900\,\rm K$ . All of these models are for solar composition and
gravity appropriate for a 30 Jupiter-mass brown dwarf. Note that the spectral
region just longward of $1\,\rm\mu m$ is particularly sensitive to the cloud
opacity.

###  4.1 Clouds

At high effective temperatures the column abundance of condensates (see Marley
2000 for a discussion of influences on cloud optical depth) is low and the
difference between models computed with and without cloud opacity is slight
(Figure 4). At lower temperatures, however, the cloud substantially alters the
temperature profile of the atmosphere and provides a continuum opacity source
that limits the depth to which the usual molecular windows probe into the
atmosphere. By the effective temperature of the mid-T dwarfs, however, most of
the flux emerges from above the cloud level and the clouds are again less
important. For the effective temperature range of the mid to late L dwarfs and
the early T dwarfs, however, clouds clearly play a very large role in
controlling the vertical structure and emergent spectra of brown dwarfs.

Thus any attempt to model brown dwarf atmospheres must include a treatment of
clouds. However clouds are the leading source of uncertainty in terrestrial
atmosphere models and are inherently difficult to model. Their influence
depends on the variation of particle size, abundance, and composition with
altitude, which in turn depend on the complex interaction of many
microphysical processes (e.g., Ackerman & Marley 2001, Helling et al. 2006).
Fits of model spectra to observational data are highly sensitive to the
treatment of clouds in the underlying atmosphere model and the approaches
taken by various modeling groups vary widely (see the comparison study in
Helling et al. 2008a). For example Cushing et al. (2008) demonstrate
reasonably accurate fits of model spectra to near- and mid-IR spectra of a
sample of L and T dwarfs, but the precise values of effective temperature and
gravity obtained from the fits depend entirely upon the cloud sedimentation
efficiency (Ackerman & Marley 2001) assumed. Since the models are highly
dependent on the cloud description, the derived effective temperature and
gravities, while plausible, are nevertheless uncertain. A similar conclusion
was reached by Helling et al. (2008b). Finding a selection of L dwarfs with
known $g$ and $T_{\rm eff}$ that could serve as calibrators of the model
spectra would be invaluable. L dwarf companions to main sequence stars with
constrained ages, L dwarf binaries with resolved orbits, and L dwarfs in
clusters of known ages are all promising targets for such work.

###  4.2 Characterizing Clouds

The spectral range of the InfraRed Spectrometer (IRS) on Spitzer includes the
$10\,\rm\mu m$ silicate feature which arises from the Si-O stretching
vibration in silicate grains. The spectral shape and importance of the
silicate feature depends on the particle size and composition of the silicate
grains. According to phase-equilibrium arguments, in brown dwarf atmospheres
the first expected silicate condensate is forsterite $\rm Mg_{2}SiO_{4}$
(Lodders 2002), at $T\sim 1700\,\rm K$ ($P=1\,\rm bar$). Since Mg and Si have
approximately equal abundances in a solar composition atmosphere, the
condensation of forsterite leaves substantial silicon, present as SiO, in the
gas phase. In equilibrium, at temperatures about 50 to 100 K cooler than the
forsterite condensation temperature, the gaseous SiO reacts with the
forsterite to form enstatite, $\rm MgSiO_{3}$ (Lodders 2002). The precise
vertical distribution of silicate species depends upon the interplay of the
atmospheric dynamics and chemistry and such details have yet to be fully
modeled, although efforts to improve the detailed cloud modeling continue
(e.g., Cooper et al. 2003; Woitke & Helling 2004; Helling & Woitke 2006;
Helling et al. 2008b).

Figure 5:  Top: Spitzer IRS spectrum of 2MASS 2224 (L4.5) and the best fitting
model from Cushing et al. (2006). Middle: Optical absorption ($Q_{\rm abs}/a$)
for amorphous enstatite ($\rm MgSiO_{3}$) and forsterite ($\rm Mg_{2}SiO_{3}$)
for three different particle sizes, 0.1, 1, and $10\,\rm\mu m$. Bottom:
Optical absorption for crystalline enstatite, also for three different
particle sizes. The deviation of the model (shifted vertically) from the data
suggests that additional small, and perhaps crystalline, silicate grains are
required to adequately account for the observed spectrum.

In brown dwarf clouds there is likely a range of particle sizes, ranging from
very small, recently condensed grains, to larger grains that have grown by
agglomeration. The mean particle size for silicate grains in L dwarf model
atmospheres is typically computed to be in the range of several to several
tens of microns (Ackerman & Marley 2001, Helling et al. 2008b) . Figure 5
compares the absorption efficiency of silicate grains of various sizes,
composition, and crystal structures to the spectrum of 2MASS J2224-0158
(L4.5). For each species the quantity $Q_{\rm abs}/a$, or Mie absorption
efficiency divided by particle radius, is shown; all else being equal, the
total cloud optical depth is proportional to this quantity (Marley 2000).
Large grain sizes tend to have a relatively flat absorption spectra (dashed
lines) across the IRS spectral range. Only grains smaller than about
$3\,\rm\mu m$ in radius show the classic $10\,\rm\mu m$ silicate feature
(Hanner et al. 1994). Figure 5 suggests that the mismatch between the models
and data may arise from a population of silicate grains that is not captured
by the cloud model used to construct the figure (Ackerman & Marley 2001). The
actual silicate cloud may contain both more small particles and a mixture of
enstatite and forsterite grains (e.g., Helling & Woitke 2006), although
detailed models for this particular dataset have not been attempted.
Furthermore the model shown in the figure employs optical properties of
amorphous silicate. It is possible, especially at the higher pressures found
in brown dwarf atmospheres, that the grains are crystalline, not amorphous.
Indeed laboratory solar-composition condensation experiments produce
crystalline, not amorphous, silicates (Toppani et al. 2004). Crystalline
grains (Figure 5) can have larger and spectrally richer absorption cross
sections.

###  4.3 The Transition from L to T

The evolutionary cooling behavior of a given substellar object can be inferred
from the field brown dwarf near-infrared color-magnitude diagram (Fig. 6a).
Over tens to hundreds of millions of years a given substellar object first
moves to redder $J-K$ colors as it cools while falling to fainter J
magnitudes. Around ${\rm M}_{J}\sim 14-15$ the $J-K$ color turns bluer and the
$J$ magnitude slightly (and counter-intuitively) brightens (Dahn et al. 2002,
Tinney et al. 2003, Vrba et al. 2004). With further cooling a given dwarf
finally falls to fainter $J$ magnitudes and apparently continues to slightly
turn somewhat bluer in $J-K$. The behavior with even greater cooling is as yet
uncertain until many more objects with $T_{\rm eff}<700\,\rm K$ are found.

The ‘L to T’ transition is the ‘horizontal branch’ of the color-magnitude
diagram as objects move from red to blue in the diagram. From Figure 3 we know
that this color (and underlying spectral) change from late-type L dwarfs with
$J-K\sim 2.5$ to blue T dwarfs with $J-K\sim-1$ happens rapidly over a small
range of effective temperature. No brown dwarf evolution model can currently
reproduce the magnitude of the observed color change over such a small range
of $T_{\rm eff}$. Various explanations have been suggested including holes
forming in the condensate cloud decks (Ackerman & Marley 2001, Burgasser et
al. 2002b), an increase in the efficiency of grain sedimentation (Knapp et al.
2004), or a change in particle size (Burrows et al. 2006). In a series of
papers Tsuji (Tsuji 2002, Tsuji & Nakajima 2003, Tsuji et al. 2004) proposed
that a physically very thin cloud could self-consistently explain the rapid L
to T transition. These models indeed exhibit a somewhat faster L- to T-like
transition, but are still not consistent with the observed rapidity of the
color change. More recently Tsuji (2005) has favored a sudden collapse of the
global cloud deck at the transition along the lines of the Knapp et al. (2004)
suggestion.

Support for a rapid increase in sedimentation efficiency at the L to T
transition has come from the model analysis of the 0.8–$14.5\,\mu$m spectra of
four transition dwarfs by Cushing et al. (2008), who find that the cloud
sedimentation efficiency (Ackerman & Marley 2001) indeed increases across the
transition. The Cushing et al. (2008) sample includes two pairs of mid- to
late-L dwarfs with very different near-IR colors. The authors find that the
redder L dwarfs have less efficient sedimentation and therefore thicker cloud
decks consisting of smaller particles, although gravity also may play a role
for one pair. For one of the red L dwarfs Cushing et al. (2006) identified a
broad absorption feature at 9–$11\,\mu$m which may be due to the presence of
small silicate grains (Figure 5).

The difficulty in characterizing the L to T transition arises from our lack of
understanding of the masses and effective temperatures of objects at various
locations in the ultra-cool dwarf color-magnitude diagram. It is not clear,
for example, if the reddest field L dwarfs are more or less massive than bluer
objects or if the turn towards the blue in $J-K$ is mass dependent (although
there are some indications that it may be (Metchev & Hillenbrand 2006)). There
are two ways in which this shortcoming in current understanding could be
addressed. First, observing the orbits of binary brown dwarfs allows the total
system mass to be measured. Secondly, photometry leading to near-infrared
color-magnitude diagrams for many clusters with a variety of ages and
metalicities will constrain the nature of transition.

To date, the cluster color magnitude diagram (CMD) has only reached the
transition in the Pleiades and perhaps the Sigma Orionis clusters. Figure 6b
shows the currently best available CMD for the Pleiades. Two objects in this
figure can be seen to have turned towards the blue. Deeper searches to fainter
magnitudes in this cluster should soon reveal the expected downward turn to
the fainter J magnitudes apparent in the field CMD. By constructing evolution
models at the age of the Pleiades, it should be possible to constrain the mass
at the turnoff from the red L sequence. Given enough clusters of different
ages, the turnoff effective temperature and gravity can be constrained, thus
illuminating the dependence of the turnoff on gravity and perhaps metallicity.

Figure 6: Near-infrared color-magnitude diagrams for field and cluster
ultracool dwarfs. (a) Black dots show single field L & T dwarfs. Green dots
are resolved components of binary systems. Dotted circles are suspected (but
unresolved) binaries (figure courtesy M. Liu based on Liu et al. (2006)). (b)
Candidate ultracool dwarfs in the Pleiades in the most sensitive current
survey (Casewell et al. 2007). Faintest objects in this plot have masses of
about $11\,\rm M_{Jup}$. Note that at a fixed magnitude the cluster members
tend to be redder than the field objects, which is likely a signature of low
gravity. JWST will obtain spectra of quality comparable to Figure 1 for the
candidate objects shown on this panel which will help calibrate evolutionary
models of the brown dwarf cooling sequence. The detection limit for NIRCam on
JWST is at about $J=22$ for this cluster or $\sim 1\,\rm M_{Jup}$ . Model
predictions for colors of objects with $J>15$ are shown in Figure 8.

JWST will be able to obtain moderate resolution ($R\sim 1000$) spectra on the
current Pleiades candidates in Figure 6 in about 2.5 hours at $S/N\sim 20$.
This spectral resolution should be sufficient to identify, for example, FeH
and $\rm CH_{4}$ bands as they vary through the spectral sequence. This
combination of evolution models and spectra should tightly constrain the
empirical cooling sequence. Although this cluster is likely too large on the
sky for efficient searching by JWST, NIRCAM could in principle find objects
with masses as low as about $1\,\rm M_{J}$. Surveys of more compact clusters
would not reach to such low masses, but should nevertheless be deep enough to
find many young T dwarfs that have already undergone the transition.

###  4.4 The Latest T Dwarfs

At this time only 16 very cool ($T_{\rm eff}<900\,\rm K$) dwarfs with types T7
and later are known, and of these only four are T8 or later. New surveys that
go fainter than 2MASS and SDSS have started or are planned, and several groups
are attempting to push to later and cooler types (e.g. Warren et al. 2007).
All but two of the very late T dwarfs are isolated (the exceptions are Gl 570
D, Burgasser et al. 2000, and HD 3651B, Mugrauer et al. 2007). Since age is
unknown for field dwarfs and brown dwarf cool with time, observed spectra must
be compared with models or spectra of fiducial objects, to constrain mass and
age. Since there are only a few T dwarfs with highly constrained properties,
accurate model analysis is crucial for secure determination of gravity and
hence mass at the bottom of the T sequence. For field brown dwarfs with ages
in the range $\sim$ 1-5 Gyr and masses of 20–50 Jupiter-masses, effective
temperature will lie in the range of $\sim$600–800 K. To understand the
physical parameters of these elusive, cold, and low-mass dwarfs requires
observation of their full spectral energy distribution.

###  4.5 Vertical Mixing and Chemical Disequilibrium

It has long been understood that the abundances of molecules in Jupiter’s
atmosphere depart from the values predicted purely from equilibrium chemistry
(Prinn & Barshay 1977; Barshay & Lewis 1978; Fegley & Prinn 1985; Noll et al.
1988; Fegley & Lodders 1994; Fegley & Lodders 1996)444In stellar atmospheres,
departures from thermochemical equilibrium can arise from interactions of
atoms and molecules with the non-thermal radiation field (Hauschildt et al.
1997; Schweitzer, Hauschildt & Baron 2000). In brown dwarf atmospheres this
effect is negligible.. Rapid upwelling can carry compounds from the deep
atmosphere up into the observable regions of the atmosphere on time scales of
hours to days. When this convective timescale is shorter than the timescale
for chemical reactions to reach equilibrium, then the atmospheric abundances
will differ from those that would be found under pure equilibrium conditions.
The canonical example of this situation is carbon monoxide in Jupiter’s
atmosphere. In Jupiter’s cold and dense upper troposphere, carbon should
almost entirely be found in the form of methane. Deeper into the atmosphere,
where temperatures and pressures are higher, CO should be the principal
carrier of carbon. Rapid vertical mixing, combined with the strong C-O
molecular bond, means that CO molecules can be transported to the observed
atmosphere faster than chemical reactions can reduce the CO into $\rm CH_{4}$.
The observed enhancement of CO, combined with (uncertain) reaction rates
places limits on the vigor of convective mixing in the atmosphere.

##### CO

Analyses of the 4.5 – $5\,\rm\mu m$ spectra of the T dwarfs Gl 229B and Gl 570
D reveal an abundance of CO that is over 3 orders of magnitude larger than
expected from chemical equilibrium calculations (Noll et al. 1997, Oppenheimer
et al. 1998, Griffith & Yelle 1999, Saumon et al. 2000), as anticipated by
Fegley & Lodders (1996). Photometry in the $M$ band, which overlaps the CO
band at $\rm 4.6\,\rm\mu m$, shows that an excess of CO may be a common
feature of T dwarfs. Golimowski et al. (2004) have found that the $M$ band
flux is lower than equilibrium models predict based on the $K$ and
$L^{\prime}$ fluxes in all of the T dwarfs in their sample. The low $M$ band
flux certainly arises from an excess of CO above that expected by equilibrium
models. Since CO is a strong absorber in the $M$ band, a brown dwarf can be
much fainter in this band than would be predicted by equilibrium chemistry.
Equilibrium models predict that brown dwarfs and cool extrasolar giant planets
should be bright at $M$ band, hence any flux decrement at this wavelength
would have implications for surveys for cool dwarfs and giant planets. The
degree to which this is a concern depends upon how the vigor of mixing
declines with following effective temperature. Hubeny & Burrows (2007)
recently have argued that this will not be a concern at temperatures below
about 500 K, however, because the vigor of mixing falls with effective
temperature.

##### NH3

Ammonia forms from $\rm N_{2}$ by the reaction $\rm
N_{2}+3\,H_{2}\Leftrightarrow 2\,NH_{3}$. $\rm N_{2}+3H_{2}$ is favored at low
pressures and high temperatures because of higher entropy. $\rm NH_{3}$ is
favored at low temperatures, but since molecular nitrogen is a strongly bound
molecule, reactions involving this molecule typically have high reaction
energies and proceed very slowly at low temperatures. Like CO, $\rm N_{2}$ is
favored at the higher temperatures found deep in brown dwarfs atmospheres.
Again, like CO, vigorous vertical transport can bring $\rm N_{2}$ in the upper
atmosphere faster than it can be converted to $\rm NH_{3}$, resulting in an
excess of $\rm N_{2}$ compared to the values expected from chemical
equilibrium. Figure 7 illustrates the effect of mixing (Saumon et al. 2007) on
the $10\,\rm\mu m$ ammonia band in the spectra of the T8 dwarf 2MASS0415-0935
(Burgasser et al. 2002a).

Figure 7: Fits of the IRS spectrum of 2MASS J0415-0935 (Saumon et al. 2007)
showing the difference between a model in chemical equilibrium and a model
that includes vertical transport that drives the nitrogen and carbon chemistry
out of equilibrium. The red thin curve is the best-fitting model in chemical
equilibrium, and the blue thin curve is the best-fitting nonequilibrium model.
The data and the noise spectrum are shown by the histograms (black). The
uncertainty on the flux calibration of the IRS spectrum is $\pm 5\%$. The
model fluxes, which have not been normalized to the data, are shown at the
resolving power of the IRS spectrum.

JWST will obtain higher resolution mid-infrared spectra than the Spitzer data
analyzed by Saumon et al. (2007). Higher spectral resolution on more targets
will allow more in depth studies of atmospheric mixing. Since the vertical
profile of mixing also influences cloud particle sizes and optical depths (in
L- and L to T transition dwarfs), mapping out the eddy diffusion coefficient
(which parameterizes mixing) as a function of mass and effective temperature
will help to shed light on cloud dynamics as well as atmospheric chemistry.

##  5 Opportunities for JWST

As we have highlighted in the above sections, there are many unsolved problems
in the study of ultracool dwarfs. In this section we will briefly summarize
some of the most promising avenues for JWST.

###  5.1 Characterizing Rosetta Stone Dwarfs

The characterization of most field brown dwarfs, particularly the L and early
T dwarfs, is hampered by the dependency of model fits on the particulars of
the cloud models used to generate model atmospheres and spectra for comparison
to data. Brown dwarfs of known mass, metallicity, and age are thus of
particular importance as calibrators for the entire brown dwarf cooling
sequence. This section discusses some opportunities for unearthing and
deciphering ‘Rosetta Stone’ ultracool dwarfs with known or easily deduced
masses and effective temperatures. Such objects could turn the page to much
greater understanding of our library of known L and T dwarfs.

#### Color-magnitude diagram for clusters to low masses

By providing cluster color-magnitude diagrams to very low masses (a Jupiter-
mass or less) JWST will revolutionize our understanding of brown dwarf cooling
in environments controlled for age and metallicity. Comparison of spectra of
objects with known properties to models will finally provide insight into the
variation in cloud properties with mass and effective temperature. In the
Pleiades, Casewell et al. (2007) have detected objects with masses as low as
11 Jupiter masses (Figure 6). In each cluster a brown dwarf of a given mass
will be found either earlier or later on its cooling track, depending on the
cluster age. Since the evolutionary cooling of brown dwarfs is well
understood, spectra of cluster objects with known masses will definitively
connect spectral features with gravity for mid- to late L dwarfs. In the
Pleiades such a project should be straightforward for JWST as NIRSPEC will be
able to obtain $R\sim 1000$ JHK spectra of the known low-mass cluster members
in as little as a few hours. A Jupiter-mass object will be about 6 magnitudes
fainter in $J$ band. Assuming JWST could survey a sufficiently large area to
find candidates, it would define the brown dwarf cooling curve to a degree
still not reached in the disk population. Model photometric predictions
(Burrows et al. 2003) for such objects are shown in Figure 8. The actual
trajectory in color-magnitude space will ultimately depend upon the interplay
of water clouds and atmospheric mixing with the emitted spectra. Comparison of
models such as those in the figure with data will test our understanding
through this as-yet unexplored range of parameter space.

Figure 8: Predicted absolute J magnitude ($M_{J}$) vs. $J-K$ color for a range
of brown dwarf masses and ages. The numbers by the symbols denote the masses
of the objects in Jupiter mass units. In the Pleiades the JWST NIRCAM
detection limit will be about 1 Jupiter mass. Figure and description from
Burrows et al. (2003); see discussion therein for greater detail.

#### Resolved Spectra of Close Binaries

Binary stars have long served as astronomical workhorses, helping to reveal
important details of stellar astrophysics. Likewise binary brown dwarfs are
also exceptionally useful. The orbits of close L- and T-dwarf binaries allow
the total system mass to be determined as Bouy et al. (2004) have done for the
binary 2MASSW J0746425+2000321. Assuming co-evality and equal metallicity
combined with the total system mass and resolved spectra of the individual
dwarfs furthermore allows such systems to elucidate the interpretation of
brown dwarf spectra. Many more tight binaries have been found by the
combination of HST and ground based adaptive optics imaging (see summary in
Bouy et al. 2008). Orbital periods for many of these systems appear to be less
than twenty years, so dynamical masses will be available during the JWST
mission lifetime. The combination of the ground-based astrometry and
photometry and resolved NIRSPEC high $R$ spectra of many of the individual
objects, particularly for the tightest binaries, in these systems should
provide important constraints on models of brown dwarf evolution, atmospheric
structure, and emergent spectra.

#### Cloud Behavior from L to T to Y

Perhaps the greatest single observational result that could drive improvements
in understanding of the L to T transition would be a direct measurement of
surface gravity (or almost equivalently, mass) and effective temperature of
late L and early T dwarfs. This would elucidate the dependence of the
initiation of the L to T transition on mass and $T_{\rm eff}$. Constraining
the turnoff absolute magnitude in a variety of clusters of known ages and
resolving the spectra of close binaries that have measurable dynamical masses
would highly constrain the nature of the L to T transition.

Likewise as brown dwarfs further cool through the T sequence a number of open
issues remain. Although T dwarfs are generally modeled as being entirely cloud
free, some models that include very thin cloud decks better reproduce the
spectra and color of the T’s. As brown dwarfs cool through about 500 K, thin
water clouds should appear high in their atmospheres. With falling effective
temperature these clouds are expected to thicken and begin to substantially
alter the spectra of the Y dwarfs. It is entirely possible that, like the
departure of clouds in the late L dwarfs, the arrival of clouds in the early Y
dwarfs will produce unexpected and perhaps rapid color and spectral changes.
Although we cannot yet identify what these changes might be, the same type of
observations noted above will also be invaluable in constraining the Y dwarfs.
Since the optical and near-IR flux of the Y dwarfs is expected to rapidly
decrease (Burrows et al. 2003), the signatures of the water clouds will be
best obtained in the mid-IR by JWST.

###  5.2 Spectra of Very Cool Objects

As of early 2008, the brown dwarf with the lowest estimated effective
temperature is ULAS J0034-00 with $T_{\rm eff}\sim 650\,\rm K$ (Warren et al.
2007). A number of ongoing and future searches will certainly find cool
objects in the solar neighborhood (e.g., UKIDSS, Pan-STARRS, and the WISE
mission). In particular WISE will have sufficient sensitivity to detect a
$T_{\rm eff}\sim 200\,\rm K$ brown dwarf ($2\,\rm M_{J}$ at 1 Gyr) at a
distance of about 2 pc. Assuming such nearby targets are found, JWST will
produce exquisite spectra that will be unobtainable from the ground. A survey
with NIRSPEC in high resolution mode with the YJH and LM gratings would nicely
sample the spectra of cool field dwarfs. For a 500K dwarf at 10 pc, we
estimate an exposure of about a minute will provide a spectrum with S/N of
about 100 in $M$ band. A one hour exposure would be required for the same
dwarf at 25 pc. Detection limits and model spectra are shown in Figure 9.

Figure 9: Spectra (flux in millijanskys) vs. wavelength (in microns) for a
range of brown dwarf masses at an age of 1 Gyr and a distance of 10 pc.
Superposed are the approximate point-source sensitivities for instruments on
Spitzer (red) and JWST (blue). The JWST/NIRCam sensitivities are $5\,\sigma$
and assume an exposure time of $5\times 10^{4}\,\rm sec$. The JWST/MIRI
sensitivity curve from 5.0 to $27\,\rm\mu m$ is $10\,\sigma$ and assumes an
exposure time of $10^{4}\,\rm sec$. Figure and description from Burrows et al.
(2003); see discussion therein for greater detail.

High quality spectra of cool dwarfs will be important for a number of reasons.
First, cold disk objects possess effective temperatures comparable to those of
middle-aged to old extrasolar giant planets. The disk population of brown
dwarfs will thus provide ground truth for the spectral features that such cold
objects exhibit (for example $\rm NH_{3}$ should appear in the near-IR (Saumon
et al. 2001, Burrows et al. 2003, Leggett et al. 2008)). These objects will
also have water clouds. The experience with the challenge of modeling silicate
and iron clouds in L dwarfs alluded to above implies that water clouds will be
no more tractable. The cold disk population will thus provide a proving ground
for exoplanet water cloud modeling, which is undoubtedly needed (see Marley et
al. 2007 and references therein).

Atmospheric mixing, long recognized in Jupiter’s atmosphere, is also important
for brown dwarfs (see §4.5), yet the most diagnostic spectral region for this
process (the CO band at $4.6\,\rm\mu m$) lies in a blind spot for Spitzer
spectroscopy and most groundbased observatories, but not JWST. Thus the nearby
disk brown dwarfs will elucidate the extent to which $M$ band flux of objects
in this effective temperature range is impacted by excess atmospheric CO.
Vertical mixing could be an important consideration for the direct detection
of giant planets around nearby stars (Golimowski et al. 2004; Marley et al.
2006; Hinz et al. 2006). After the discovery of Gl 229B, Marley et al. (1996)
suggested that a substantial 4 to $5\,\rm\mu m$ flux peak should be a
universal feature of giant planets and brown dwarfs. This expectation,
combined with a favorable planet/star flux ratio, has made the band a favorite
for planet detection (Burrows et al. 2005). However, groundbased and IRAC
photometry suggests that cool dwarfs are fainter in this region—and the $L$
band region is brighter—than predicted by equilibrium chemistry. Given these
and other considerations, Leggett et al. (2007) suggested that the comparative
advantage of ground-based searches for young, bright giant planets at $M$ band
might, such as the searches planned with the JWST coronagraph, might be
somewhat less than currently expected (see also Marley et al. 2007).

Hubeny et al. (2007), however, recently predicted that the vigor of
atmospheric mixing will decline with effective temperature. If this is indeed
the case then $M$ band will remain a fruitful hunting ground for extrasolar
giant planet coronagraphic imaging. NIRCAM photometry of cold brown dwarfs
will certainly illuminate this issue.

###  5.3 Mid-IR Spectra Beyond Spitzer

Examples of some of the best available Spitzer IRS mid-infrared spectra of L
and T dwarfs are shown in Figures 5 and 7. The spectral region between about 6
and $15\,\rm\mu m$ is important for a number of reasons. First there are
several strong molecular bands in this region, including water, methane, and
ammonia. As recounted above methane and ammonia are particularly sensitive to
atmospheric mixing. The Si-O vibrational band, seen in the opacity of small
silicate grains (Figure 5), also may trace the arrival of silicate clouds.
Finally a number of other molecules, not yet detected in brown dwarf spectra,
have absorption features in this range (Mainzer et al. 2007).

MIRI will produce much higher resolution and S/N spectra than the best data
from IRS (see the sensitivity curve in Figure 9). Higher quality spectra will
allow for more robust detection of silicate features, perhaps including the
sort of fine structure in the grain opacity seen in the lower panel of Figure
5, as well as for fine detail on the molecular features (see model prediction
in Burrows et al. (2003)). If silicates are indeed detected, high resolution
spectra could in principle differentiate between the particular silicate
species, including forsterite, enstatite, and even quartz ($\rm SiO_{2}$,
Helling et al. 2006), and whether the grains are in crystalline or amorphous
form (Figure 5).

##  6 Conclusions

The next decade holds the potential to substantially improve our fundamental
understanding of ultracool dwarfs. Ground and space based surveys for very
cool dwarfs, including UKIDSS, Pan-STARRS, and the WISE mission as well as
deep surveys of young clusters will provide a host of ultracool dwarf targets
for JWST. For cold nearby dwarfs JWST will provide unparalleled near- and
especially mid-infrared spectra. These observations will constrain the water
clouds expected to be present in objects with $T_{\rm eff}<500\,\rm K$ and
measure the degree of atmospheric mixing. Both types of observations are
highly relevant to the ultimate direct detection and characterization of
extrasolar giant planets by coronagraphy.

In young clusters (e.g, the Pleiades and younger) JWST will provide
exceptional quality spectra of many known cluster members and will have the
capability of imaging objects down to about one Jupiter mass and below. Such
observations will constrain the evolutionary cooling tracks for dwarfs with
lower gravities than most field objects and will tightly constrain the nature
of the L to T transition by revealing its dependence on gravity.

Many other opportunities, including producing resolved spectra of tight binary
dwarfs and searching for spectral signatures of condensates and low abundance
gasses are also possible. Combined with the inevitable unexpected discoveries,
there is no doubt that JWST will bring brown dwarf astrophysics into the same
highly constrained realm as stellar astrophysics. Interpretting these expected
datasets will undoubtedly require substantial improvements to atmosphere and
evolution modeling, particularly cloud and chemical transport modeling of
ultracool dwarf atmospheres.

###### Acknowledgements.

The authors thank M. Cushing, M. Liu, & D. Saumon, for helpful conversations
on the future of brown dwarf science and A. Burrows, Ch. Helling, X. Tielens
and K. Zahnle for thoughtful comments on the manuscript. We thank M. Cushing,
M. Liu, and K. Lodders for preparing Figures 1, 6a & 2, respectively and C.
Nixon for kindly providing the Cassini CIRS spectrum of Jupiter for Figure 1.

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