To appear in the Astronomical Journal, July 2007
Periodic accretion from a circumbinary disk in the young binary
UZ Tau E
Eric L. N. Jensen1, Saurav Dhital1,2, Keivan G. Stassun2, Jenny Patience3, William
Herbst4, Frederick M. Walter5, Michal Simon5, Gibor Basri6
ABSTRACT
Close pre–main-sequence binary stars are expected to clear central holes in
their protoplanetary disks, but the extent to which material can flow from the
circumbinary disk across the gap onto the individual circumstellar disks has been
unclear. In binaries with eccentric orbits, periodic perturbation of the outer
disk is predicted to induce mass flow across the gap, resulting in accretion that
varies with the binary period. This accretion may manifest itself observationally
as periodic changes in luminosity. Here we present a search for such periodic
accretion in the pre–main-sequence spectroscopic binary UZ Tau E. We present
BV RI photometry spanning three years; we find that the brightness of UZ Tau E
is clearly periodic, with a best-fit period of 19.16± 0.04 days. This is consistent
with the spectroscopic binary period of 19.13 days, refined here from analysis
of new and existing radial velocity data. The brightness of UZ Tau E shows
significant random variability, but the overall periodic pattern is a broad peak
in enhanced brightness, spanning more than half the binary orbital period. The
variability of the Hα line is not as clearly periodic, but given the sparseness of
the data, some periodic component is not ruled out. The photometric variations
are in good agreement with predictions from simulations of binaries with orbital
1Swarthmore College Department of Physics and Astronomy, 500 College Ave., Swarthmore PA 19081;
ejensen1@swarthmore.edu.
2Department of Physics and Astronomy, Vanderbilt University, Nashville, TN 37235
3Division of Physics, Mathematics, and Astronomy, California Institute of Technology, Pasadena, CA
91125
4Astronomy Department, Wesleyan University, Middletown, CT 06459
5Department of Physics and Astronomy, State University of New York, Stony Brook, NY 11794-3800
6Astronomy Department, University of California, 521 Campbell Hall, Berkeley, CA 94720
http://arxiv.org/abs/0704.0307v2
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parameters similar to those of UZ Tau E, suggesting that periodic accretion does
occur from circumbinary disks, replenishing the inner circumstellar disks and
possibly extending the timescale over which they might form planets.
Subject headings: accretion disks — circumstellar matter — stars: binaries: spec-
troscopic — stars: individual (UZ Tau E) — stars: pre-main sequence — plane-
tary systems: protoplanetary disks
1. Introduction
It is now well-established that most stars are members of binary systems at birth, and
that many of these stars are surrounded by disks similar to those found around young single
stars (see, e.g., the recent review by Monin et al. 2006). Thus, understanding the origin of
binaries is vital to understanding the star formation process. The predominance of binaries
also means that, based on number of systems alone, most potential sites of planet formation
lie in multiple systems. However, interactions between stars and disks in binary systems
can alter disk structure (Beckwith et al. 1990; Jensen et al. 1994; Osterloh & Beckwith 1995;
Jensen et al. 1996a,b; Jensen & Mathieu 1997), resulting in a more complicated environment
for planet formation. Nonetheless, the discovery of planets in relatively close binary systems
(Eggenberger et al. 2004) shows that binary systems are viable sites of planet formation.
Understanding the extent to which binaries modify the structure of their surrounding disks
is important for understanding the possible diversity of planetary system environments.
In addition, though the mass ratios are different, the interactions between stellar binary
companions and disks involve the same physics as those between planetary companions and
disks (e.g., D’Angelo et al. 2006) but are more easily observable. Thus, an understanding
of binary-disk interactions may help us understand planet formation around single stars as
well.
A binary star system may have up to three disks: two circumstellar disks, one each
around the primary and secondary; and a circumbinary disk outside the binary orbit. Both
analytic calculations and numerical simulations show that the region between these disks
is not stable for orbiting disk material. However, the question of how easily material can
flow from the outer, circumbinary disk across the gap to the circumstellar disks has not
been clear, either from an observational or a theoretical standpoint. The spectral energy
distributions of some young binaries show the clear signature of a cleared central region, while
other, apparently similar systems do not (Jensen & Mathieu 1997). Theoretical analyses by
Lin & Papaloizou (1993) and Artymowicz & Lubow (1994) suggested that material in the
region around the binary orbit is cleared, creating a quasi-equilibrium structure with three
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distinct disks and a cleared region between them. If the gap between disks is impermeable,
the disks evolve independently of each other. Since the presence of a binary companion may
increase the rate of accretion (Clarke 1992; Ostriker et al. 1992) and the binary orbit presents
a constraint on the size of each circumstellar disk, the circumstellar disks would be exhausted
much more quickly than disks around single stars. However, smoothed particle hydrodynamic
simulations by Artymowicz & Lubow (1996) (hereafter AL96; see also Günther & Kley 2002)
predicted that material may indeed flow from the circumbinary disk to the circumstellar
environment, with the accretion rate varying with the phase of the binary orbit.
If such periodic accretion occurs in young binaries, it may be detectable observationally
by a periodic brightening of the system as the material flowing from the circumbinary disk
shocks when it collides with the circumstellar disk(s) or accretes onto the stellar surface(s).
Observations of the T Tauri spectroscopic binary DQ Tau by Mathieu et al. (1997) showed
such brightening, occurring at the binary orbital period. In addition, DQ Tau shows periodic
variations in spectral veiling and emission line intensities with orbital phase (Basri et al.
1997), providing strong support for the broad picture of mass flow across gaps suggested by
AL96. However, subsequent searches in other young, short-period binary systems surrounded
by disks have yielded mixed results. Alencar et al. (2003) did not find periodic photometric
variations in AK Sco, but they did find that the blue wing of the Hα line, and both the
blue and red wings of the Hβ line, vary with the binary orbital period, as does the total
Hα equivalent width. V4046 Sgr has shown periodic photometric variations at the binary
orbital period (Quast et al. 2000; Mekkaden 2000), though an earlier study did not find
such variations (Byrne 1986). Like AK Sco, however, V4046 Sgr does show variations in the
equivalent width and shape of Balmer lines as a function of orbital phase (Stempels & Gahm
2004). No dedicated photometric monitoring of UZ Tau E has been reported in the literature
to date, but in previous spectroscopic observations, neither Hα equivalent width nor spectral
veiling has shown any obvious dependence on binary orbital phase (Mart́ın et al. 2005).
Motivated by previous observational work and a desire to understand accretion in binary
systems, we have undertaken a photometric monitoring campaign for the pre–main-sequence
(PMS) spectroscopic binary UZ Tau E. UZ Tau, in the Taurus-Auriga star-forming region,
was first discovered to be variable by Bohlin during a bright outburst in 1921 (Bailey 1921;
Bohlin 1923) and identified as a ∼ 3.′′7 binary by Joy & van Biesbroeck (1944), one of the first
pre–main-sequence binaries to be identified. Subsequently, both components of the binary
were found to be binary themselves, making this a quadruple system. Simon et al. (1992)
and Ghez et al. (1993) identified UZ Tau W as a binary system, and Mathieu et al. (1996)
identified UZ Tau E to be a single-lined spectroscopic binary with a 19.1-day period. UZ Tau
W, a 0.′′34 binary (47.6 AU assuming a distance of 140 pc to Taurus-Auriga; Kenyon et al.
1994), is separated from UZ Tau E by 3.′′78 (530 AU; Simon et al. 1995a). Prato et al. (2002)
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detected absorption lines of the secondary star in the near infrared spectrum of UZ Tau E,
measuring the mass ratio M2/M1 = 0.28 ± 0.01. Mart́ın et al. (2005) presented additional
radial velocity data for UZ Tau E; they found a binary orbital period of 18.979 days and an
eccentricity of 0.14. UZ Tau E shows strong Hα emission, indicative of ongoing accretion, and
strong infrared and millimeter excess emission from circumstellar and circumbinary disks;
the circumbinary disk has been resolved at λ = 1.3 mm (Jensen et al. 1996a) and 2.6 mm
(Dutrey et al. 1996) with a size and mass comparable to disks around single stars, showing
that the close spectroscopic companion has not significantly decreased the disk mass, in
contrast to the 50-AU pair in UZ Tau W, where the presence of a companion at a separation
comparable to typical disk sizes has greatly reduced the presence of circumstellar material.
In this paper, we present our new photometric observations, as well as a re-determination
of the binary orbital parameters from new and existing radial velocity observations. We
then examine periodicities in the data, showing that the photometric data vary at the binary
orbital period. Finally, we interpret the results in the context of the model of pulsed accretion
in binary systems.
2. Observations
2.1. Photometry
In order to search for periodic photometric variations, we have obtained new photometry
of UZ Tau E. Our photometric observations were made with the 0.6-m Perkin Telescope
at the Van Vleck Observatory (VVO) at Wesleyan University, and with ANDICAM on the
1.3-m SMARTS Telescope at CTIO. See Table 1 for details of the observations. The data
were reduced using standard techniques.
Since UZ Tau E and UZ Tau W are separated by 3.′′78, we sought to minimize con-
tamination of the UZ Tau E photometry by light from UZ Tau W by rejecting images with
FWHM greater than 7 pixels, and by using a relatively small (3-pixel-radius) photometric
aperture. Light curves of UZ Tau E and W show no correlation, indicating that the UZ Tau
E photometry is uncontaminated.
We performed differential photometry on UZ Tau E using USNO-B1.0 1158-0057597
as a comparison star. The UZ Tau field is relatively sparse, and in some of the SMARTS
images, this was the only star that was sufficiently bright to serve as a comparison star,
so we used it as the sole comparison in all of our photometry. The star was verified to be
non-variable at the few percent level, showing a standard deviation of 0.02 magnitudes by
comparison with several other stars of similar brightness using the wider-field VVO images.
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The USNO-B1.0 Catalog (Monet et al. 2003) gives magnitudes for this star from the Second
Palomar Sky Survey of B = 16.39, R = 13.75, and I = 12.8. Although these filters are not
identical to those used in our CCD observations, we adopted these values for the magnitude
of the reference star to set the zero point of our light curves. In addition, we adopted
V = 14.99 by noting that the R − I color for the comparison star suggests a spectral type
of ∼ M0, and adopting a corresponding B − V color. Adopting these values allows us to
determine approximate colors for UZ Tau E, though we caution that the absolute scaling
of both the individual magnitudes and the colors is uncertain. The differential photometry
and the color changes, which form the basis of our analysis, are both unaffected by this
systematic uncertainty in the zero point.
2.2. Spectroscopy
We did not acquire new spectra of UZ Tau E solely for this program, but we did
make new measurements of a number of spectra taken during the course of other programs.
Some of these spectra were kindly supplied by Marcos Huerta. These are echelle spectra
from McDonald observatory spanning 14 nights in January 2002, with R = 46, 000 and
wavelength coverage of 5460–6760 Å. The observations and data reduction are described in
detail in Huerta et al. (2005). Additional echelle spectra of UZ Tau E were taken at Keck
(R = 31, 000) with the setup described in Basri & Reiners (2006) and at Lick (R = 48, 000)
with the setup described in Alencar & Basri (2000).
We used the spectra from Huerta to measure radial velocities of UZ Tau E. Spectra of
the weak-lined T Tauri star V819 Tau were used as a radial velocity standard. By cross-
correlating the UZ Tau E spectra against the V819 Tau spectra, we measured heliocentric
radial velocities of UZ Tau E, assuming vhelio = 14.4±1.5 km s
−1 for V819 Tau (Walter et al.
1988). The resulting velocities are given in Table 2. Radial velocities were measured using
several different echelle orders with strong absorption lines; the quoted uncertainties reflect
the dispersion in these different measurements, as well as the uncertainty of V819 Tau’s
radial velocity.
In addition, we also measured the equivalent width of the Hα line in both sets of spectra
in order to track changes in accretion rate over time. Equivalent widths are given in Table
2, with an estimated uncertainty of 10%.
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3. Binary orbital parameters
In order to assess whether or not any periodic photometric variations detected in UZ
Tau E are synchronized with the binary orbit, we need to have an accurate knowledge of
the orbital parameters. These have been determined previously by Mathieu et al. (1996),
Prato et al. (2002), and Mart́ın et al. (2005), but the number of radial velocity points avail-
able is still relatively small, especially for the secondary, leaving open the prospect of further
improvements to the orbital parameters. To that end, we have re-analyzed the spectroscopic
orbit using data published in Prato et al. (2002) and Mart́ın et al. (2005) as well as our new
radial velocity measurements (Section 2.2).
We fit the radial velocity data using the Binary Star Combined Solution software
(Gudehus 2001), the ORBIT code (Forveille et al. 1999), and our own custom-written IDL
code; all gave the same solution. The best-fit phased radial velocity curve is shown in Figure
1 and the orbital parameters are given in Table 3.
The best-fit period of 19.131±0.003 days is inconsistent with the value of 18.979±0.007
days found by Mart́ın et al. (2005). Examination of the power spectrum of the velocity data
used by Mart́ın et al. shows that the 18.979-day period appears to be an alias of the true
period, caused by beating of the 19.131-day period with two six-year gaps in the radial
velocity data; there is a corresponding alias at 19.3 days. Re-fitting only the data used by
Mart́ın et al., we find that the two periods both correspond to local minima in χ2 space, with
reduced χ2 = 8.8 for P = 19.131 days and reduced χ2 = 10.7 for P = 18.979 days. When
the new radial velocity data are added, the fit for P = 19.131 days improves to reduced
χ2 = 8.1 while that for P = 18.979 days worsens to χ2 = 11.7, as expected if 19.131 days is
the correct period.
We note that we have not added any additional radial velocity measurements of the
secondary, and thus the mass ratio remains more uncertain than the other orbital elements,
resting on the six secondary radial velocities presented by Prato et al. (2002).
UZ Tau E is one of only a handful of pre–main-sequence systems with measured stellar
masses (see Mathieu et al. 2006 for a recent review). Because the total system mass has been
measured (Simon et al. 2000), the spectroscopic orbital parameter M sin3 i can be used to
determined the orbital inclination. This can then be compared with the observed inclination
of the circumbinary disk. While this was done by Simon et al. (2000) and Prato et al.
(2002), we revisit this issue here using our newly-determined orbital parameters for UZ
Tau E. Combined with M = 1.31 ± 0.08 M⊙ (Simon et al. 2000), our orbital parameters
give sin iorbit = 0.81± 0.05, or iorbit = 54
◦ ± 5◦. This is in excellent agreement with the disk
inclinations of 54◦±3◦ and 56◦±2◦ measured from interferometric images of the λ = 1.3 mm
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continuum emission and the CO line emission, respectively (Simon et al. 2000). Thus, the
binary orbit and the circumbinary disk are coplanar. Since the disk inclination is measured
at scales of ∼ 100 AU and the binary orbit is only a few tenths of an AU, this co-planarity
apparently extends over the entire disk. Though there are several theoretical studies of
how tilted circumstellar disks interact with a binary system (e.g., Papaloizou & Terquem
1995; Larwood et al. 1996; Bate et al. 2000; Lubow & Ogilvie 2000), we know of no studies
of the timescale for alignment of a circumbinary disk if it is initially tilted relative to the
binary orbit. Studies of the effects of a planetary-mass companion on a tilted external
disk (Lubow & Ogilvie 2001) suggest that such disk tilts do decay over time, however. The
example of UZ Tau E shows that co-planarity over the entire disk exists already by an age
of a few Myr, suggesting that circumbinary disks either form already aligned with the orbit,
or come into alignment very quickly. This is similar to the result of Jensen et al. (2004),
who found that circumstellar disks in young binaries tend to be aligned with each other, and
thus presumably with the binary orbit.
4. Periodic variations
4.1. Photometry
In order to determine whether the system is varying in phase with the binary orbit, or
in any other systematic way, we have searched the photometric data for periodic signals. We
begin by searching for evidence of periodicity, without pre-supposing a particular period.
A Lomb-Scargle periodogram (Scargle 1982) of the I-band data (the band with the largest
number of points and best time coverage; Table 1) is shown in Figure 2. There is a strong peak
at a period of 19.20± 0.03 days. The false-alarm probability (FAP) of this peak is less than
0.001 according to the formulation of Horne & Baliunas (1986). While this FAP calculation
is strictly applicable only to evenly-spaced data, a Monte Carlo bootstrapping method (e.g.,
Stassun et al. 1999) confirms that this period is statistically significant at better than 99.9%
confidence. The period uncertainty reported above, which follows from the formulation of
Kovacs (1981), is probably underestimated as it assumes that the underlying signal is well
described by a single sinusoid.
Though it is slightly off the main peak of the power spectrum, there is significant power
at the binary period of 19.131 days. Periodograms of the B, V , and R data are similar
(Figure 3), showing peaks near the binary period, but with broader peaks and higher false-
alarm probabilities, perhaps due to the more limited time coverage of the data in those
bands.
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To refine the period and to better estimate its uncertainty we next performed a phase
dispersion minimization (PDM) analysis (Stellingwerf 1978), which is particularly well-suited
to periodic variability that is highly non-sinusoidal and/or to data with large intrinsic scatter,
both of which apply to the photometry of UZ Tau E. The PDM search of the I-band data
yields a best period of P = 19.17±0.05 d, where the uncertainty was determined empirically
from the 1/e folding scale of the PDM merit function. The same analysis on the V -band
data gives P = 19.15± 0.04 d.
Schwarzenberg-Czerny (1989) argues that a related test, the one-way analysis of vari-
ance, is the most powerful statistic of this kind for detection of periodic signals. Applying
that test to our data yields P = 19.16± 0.03 d for the I-band data, and P = 19.15± 0.05 d
for the V -band data. Following Schwarzenberg-Czerny (1989), the period uncertainty was
determined using a “post-mortem” analysis that measures the 1-σ confidence interval of the
primary periodogram peak, defined by its width at the mean noise power level of the peri-
odogram in the vicinity of this peak. As above, a Monte Carlo permutation analysis of the
light curves confirms that this period is statistically significant at better than 99.9% confi-
dence. Combining these estimates, our best-fit photometric period is P = 19.16 ± 0.04 d,
consistent with the binary orbital period.
Figure 4 shows I-band light curves for all three observing seasons, folded at the binary
orbital period of 19.131 days. As suggested by the periodogram analysis, all show indications
of periodic behavior, with a broad minimum near orbital phase 0.5. The data from the 2004–
2005 season show the smoothest variability, but this is also the season with the smallest
number of data points. Clearly there is significant random variability as well, with scatter
of roughly 0.6 magnitudes at all orbital phases.
Figure 5 shows folded B, V , R, and I lightcurves from the 2003–2004 and 2005–2006
seasons. Broadly speaking, the BV R data show the same behavior as the I-band light curves,
with large-amplitude variability that appears to have both periodic and random components.
We note that the R band includes the Hα line, which may complicate the interpretation of
the light curve.
All four bands show a gradual increase in brightness over the three years of our ob-
servations, with the mean magnitude changing by 0.3 mags at I band from 2003–2004 to
2005–2006. To separate long-term variations from the shorter-term variations of interest
here, a linear trend (with a slope of roughly 0.15 mag / yr at I band) has been fit to each
band. The folded light curves using these de-trended data, and combining all three observing
seasons, are shown in Figure 6.
In addition to our photometric data, previous data on UZ Tau have shown some evidence
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of both long-term trends and periodic variations. Bohlin (1923), in one of the first papers
to mention UZ Tau, reported on a major flare and then a four-magnitude overall decline in
brightness from 1921–1923. He also noted that there was a short-period variation with a
period of 10–20 days, which encompasses the period of the variations reported here. Bohlin’s
measurements are for the entire UZ Tau system, but later examination of Bohlin’s plates by
Herbig (1977) showed that it was UZ Tau E that brightened dramatically in 1921.
Variations in color of the system can also give clues about the cause of the variability.
Figure 7 shows the V − I color as a function of I magnitude and orbital phase. The system
shows a behavior commonly seen in T Tauri stars, appearing redder when fainter and bluer
when brighter (Herbst et al. 1994). This behavior is consistent either with periodic changes
in extinction (causing both dimming and reddening when the extinction is higher) or in
accretion (adding additional blue light when the accretion rate is higher).
4.2. Spectroscopy
If the periodic photometric variations are due to changes in accretion rate, one might
expect accompanying variations in Hα emission or spectral veiling, common tracers of accre-
tion. Mart́ın et al. (2005) searched for both of these in UZ Tau E and did not find evidence
of either. With our new spectra, we can revisit the question of Hα variability.
Figure 8 shows the equivalent width of the Hα emission line as a function of binary
orbital phase. The variations do not appear to be strongly correlated with orbital phase.
There is some evidence for lower Hα equivalent widths around phase 0.4–0.8, as seen in the
photometric data, but the data are relatively sparse in that phase range as well.
While a lack of periodic variability would be at odds with the photometric data, we note
that periodicity may not be as obvious in the spectroscopic data, since the two datasets differ
in two important respects. First, the Hα data have much sparser sampling; they span a total
of eight years (1994–2002), but with only a handful of points during a given year. Second,
they do not overlap at all with the photometric data. Given that the photometric data show
both long-term trends and short-term scatter in addition to the periodic variations, and that
T Tauri stars in general are known to show significant random variability at Hα, it may be
difficult to separate random and periodic variations (if any) without a dedicated monitoring
campaign, preferably one that includes simultaneous photometric and spectroscopic mea-
surements. We conclude that while periodic spectroscopic variations similar to those seen
in the photometry are not definitively present in these spectroscopic data, neither are they
ruled out.
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5. Discussion
The photometric data (and possibly the spectroscopic data) show periodic variability
at the binary orbital period, suggesting that there is a link between the variability and in-
teractions of the binary with its circumstellar and/or circumbinary material. In this section,
we first argue that the periodic variations are unlikely to be due to stellar rotation, and
then we examine how well the observed behavior matches what is expected from the pulsed
accretion model of Artymowicz & Lubow (1996). Finally, we examine the available data for
other spectroscopic binaries to assess whether or not there is evidence for periodic accretion
as a general phenomenon.
5.1. Could the variations be due to rotation?
Periodic variability is not uncommon in photometric studies of PMS stars. Indeed, dedi-
cated monitoring surveys of rich star-forming regions (e.g., Mandel & Herbst 1991; Attridge & Herbst
1992; Choi & Herbst 1996; Stassun et al. 1999; Rebull 2001; Herbst et al. 2002a) have now
discovered hundreds of PMS stars exhibiting periodic variability, the result of surface-
brightness inhomogeneities (i.e. starspots) that rotate in and out of view with the stellar
rotation period. The periodic variability observed in UZ Tau E is very unlikely to be the
result of such rotationally modulated spot signals, for several reasons. First, the rotation
periods of low-mass PMS stars are nearly always shorter than about 12 days, while the pe-
riod of the variations reported here is 19 days. Among 150 low-mass PMS stars in the Orion
Nebula Cluster, only two stars (∼ 1%) have Prot > 15 d (Herbst & Mundt 2005). Second,
rotationally modulated spot signals are typically sinusoidal, and stable over many cycles or,
in some cases, many years. In contrast, the periodic signal we have found in UZ Tau E
is decidedly non-sinusoidal, with considerable scatter; the “bright” state has a duty cycle
of ∼ 60%. Thus, it is either intrinsically non-sinusoidal, or shows substantial phase shift-
ing from one cycle to the next; neither of these is consistent with rotationally-modulated
variability.
The rotation period distributions discussed above are presumably dominated by single
stars or member of wide binaries, while tidal interactions between the stars in a close binary
system can synchronize the orbital and rotational periods. However, in the case of eccentric
systems like UZ Tau E, pseudo-synchronization (in which the stellar angular velocity is
synchronized with the orbital angular velocity at periastron) occurs instead, since the tidal
interactions are strongest around periastron (Hut 1981). The predicted pseudo-synchronous
rotation period for UZ Tau E, using the weak friction formulation of Hut (1981) and the
orbital parameters in Table 3, is 11.4 ± 1.2 d, inconsistent with the observed variability
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period.
We can also estimate the rotation period of the UZ Tau E primary directly if three
quantities are known: the inclination irot of the star’s rotation axis, the star’s projected
rotational velocity v sin irot, and the stellar radius. Of these, the inclination is typically
impossible to measure, except under special circumstances.
As noted in Section 3, the dynamical mass measurement of UZ Tau E allows us to
determine the binary orbital inclination iorbit. Based on studies of other binary systems, it
is reasonable to assume that this inclination is the same as that of the stellar rotation axis,
irot. The most detailed study comparing the orientations of these axes in binary systems
is that of Hale (1994). Considering spectral types of F5–K5, he finds that binaries with
separations less than 30–40 AU tend to exhibit co-planarity between rotational equators and
orbital planes, while wider binaries have random orientations. Using a similar method, Weis
(1974) found a tendency for the stellar rotational equators to align with the binary orbit
among primaries in F star binaries. Interestingly, Weis (1974) did not find a tendency toward
co-planarity between rotational and orbital planes among A stars, suggesting that caution
is necessary when comparing stars of different masses. Similarly, Guthrie (1985) found no
correlation between orbital inclination and v sin i among 23 A2–A9 binaries with semi-major
axes of 10–70 AU. The low mass and short period of UZ Tau E suggest, however, that the
conclusions of Hale (1994) are most applicable here.
Prato et al. (2002) find L = 0.63+0.19
−0.17
L⊙ and Teff = 3700 ± 150 K for the primary in
UZ Tau E. Combining these values yields R = 1.9 ± 0.2 R⊙. Hartmann & Stauffer (1989)
find v sin i = 15.9±4.0 km s−1 for UZ Tau E using optical spectra, consistent with the value
v sin i = 16 ± 2 km s−1, which we measure from our new spectra and adopt here. Since
absorption lines of the secondary of UZ Tau E have only been seen in near-infrared spectra
and are not evident in any of our optical spectra, we take this to be the projected rotation
velocity of the primary. Combining these measurements with sin iorbit = 0.81±0.05 (Section
3), and assuming iorbit = irot, we find Prot = 4.9± 0.8 d. If iorbit 6= irot, we find Prot ≤ 6± 1 d
since sin i ≤ 1. Thus, uncertainty on the inclination cannot reconcile the photometric period
with the inferred rotation period.
The most uncertain remaining quantity is v sin i, but since Hartmann & Stauffer (1989)
measured v sin i from 11 different spectra of UZ Tau E, with self-consistent results from two
different parts of the spectrum (including spectra near λ = 5200 Å) and consistency with our
new v sin i measurement, it is unlikely that line broadening from photospheric lines of the
faint, red secondary could lead to an overestimate of v sin i by a factor of four. Similarly, given
the uncertainties on L and Teff , it is difficult to see how the radius could be underestimated
by a factor of four. Thus, we conclude that the observed periodic variations are unlikely to
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be due to stellar rotation.
5.2. Evidence for pulsed accretion
We have shown above that UZ Tau E exhibits periodic photometric variations that have
the same period as the binary orbit, and that these variations are unlikely to be caused by
stellar rotation. Here, we examine the predictions made by the pulsed accretion model of
AL96 and compare them to our observations.
5.2.1. What are the predictions?
Broadly speaking, Artymowicz & Lubow (1996) predict that a binary with an eccentric
orbit and a circumbinary disk will have an accretion flow from the circumbinary disk, and
thus onto the circumstellar disks or stellar surfaces, that varies periodically at the binary
orbital period.
The exact behavior of the accretion rate with orbital phase depends on the binary orbital
parameters. AL96 show the results of two simulations, one for mass ratio M2/M1 = 0.43 and
eccentricity e = 0.1, and another for M2/M1 = 0.79 and e = 0.5. The former shows accretion
that varies relatively smoothly over the orbital period, while the latter is strongly peaked
at periastron. As noted by AL96, the exact timing of the accretion variability depends
on the orbital parameters, most strongly on e. Some previous observational studies of T
Tauri spectroscopic binaries have focused specifically on looking for enhanced accretion near
periastron; however, we note here that the actual prediction of the model is more general
than that, and that the peak accretion rate need not come near periastron.
5.2.2. How well do the data match the predictions?
First, we note that our observations match the general predictions of AL96 quite well,
in that there are indeed periodic photometric variations at the binary orbital period, which
are readily interpretable as a variable accretion rate. The comparison with the spectroscopic
data is more ambiguous; if more intensive monitoring of the Hα line in UZ Tau E were to
show that there are no orbit-modulated Hα variations, it would present a problem for the
model.
For a more specific comparison with our data, Figure 9 shows the variations of accretion
– 13 –
with orbital phase predicted by AL96 for a binary with M2/M1 = 0.43, e = 0.1. UZ Tau E
has a more extreme mass ratio (M2/M1 = 0.30) and larger eccentricity (e = 0.33) than this,
but these parameters are closer to those of UZ Tau E than those of the other simulation in
AL96. AL96 do note that the timing of the maxima of the accretion depend largely on e
rather than M2/M1. Since e for UZ Tau E is intermediate between the two models calculated
by AL96, we might then expect the maximum accretion to come between the phase of ∼ 0.75
they calculate for the low-e case and the phase of ∼ 1 for the high-e case.
For comparison with our data, we have taken the logarithm of the variations of accretion
rate predicted by AL96 to shift them onto a “magnitude-like” scale, and added an arbitrary
offset and scale factor to match the mean of the data and amplitude of the variations. The
phase of the minimum predicted by this simulation does not match our data well; when the
model is given a shift of +0.2 in orbital phase, there is better agreement between the model
predictions and the data. This scaling and shifting to match the data is obviously ad hoc,
but it allows us to compare the phase width of the observed variations, which appear to
match the predictions relatively well. In addition, this shifted position of the maximum is
indeed between the two cases calculated by AL96, as expected if eccentricity is the dominant
factor in determining the timing of maximum accretion.
5.3. Evidence for periodic accretion in other T Tauri binaries
The discussion and data above show that looking for evidence of periodic accretion can
be complicated, with other sources of variability perhaps being important and masking the
effect in small datasets, and with the exact behavior expected to be a function of the specific
binary orbital parameters. That said, is evidence for pulsed accretion seen in other young
binary systems? In Table 4 we present characteristics of young binaries with periods of
less than one year and evidence of circumbinary material, in order of increasing eccentricity.
Below, we examine the observational data for some of these systems, attempting to relate
them to what we see in UZ Tau and exploring similarities and differences. Unfortunately,
the small number of systems and their somewhat heterogeneous properties means that it is
difficult to generalize, so we offer these comments in the spirit of attempting to pull together
the existing data, rather than arguing one way or the other for the validity of the AL96
model for the sample as a whole.
– 14 –
5.3.1. DQ Tau
DQTau was the first system to be scrutinized for evidence of pulsed accretion. Mathieu et al.
(1997) showed that the photometric variations are modulated at the binary orbital period,
and Basri et al. (1997) showed that the Hα line and spectral veiling are as well. Fortuitously,
the mass ratio and eccentricity of DQ Tau are quite similar to those of the high-eccentricity
case modeled by AL96, allowing for specific comparison with the theory. The timing and
phase width of the photometric and spectroscopic variations match the predictions well,
being sharply peaked near periastron. However, the DQ Tau observations did show consid-
erable orbit-to-orbit variation, with the periastron brightening being seen roughly 65% of
the time. This is reminiscent of the large scatter that we see in the UZ Tau light curves;
clearly the periodic accretion process is not exactly repeatable from orbit to the next, nor is
it the only source of variability.
5.3.2. AK Sco
The orbital eccentricity and binary mass ratio of AK Sco are quite similar to those of DQ
Tau, and indeed, simulations by Günther & Kley (2002) for a binary with AK Sco’s orbital
parameters predict pulsed accretion. Thus, it comes as some surprise that the system does
not show periodic photometric variability, despite extensive monitoring (Alencar et al. 2003).
The overall variability is large (up to 1.5 mags in y), but apparently random. There are
periodic variations in the Balmer lines, but they are not sharply peaked around periastron.
Examining Table 4, we note two properties of AK Sco that are quite different from those
of DQ Tau or UZ Tau. First, AK Sco is considerably hotter and more luminous. Thus,
accretion variations of the same luminosity as those occurring in UZ Tau and DQ Tau would
result in substantially smaller magnitude changes, which could be swamped by the large
random variability. Second, AK Sco has considerably lower millimeter flux than either of
the other two systems. If the systems are fit with similar disk models (in which the disk
is assumed to be optically thin at millimeter wavelengths in its outer regions), AK Sco’s
disk mass is an order of magnitude smaller than that of DQ Tau or UZ Tau E (Jensen et al.
1996a; Jensen & Mathieu 1997; Mathieu et al. 1997). Alencar et al. (2003) fit AK Sco with
an optically-thick disk model that has a comparable mass to the disk models fit to DQ Tau
and UZ Tau E. However, such disk models have not been fit to DQ Tau or UZ Tau E, and
would presumably result in even larger disk masses for those systems. In a direct comparison
of λ = 1.1 mm flux, DQ Tau and UZ Tau E are 3 and 5 times brighter than AK Sco at
roughly the same distance, presumably reflecting larger disk masses. It is possible that a
somewhat lower-mass disk has different dynamics, and that the accretion flow in the AK Sco
– 15 –
system is fundamentally different than that in the other systems with more massive disks.
5.3.3. GW Ori
This system has a near-circular orbit, and thus would not be expected to show pulsed
accretion under the model set forth by AL96. However, Stempels & Gahm (2004) quote
Artymowicz, private communication, as saying that pulsed accretion is possible for systems
with circular orbits as well, and indeed D’Angelo et al. (2006) show that this occurs for giant
planets embedded in disks. Thus, pulsed accretion appears to be possible for at least some
circular-orbit systems and thus may be for GW Ori as well, though the larger gap cleared by
a stellar companion (Artymowicz & Lubow 1994) will clear some of the disk resonances that
might contribute to disk eccentricity growth in a system with a planetary-mass companion.
Like AK Sco, GW Ori is very luminous, and shows significant random variability, though
no obvious periodic variability. It does have a much more massive disk than AK Sco, however,
and indeed than any of the systems considered here. Because of its much larger semimajor
axis, and to some extent its circular orbit, GW Ori is much more likely to have significant
circumstellar disks, as the stars do not approach each other very closely at periastron. Thus,
material flowing from the circumbinary disk may merge with the circumstellar disks and then
accrete more gradually onto the stars, rather than falling directly on (or near) the stellar
surfaces as is expected to happen in the shorter-period systems. If the infalling material does
not shock strongly as it merges with the circumstellar disk, and if any density enhancements
are smoothed out somewhat by the time the material reaches the stellar surface, then any
photometric signature of the periodic infall would be weakened. We note that UZ Tau E
likely has circumstellar disks as well (Jensen et al. 1996a), so a similar effect could be at work
in reducing the amplitude of the periodic variability relative to the stochastic variability.
5.3.4. V4046 Sgr
Like GW Ori, V4046 Sgr has a nearly circular orbit. However, V4046 Sgr has shown
periodic photometric variations at the binary orbital period (Quast et al. 2000; Mekkaden
2000). These variations persist over several years and are relatively sinusoidal (Walter,
unpublished data, 2003–2005). Unlike the other binaries discussed here, in this case stellar
rotation is a plausible explanation of the observed variations. It is common for stellar
rotational periods to become synchronized with the binary orbital period, particularly for
short-period binaries like V4046 Sgr. Given the short period (resulting in stronger tidal
– 16 –
interactions and a shorter synchronization time scale) and the somewhat older age of this
system (∼ 10 Myr), synchronization is plausible, and indeed is supported by detailed analysis
of the system (Stempels & Gahm 2004). However, rotation does not explain the periodic
Balmer line variations observed, which Stempels & Gahm (2004) attribute to accumulations
of gas co-rotating with the binary orbit.
5.3.5. ROXs 42 and ROXs 43B
These two spectroscopic binaries are both weak-lined T Tauri stars (Bouvier & Appenzeller
1992; Walter et al. 1994), indicating less-active accretion than some of the other systems dis-
cussed here. Neither has been detected at millimeter wavelengths, yielding only an upper
limit on the disk masses (Skinner et al. 1991; Jensen et al. 1996b). Both systems show
mid-infrared excesses, indicating the presence of circumbinary material, and a lack of near-
infrared excess, which can be modeled as a cleared central region in the disk (Jensen & Mathieu
1997). The fact that both are higher-order multiple systems complicates matters; ROXs 42
(NTTS 162814−2427) is a triple system with a separation of 0.′′15 (Lee 1992; Ghez et al.
1993), while ROXs 43B (NTTS 162819−2423S) has a wide companion at 4.′′8 which is itself a
close binary system (Walter et al. 1994; Simon et al. 1995b). Since the evidence for the pres-
ence of a substantial disk rests on the low-spatial-resolution IRAS detections, it is possible
that the excess is associated with the wider companions rather than arising from circumbi-
nary disks around the spectroscopic binaries. In any case, the lack of millimeter detections
indicates that there is less disk mass in these two systems than in the others discussed here.
Neither system has been intensively monitored over timespans that would be necessary to
detect periodic photometric variations at the relatively long orbital periods. ROXs 42 shows
evidence for some semi-regular variations over roughly 1.5 orbital periods (Zakirov et al.
1993), while the combined light of the ROX 43 system shows only a 0.1-magnitude varia-
tion, with evidence of a 1.5-day or 3-day periodicity, presumably attributable to rotation of
one or more of the stars (Shevchenko & Herbst 1998).
5.3.6. KH 15D
The unusual pre–main-sequence system KH 15D (V582 Mon) is a spectroscopic bi-
nary that undergoes deep (∆I ∼ 3.5 mag) eclipses, thought to arise due to occultation
from a circumbinary disk (Hamilton et al. 2001; Herbst et al. 2002b; Hamilton et al. 2005;
Winn et al. 2006 and references therein). While the system has an eccentricity and mass
ratio that would suggest that pulsed accretion might be present, the photometric variations
– 17 –
at the binary orbital period are dominated by the deep eclipses. Furthermore, the depth and
detailed shape of these eclipses are evolving with time (Winn et al. 2003; Johnson & Winn
2004; Maffei et al. 2005; Johnson et al. 2005; Winn et al. 2006), making it very difficult to
determine whether there might currently be an additional, smaller-amplitude component
with the same period that is related to accretion rather than occultation. Winn et al. (2003)
showed that the current deep eclipses did not occur during the first half of the twentieth
century, raising the possibility of searching for evidence of accretion-related variability at
earlier epochs. Their limit of one mag on the variability during that time does not preclude
accretion-related variations like those seen in UZ Tau E. The ∼ 0.9-mag periodic varia-
tions seen from the 1960’s through the 1980’s (Johnson & Winn 2004; Maffei et al. 2005;
Johnson et al. 2005), however, are relatively smooth and are well-fit by the eclipse model
(Winn et al. 2006), placing a limit on how much any accretion-related component was con-
tributing to the variability during that time. Since the inferred mass ratio and eccentricity
for KH 15D are similar to those of DQ Tau (Table 4), we might expect accretion-related vari-
ability to be strongly peaked around periastron, which is also when the current deep eclipses
occur. This might help explain several anomalously bright points seen during eclipses in the
late 1990’s that are not well-fit by the model of Winn et al. (2006).
The precessing circumbinary disk occultation model of Winn et al. (2006) is quite suc-
cessful in reproducing the shape and ongoing evolution of the light curve, and we do not
suggest that accretion explains most of the photometric variations. We note, however, the
possibility that such an additional component might be sporadically present (with the same
period) and that, if it is, this could complicate the modeling of the historical evolution of
the light curve, especially during earlier, more-sparsely-sampled epochs.
6. Conclusions
We have shown that the pre–main-sequence binary UZ Tau E shows clear photometric
variability at the binary orbital period of 19.13 days. This variability is consistent with a
model in which material in the circumbinary disk is periodically perturbed by the binary in
its eccentric orbit and falls from the outer disk, across the cleared central gap and onto the
stars or their circumstellar disks. There is significant scatter in the light curves, indicating
that this “pulsed accretion” may not occur during every binary orbit. Hα equivalent widths
show some suggestion of periodic variability, but it is not definitive.
The apparently intermittent behavior of the accretion, and the presence of other, random
sources of variability, suggest that searches for this sort of accretion signature require well-
sampled datasets with long time baselines in order to detect any periodic component. In
– 18 –
particular, simultaneous photometric and spectroscopic monitoring of UZ Tau E in the future
will help determine whether the Hα variations show a periodic component, as the photometric
variations do.
The good overall agreement between theory and observations suggests that resonant
interactions between stars (and, by extension, planets) and disks are indeed important in
determining disk structure and dynamics, while the random component of the observed
behavior shows that there is still work to be done in understanding the full complexity of
these interactions.
We gratefully acknowledge the support of the National Science Foundation through
grant AST-0307830. We thank the referee, Steve Lubow, for useful comments that improved
this paper. We are grateful to Michael Meyer, David Cohen, and Larry Marschall for useful
discussions; to Marcos Huerta and Pat Hartigan for use of their spectra of UZ Tau E; to
Matthew Richardson for assistance with data reduction; to Peter Collings for translating
early papers on UZ Tau from German to English; and to Thierry Forveille for use of his
ORBIT code. MS and FW are grateful for Stony Brook University’s partial support of
their participation in the SMARTS consortium. This research has made use of the SIMBAD
database, operated at CDS, Strasbourg, France, and of NASA’s Astrophysics Data System.
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This preprint was prepared with the AAS LATEX macros v5.2.
– 23 –
Table 1. Observations of UZ Tau E
Telescope Exp. time (s) Filter(s) Season # of nights Timespan in days
SMARTS (1.3m) 5 BV RI 2003–2004 63 170
2005–2006 6 42
30 V RI 2005–2006 9 33
VVO (0.6m) 60 I 2004–2005 16 126
2005–2006 20 128
– 24 –
Table 2. Radial Velocities and Hα EW
Julian Date vhelio (km s
−1) Hα EW (Å)a
2450416.82 · · · 88.4
2450783.93 · · · 42
2450783.94 · · · 45.1
2450784.96 · · · 38
2450785.11 · · · 57
2450835.65 · · · 54:
2451060.99 · · · 45.6
2451061.00 · · · 39.8
2451077.14 · · · 63.8
2451120.92 · · · 101
2451137.94 · · · 51.3
2451138.88 · · · 62.6
2451162.86 · · · 69.3
2451163.82 · · · 57.8
2451164.79 · · · 57.5
2451165.78 · · · 58.3
2451166.79 · · · 61.7
2451169.84 · · · 48.9
2451504.73 · · · 35
2451507.68 · · · 25:
2451508.67 · · · 49.1
2451509.71 · · · 57
2451510.66 · · · 71
2451517.98 · · · 75
2451523.68 · · · 42.7
2451524.74 · · · 49.3
2451525.76 · · · 40.8
2451527.68 · · · 46.6
2451528.68 · · · 58.9
2451529.66 · · · 61.8
2451530.60 · · · 69.1
– 25 –
Table 2—Continued
Julian Date vhelio (km s
−1) Hα EW (Å)a
2452280.60 −4.3 ± 2.1 44
2452281.74 −3.1 ± 1.5 51
2452282.76 −5.8 ± 1.8 81
2452283.66 −4.9 ± 4.6 78
2452284.71 · · · b 77
2452286.66 6.2 ± 2.6 90
2452287.65 17.1 ± 1.7 87
2452288.68 27.2 ± 2.1 59
2452289.69 29.6 ± 3.0 65
2452290.72 37.8 ± 7.0 67
2452291.67 28.5 ± 2.1 58
2452292.64 29.1 ± 3.2 50
2452293.67 22.0 ± 3.6 45
2452579.11 · · · 50
aPositive values denote emission.
bThe spectrum on this date was too noisy
to allow measurement of an accurate radial
velocity.
– 26 –
Table 3. Binary orbital parameters for UZ Tau E
Period (days) 19.131± 0.003
e 0.33± 0.04
JD of periastron 2451328.3± 0.5
ω 239◦ ± 9◦
a sin i (AU) 0.124± 0.003
γ (km s−1) 13.9± 0.7
K1 (km s
−1) 17.3± 1.4
K2 (km s
−1) 57.4± 4.7
M sin3 i (M⊙) 0.69± 0.13
M2/M1 0.30± 0.03
Table 4. CTTS spectroscopic binaries
Binary Period e M2/M1 Spectral L Disk Mass
a Photometric Balmer line References
System (days) Type (L⊙) (M⊙) periodicity? periodicity?
V4046 Sgr 2.421 ≤ 0.01 0.94 K5 0.82 0.0085 Yes (∆B≈0.1) Yes 1, 2, 3
GW Ori 241.9 0.04 SB1 G0 26 0.3 ? (∆V≈0.7) ? 4, 5
UZ Tau E 19.131 0.33 0.30 M1 0.91 0.063 Yes (∆I≈0.8) Maybe 6, 7, 8
ROXs 43B 89.1 0.41 SB1 G0 0.4 < 0.00037 ? (∆V=0.1) ? 1, 9, 10, 11
AK Sco 13.609 0.47 0.99 F5 8.40 0.002 No (∆y≈1.5) Yes 1, 12, 13
ROXs 42 35.95 0.48 0.92 K4 0.4 < 0.00025 ? (∆V=0.4) ? 1, 10, 11, 14, 15
DQ Tau 15.804 0.56 0.97 K7–M1 0.95 0.020 Yes (∆V≈0.5) Yes 16, 17
KH 15D 48.38 0.57–0.65 0.83b K7c 0.4c · · · Eclipse (∆I≈3.5) ? 18, 19, 20
aDisk mass estimates were made assuming that the disk is at least partially optically thin at mm wavelengths. An optically-
thick model for AK Sco (Alencar et al. 2003) yields a disk mass of 0.02 M⊙.
bDerived from the stellar luminosity ratio that best fits the eclipse data (Winn et al. 2006).
cProperties of the secondary star, since the primary is never visible.
References. — 1. Jensen & Mathieu (1997). 2. Quast et al. (2000). 3. Mekkaden (2000). 4. Mathieu et al. (1991). 5.
Mathieu et al. (1995). 6. This work. 7. Prato et al. (2002). 8. Jensen et al. (1996a). 9. Shevchenko & Herbst (1998).
10. Manset & Bastien (2003). 11. Bouvier & Appenzeller (1992). 12. Alencar et al. (2003). 13. Manset et al. (2005). 14.
Lee (1992) 15. Walter et al. (1994). 16. Mathieu et al. (1997). 17. Basri et al. (1997). 18. Hamilton et al. (2001). 19.
Hamilton et al. (2005). 20. Winn et al. (2006).
– 28 –
Fig. 1.— The best-fit spectroscopic orbit for UZ Tau E. Crosses show velocities of the pri-
mary; those enclosed in boxes (in red in the on-line edition) show new radial velocity mea-
surements presented here. Open diamonds show velocities of the secondary from Prato et al.
(2002).
– 29 –
Fig. 2.— The Lomb-Scargle periodogram for the I-band data. The periodogram peaks at
a period of 19.20 days, with a false-alarm probability of 0.001. The dashed line shows the
binary orbital period of 19.131 days. The smaller peaks visible flanking the main peak (lower
panel) are near the alias periods expected for beat periods between one-year and two-year
periods (caused by the seasonal gaps in the data) and the binary period.
– 30 –
Fig. 3.— The Lomb-Scargle periodograms for the B, V , R, and I-band data. All show
significant power near the binary orbital period.
– 31 –
Fig. 4.— The I-band magnitude for UZ Tau E folded at the binary orbital period of 19.131
days and plotted against the binary orbital phase. Top to bottom, data from 2003–2004,
2004–2005, and 2005–2006.
– 32 –
Fig. 5.— The BV RI magnitudes for UZ Tau E folded at the binary orbital period and
plotted against the binary orbital phase. Left, 2003–2004; right, 2005-2006. The open circles
in the lower right plot show the VVO I-band data, which do not have corresponding B, V,
and R data.
– 33 –
Fig. 6.— The BV RI magnitudes for UZ Tau E folded at the binary orbital period and
plotted against the binary orbital phase, after removing a long-term linear trend from each
band.
– 34 –
Fig. 7.— V − I color vs. I magnitude and vs. orbital phase for 2003–2004 (closed circles)
and 2005–2006 (open circles). The system is redder when fainter and bluer when brighter,
the expected behavior either for changes in extinction or for brightening due to increased
accretion.
– 35 –
Fig. 8.— Top: Equivalent width of the Hα line as a function of binary orbital phase. Squares
are our new measurements; triangles are from Mart́ın et al. (2005). Bottom: For comparison,
the phased I-band data. There is some suggestion of reduced Hα equivalent width at phases
of 0.4–0.8 as seen in the photometric data, but the data are too sparse there for there to be
clear evidence for periodic variability of the Hα emission.
– 36 –
Fig. 9.— Left: The theoretical predictions of Artymowicz & Lubow (1996) for the depen-
dence of accretion rate on binary orbital phase in a binary with mass ratio M2/M1 = 0.43,
e = 0.1. The top curve shows the total accretion, while the lower curves show accretion onto
the secondary (higher dark curve) and primary (lower dark curve). Right: The same total
accretion curve, but placed onto a logarithmic scale and shifted vertically for comparison
with the phased I-band data. The model here has been given an ad hoc shift of +0.2 in
phase, roughly what is expected given the binary eccentricity (see Section 5.2.2).
	Introduction
	Observations
	Photometry
	Spectroscopy
	Binary orbital parameters
	Periodic variations
	Photometry
	Spectroscopy
	Discussion
	Could the variations be due to rotation?
	Evidence for pulsed accretion
	What are the predictions?
	How well do the data match the predictions?
	Evidence for periodic accretion in other T Tauri binaries
	DQ Tau
	AK Sco
	GW Ori
	V4046 Sgr
	ROXs 42 and ROXs 43B
	KH 15D
	Conclusions

  Close pre-main-sequence binary stars are expected to clear central holes in
their protoplanetary disks, but the extent to which material can flow from the
circumbinary disk across the gap onto the individual circumstellar disks has
been unclear. In binaries with eccentric orbits, periodic perturbation of the
outer disk is predicted to induce mass flow across the gap, resulting in
accretion that varies with the binary period. This accretion may manifest
itself observationally as periodic changes in luminosity. Here we present a
search for such periodic accretion in the pre-main-sequence spectroscopic
binary UZ Tau E. We present BVRI photometry spanning three years; we find that
the brightness of UZ Tau E is clearly periodic, with a best-fit period of 19.16
+/- 0.04 days. This is consistent with the spectroscopic binary period of 19.13
days, refined here from analysis of new and existing radial velocity data. The
brightness of UZ Tau E shows significant random variability, but the overall
periodic pattern is a broad peak in enhanced brightness, spanning more than
half the binary orbital period. The variability of the H-alpha line is not as
clearly periodic, but given the sparseness of the data, some periodic component
is not ruled out. The photometric variations are in good agreement with
predictions from simulations of binaries with orbital parameters similar to
those of UZ Tau E, suggesting that periodic accretion does occur from
circumbinary disks, replenishing the inner disks and possibly extending the
timescale over which they might form planets.

Introduction
It is now well-established that most stars are members of binary systems at birth, and
that many of these stars are surrounded by disks similar to those found around young single
stars (see, e.g., the recent review by Monin et al. 2006). Thus, understanding the origin of
binaries is vital to understanding the star formation process. The predominance of binaries
also means that, based on number of systems alone, most potential sites of planet formation
lie in multiple systems. However, interactions between stars and disks in binary systems
can alter disk structure (Beckwith et al. 1990; Jensen et al. 1994; Osterloh & Beckwith 1995;
Jensen et al. 1996a,b; Jensen & Mathieu 1997), resulting in a more complicated environment
for planet formation. Nonetheless, the discovery of planets in relatively close binary systems
(Eggenberger et al. 2004) shows that binary systems are viable sites of planet formation.
Understanding the extent to which binaries modify the structure of their surrounding disks
is important for understanding the possible diversity of planetary system environments.
In addition, though the mass ratios are different, the interactions between stellar binary
companions and disks involve the same physics as those between planetary companions and
disks (e.g., D’Angelo et al. 2006) but are more easily observable. Thus, an understanding
of binary-disk interactions may help us understand planet formation around single stars as
well.
A binary star system may have up to three disks: two circumstellar disks, one each
around the primary and secondary; and a circumbinary disk outside the binary orbit. Both
analytic calculations and numerical simulations show that the region between these disks
is not stable for orbiting disk material. However, the question of how easily material can
flow from the outer, circumbinary disk across the gap to the circumstellar disks has not
been clear, either from an observational or a theoretical standpoint. The spectral energy
distributions of some young binaries show the clear signature of a cleared central region, while
other, apparently similar systems do not (Jensen & Mathieu 1997). Theoretical analyses by
Lin & Papaloizou (1993) and Artymowicz & Lubow (1994) suggested that material in the
region around the binary orbit is cleared, creating a quasi-equilibrium structure with three
– 3 –
distinct disks and a cleared region between them. If the gap between disks is impermeable,
the disks evolve independently of each other. Since the presence of a binary companion may
increase the rate of accretion (Clarke 1992; Ostriker et al. 1992) and the binary orbit presents
a constraint on the size of each circumstellar disk, the circumstellar disks would be exhausted
much more quickly than disks around single stars. However, smoothed particle hydrodynamic
simulations by Artymowicz & Lubow (1996) (hereafter AL96; see also Günther & Kley 2002)
predicted that material may indeed flow from the circumbinary disk to the circumstellar
environment, with the accretion rate varying with the phase of the binary orbit.
If such periodic accretion occurs in young binaries, it may be detectable observationally
by a periodic brightening of the system as the material flowing from the circumbinary disk
shocks when it collides with the circumstellar disk(s) or accretes onto the stellar surface(s).
Observations of the T Tauri spectroscopic binary DQ Tau by Mathieu et al. (1997) showed
such brightening, occurring at the binary orbital period. In addition, DQ Tau shows periodic
variations in spectral veiling and emission line intensities with orbital phase (Basri et al.
1997), providing strong support for the broad picture of mass flow across gaps suggested by
AL96. However, subsequent searches in other young, short-period binary systems surrounded
by disks have yielded mixed results. Alencar et al. (2003) did not find periodic photometric
variations in AK Sco, but they did find that the blue wing of the Hα line, and both the
blue and red wings of the Hβ line, vary with the binary orbital period, as does the total
Hα equivalent width. V4046 Sgr has shown periodic photometric variations at the binary
orbital period (Quast et al. 2000; Mekkaden 2000), though an earlier study did not find
such variations (Byrne 1986). Like AK Sco, however, V4046 Sgr does show variations in the
equivalent width and shape of Balmer lines as a function of orbital phase (Stempels & Gahm
2004). No dedicated photometric monitoring of UZ Tau E has been reported in the literature
to date, but in previous spectroscopic observations, neither Hα equivalent width nor spectral
veiling has shown any obvious dependence on binary orbital phase (Mart́ın et al. 2005).
Motivated by previous observational work and a desire to understand accretion in binary
systems, we have undertaken a photometric monitoring campaign for the pre–main-sequence
(PMS) spectroscopic binary UZ Tau E. UZ Tau, in the Taurus-Auriga star-forming region,
was first discovered to be variable by Bohlin during a bright outburst in 1921 (Bailey 1921;
Bohlin 1923) and identified as a ∼ 3.′′7 binary by Joy & van Biesbroeck (1944), one of the first
pre–main-sequence binaries to be identified. Subsequently, both components of the binary
were found to be binary themselves, making this a quadruple system. Simon et al. (1992)
and Ghez et al. (1993) identified UZ Tau W as a binary system, and Mathieu et al. (1996)
identified UZ Tau E to be a single-lined spectroscopic binary with a 19.1-day period. UZ Tau
W, a 0.′′34 binary (47.6 AU assuming a distance of 140 pc to Taurus-Auriga; Kenyon et al.
1994), is separated from UZ Tau E by 3.′′78 (530 AU; Simon et al. 1995a). Prato et al. (2002)
– 4 –
detected absorption lines of the secondary star in the near infrared spectrum of UZ Tau E,
measuring the mass ratio M2/M1 = 0.28 ± 0.01. Mart́ın et al. (2005) presented additional
radial velocity data for UZ Tau E; they found a binary orbital period of 18.979 days and an
eccentricity of 0.14. UZ Tau E shows strong Hα emission, indicative of ongoing accretion, and
strong infrared and millimeter excess emission from circumstellar and circumbinary disks;
the circumbinary disk has been resolved at λ = 1.3 mm (Jensen et al. 1996a) and 2.6 mm
(Dutrey et al. 1996) with a size and mass comparable to disks around single stars, showing
that the close spectroscopic companion has not significantly decreased the disk mass, in
contrast to the 50-AU pair in UZ Tau W, where the presence of a companion at a separation
comparable to typical disk sizes has greatly reduced the presence of circumstellar material.
In this paper, we present our new photometric observations, as well as a re-determination
of the binary orbital parameters from new and existing radial velocity observations. We
then examine periodicities in the data, showing that the photometric data vary at the binary
orbital period. Finally, we interpret the results in the context of the model of pulsed accretion
in binary systems.
2. Observations
2.1. Photometry
In order to search for periodic photometric variations, we have obtained new photometry
of UZ Tau E. Our photometric observations were made with the 0.6-m Perkin Telescope
at the Van Vleck Observatory (VVO) at Wesleyan University, and with ANDICAM on the
1.3-m SMARTS Telescope at CTIO. See Table 1 for details of the observations. The data
were reduced using standard techniques.
Since UZ Tau E and UZ Tau W are separated by 3.′′78, we sought to minimize con-
tamination of the UZ Tau E photometry by light from UZ Tau W by rejecting images with
FWHM greater than 7 pixels, and by using a relatively small (3-pixel-radius) photometric
aperture. Light curves of UZ Tau E and W show no correlation, indicating that the UZ Tau
E photometry is uncontaminated.
We performed differential photometry on UZ Tau E using USNO-B1.0 1158-0057597
as a comparison star. The UZ Tau field is relatively sparse, and in some of the SMARTS
images, this was the only star that was sufficiently bright to serve as a comparison star,
so we used it as the sole comparison in all of our photometry. The star was verified to be
non-variable at the few percent level, showing a standard deviation of 0.02 magnitudes by
comparison with several other stars of similar brightness using the wider-field VVO images.
– 5 –
The USNO-B1.0 Catalog (Monet et al. 2003) gives magnitudes for this star from the Second
Palomar Sky Survey of B = 16.39, R = 13.75, and I = 12.8. Although these filters are not
identical to those used in our CCD observations, we adopted these values for the magnitude
of the reference star to set the zero point of our light curves. In addition, we adopted
V = 14.99 by noting that the R − I color for the comparison star suggests a spectral type
of ∼ M0, and adopting a corresponding B − V color. Adopting these values allows us to
determine approximate colors for UZ Tau E, though we caution that the absolute scaling
of both the individual magnitudes and the colors is uncertain. The differential photometry
and the color changes, which form the basis of our analysis, are both unaffected by this
systematic uncertainty in the zero point.
2.2. Spectroscopy
We did not acquire new spectra of UZ Tau E solely for this program, but we did
make new measurements of a number of spectra taken during the course of other programs.
Some of these spectra were kindly supplied by Marcos Huerta. These are echelle spectra
from McDonald observatory spanning 14 nights in January 2002, with R = 46, 000 and
wavelength coverage of 5460–6760 Å. The observations and data reduction are described in
detail in Huerta et al. (2005). Additional echelle spectra of UZ Tau E were taken at Keck
(R = 31, 000) with the setup described in Basri & Reiners (2006) and at Lick (R = 48, 000)
with the setup described in Alencar & Basri (2000).
We used the spectra from Huerta to measure radial velocities of UZ Tau E. Spectra of
the weak-lined T Tauri star V819 Tau were used as a radial velocity standard. By cross-
correlating the UZ Tau E spectra against the V819 Tau spectra, we measured heliocentric
radial velocities of UZ Tau E, assuming vhelio = 14.4±1.5 km s
−1 for V819 Tau (Walter et al.
1988). The resulting velocities are given in Table 2. Radial velocities were measured using
several different echelle orders with strong absorption lines; the quoted uncertainties reflect
the dispersion in these different measurements, as well as the uncertainty of V819 Tau’s
radial velocity.
In addition, we also measured the equivalent width of the Hα line in both sets of spectra
in order to track changes in accretion rate over time. Equivalent widths are given in Table
2, with an estimated uncertainty of 10%.
– 6 –
3. Binary orbital parameters
In order to assess whether or not any periodic photometric variations detected in UZ
Tau E are synchronized with the binary orbit, we need to have an accurate knowledge of
the orbital parameters. These have been determined previously by Mathieu et al. (1996),
Prato et al. (2002), and Mart́ın et al. (2005), but the number of radial velocity points avail-
able is still relatively small, especially for the secondary, leaving open the prospect of further
improvements to the orbital parameters. To that end, we have re-analyzed the spectroscopic
orbit using data published in Prato et al. (2002) and Mart́ın et al. (2005) as well as our new
radial velocity measurements (Section 2.2).
We fit the radial velocity data using the Binary Star Combined Solution software
(Gudehus 2001), the ORBIT code (Forveille et al. 1999), and our own custom-written IDL
code; all gave the same solution. The best-fit phased radial velocity curve is shown in Figure
1 and the orbital parameters are given in Table 3.
The best-fit period of 19.131±0.003 days is inconsistent with the value of 18.979±0.007
days found by Mart́ın et al. (2005). Examination of the power spectrum of the velocity data
used by Mart́ın et al. shows that the 18.979-day period appears to be an alias of the true
period, caused by beating of the 19.131-day period with two six-year gaps in the radial
velocity data; there is a corresponding alias at 19.3 days. Re-fitting only the data used by
Mart́ın et al., we find that the two periods both correspond to local minima in χ2 space, with
reduced χ2 = 8.8 for P = 19.131 days and reduced χ2 = 10.7 for P = 18.979 days. When
the new radial velocity data are added, the fit for P = 19.131 days improves to reduced
χ2 = 8.1 while that for P = 18.979 days worsens to χ2 = 11.7, as expected if 19.131 days is
the correct period.
We note that we have not added any additional radial velocity measurements of the
secondary, and thus the mass ratio remains more uncertain than the other orbital elements,
resting on the six secondary radial velocities presented by Prato et al. (2002).
UZ Tau E is one of only a handful of pre–main-sequence systems with measured stellar
masses (see Mathieu et al. 2006 for a recent review). Because the total system mass has been
measured (Simon et al. 2000), the spectroscopic orbital parameter M sin3 i can be used to
determined the orbital inclination. This can then be compared with the observed inclination
of the circumbinary disk. While this was done by Simon et al. (2000) and Prato et al.
(2002), we revisit this issue here using our newly-determined orbital parameters for UZ
Tau E. Combined with M = 1.31 ± 0.08 M⊙ (Simon et al. 2000), our orbital parameters
give sin iorbit = 0.81± 0.05, or iorbit = 54
◦ ± 5◦. This is in excellent agreement with the disk
inclinations of 54◦±3◦ and 56◦±2◦ measured from interferometric images of the λ = 1.3 mm
– 7 –
continuum emission and the CO line emission, respectively (Simon et al. 2000). Thus, the
binary orbit and the circumbinary disk are coplanar. Since the disk inclination is measured
at scales of ∼ 100 AU and the binary orbit is only a few tenths of an AU, this co-planarity
apparently extends over the entire disk. Though there are several theoretical studies of
how tilted circumstellar disks interact with a binary system (e.g., Papaloizou & Terquem
1995; Larwood et al. 1996; Bate et al. 2000; Lubow & Ogilvie 2000), we know of no studies
of the timescale for alignment of a circumbinary disk if it is initially tilted relative to the
binary orbit. Studies of the effects of a planetary-mass companion on a tilted external
disk (Lubow & Ogilvie 2001) suggest that such disk tilts do decay over time, however. The
example of UZ Tau E shows that co-planarity over the entire disk exists already by an age
of a few Myr, suggesting that circumbinary disks either form already aligned with the orbit,
or come into alignment very quickly. This is similar to the result of Jensen et al. (2004),
who found that circumstellar disks in young binaries tend to be aligned with each other, and
thus presumably with the binary orbit.
4. Periodic variations
4.1. Photometry
In order to determine whether the system is varying in phase with the binary orbit, or
in any other systematic way, we have searched the photometric data for periodic signals. We
begin by searching for evidence of periodicity, without pre-supposing a particular period.
A Lomb-Scargle periodogram (Scargle 1982) of the I-band data (the band with the largest
number of points and best time coverage; Table 1) is shown in Figure 2. There is a strong peak
at a period of 19.20± 0.03 days. The false-alarm probability (FAP) of this peak is less than
0.001 according to the formulation of Horne & Baliunas (1986). While this FAP calculation
is strictly applicable only to evenly-spaced data, a Monte Carlo bootstrapping method (e.g.,
Stassun et al. 1999) confirms that this period is statistically significant at better than 99.9%
confidence. The period uncertainty reported above, which follows from the formulation of
Kovacs (1981), is probably underestimated as it assumes that the underlying signal is well
described by a single sinusoid.
Though it is slightly off the main peak of the power spectrum, there is significant power
at the binary period of 19.131 days. Periodograms of the B, V , and R data are similar
(Figure 3), showing peaks near the binary period, but with broader peaks and higher false-
alarm probabilities, perhaps due to the more limited time coverage of the data in those
bands.
– 8 –
To refine the period and to better estimate its uncertainty we next performed a phase
dispersion minimization (PDM) analysis (Stellingwerf 1978), which is particularly well-suited
to periodic variability that is highly non-sinusoidal and/or to data with large intrinsic scatter,
both of which apply to the photometry of UZ Tau E. The PDM search of the I-band data
yields a best period of P = 19.17±0.05 d, where the uncertainty was determined empirically
from the 1/e folding scale of the PDM merit function. The same analysis on the V -band
data gives P = 19.15± 0.04 d.
Schwarzenberg-Czerny (1989) argues that a related test, the one-way analysis of vari-
ance, is the most powerful statistic of this kind for detection of periodic signals. Applying
that test to our data yields P = 19.16± 0.03 d for the I-band data, and P = 19.15± 0.05 d
for the V -band data. Following Schwarzenberg-Czerny (1989), the period uncertainty was
determined using a “post-mortem” analysis that measures the 1-σ confidence interval of the
primary periodogram peak, defined by its width at the mean noise power level of the peri-
odogram in the vicinity of this peak. As above, a Monte Carlo permutation analysis of the
light curves confirms that this period is statistically significant at better than 99.9% confi-
dence. Combining these estimates, our best-fit photometric period is P = 19.16 ± 0.04 d,
consistent with the binary orbital period.
Figure 4 shows I-band light curves for all three observing seasons, folded at the binary
orbital period of 19.131 days. As suggested by the periodogram analysis, all show indications
of periodic behavior, with a broad minimum near orbital phase 0.5. The data from the 2004–
2005 season show the smoothest variability, but this is also the season with the smallest
number of data points. Clearly there is significant random variability as well, with scatter
of roughly 0.6 magnitudes at all orbital phases.
Figure 5 shows folded B, V , R, and I lightcurves from the 2003–2004 and 2005–2006
seasons. Broadly speaking, the BV R data show the same behavior as the I-band light curves,
with large-amplitude variability that appears to have both periodic and random components.
We note that the R band includes the Hα line, which may complicate the interpretation of
the light curve.
All four bands show a gradual increase in brightness over the three years of our ob-
servations, with the mean magnitude changing by 0.3 mags at I band from 2003–2004 to
2005–2006. To separate long-term variations from the shorter-term variations of interest
here, a linear trend (with a slope of roughly 0.15 mag / yr at I band) has been fit to each
band. The folded light curves using these de-trended data, and combining all three observing
seasons, are shown in Figure 6.
In addition to our photometric data, previous data on UZ Tau have shown some evidence
– 9 –
of both long-term trends and periodic variations. Bohlin (1923), in one of the first papers
to mention UZ Tau, reported on a major flare and then a four-magnitude overall decline in
brightness from 1921–1923. He also noted that there was a short-period variation with a
period of 10–20 days, which encompasses the period of the variations reported here. Bohlin’s
measurements are for the entire UZ Tau system, but later examination of Bohlin’s plates by
Herbig (1977) showed that it was UZ Tau E that brightened dramatically in 1921.
Variations in color of the system can also give clues about the cause of the variability.
Figure 7 shows the V − I color as a function of I magnitude and orbital phase. The system
shows a behavior commonly seen in T Tauri stars, appearing redder when fainter and bluer
when brighter (Herbst et al. 1994). This behavior is consistent either with periodic changes
in extinction (causing both dimming and reddening when the extinction is higher) or in
accretion (adding additional blue light when the accretion rate is higher).
4.2. Spectroscopy
If the periodic photometric variations are due to changes in accretion rate, one might
expect accompanying variations in Hα emission or spectral veiling, common tracers of accre-
tion. Mart́ın et al. (2005) searched for both of these in UZ Tau E and did not find evidence
of either. With our new spectra, we can revisit the question of Hα variability.
Figure 8 shows the equivalent width of the Hα emission line as a function of binary
orbital phase. The variations do not appear to be strongly correlated with orbital phase.
There is some evidence for lower Hα equivalent widths around phase 0.4–0.8, as seen in the
photometric data, but the data are relatively sparse in that phase range as well.
While a lack of periodic variability would be at odds with the photometric data, we note
that periodicity may not be as obvious in the spectroscopic data, since the two datasets differ
in two important respects. First, the Hα data have much sparser sampling; they span a total
of eight years (1994–2002), but with only a handful of points during a given year. Second,
they do not overlap at all with the photometric data. Given that the photometric data show
both long-term trends and short-term scatter in addition to the periodic variations, and that
T Tauri stars in general are known to show significant random variability at Hα, it may be
difficult to separate random and periodic variations (if any) without a dedicated monitoring
campaign, preferably one that includes simultaneous photometric and spectroscopic mea-
surements. We conclude that while periodic spectroscopic variations similar to those seen
in the photometry are not definitively present in these spectroscopic data, neither are they
ruled out.
– 10 –
5. Discussion
The photometric data (and possibly the spectroscopic data) show periodic variability
at the binary orbital period, suggesting that there is a link between the variability and in-
teractions of the binary with its circumstellar and/or circumbinary material. In this section,
we first argue that the periodic variations are unlikely to be due to stellar rotation, and
then we examine how well the observed behavior matches what is expected from the pulsed
accretion model of Artymowicz & Lubow (1996). Finally, we examine the available data for
other spectroscopic binaries to assess whether or not there is evidence for periodic accretion
as a general phenomenon.
5.1. Could the variations be due to rotation?
Periodic variability is not uncommon in photometric studies of PMS stars. Indeed, dedi-
cated monitoring surveys of rich star-forming regions (e.g., Mandel & Herbst 1991; Attridge & Herbst
1992; Choi & Herbst 1996; Stassun et al. 1999; Rebull 2001; Herbst et al. 2002a) have now
discovered hundreds of PMS stars exhibiting periodic variability, the result of surface-
brightness inhomogeneities (i.e. starspots) that rotate in and out of view with the stellar
rotation period. The periodic variability observed in UZ Tau E is very unlikely to be the
result of such rotationally modulated spot signals, for several reasons. First, the rotation
periods of low-mass PMS stars are nearly always shorter than about 12 days, while the pe-
riod of the variations reported here is 19 days. Among 150 low-mass PMS stars in the Orion
Nebula Cluster, only two stars (∼ 1%) have Prot > 15 d (Herbst & Mundt 2005). Second,
rotationally modulated spot signals are typically sinusoidal, and stable over many cycles or,
in some cases, many years. In contrast, the periodic signal we have found in UZ Tau E
is decidedly non-sinusoidal, with considerable scatter; the “bright” state has a duty cycle
of ∼ 60%. Thus, it is either intrinsically non-sinusoidal, or shows substantial phase shift-
ing from one cycle to the next; neither of these is consistent with rotationally-modulated
variability.
The rotation period distributions discussed above are presumably dominated by single
stars or member of wide binaries, while tidal interactions between the stars in a close binary
system can synchronize the orbital and rotational periods. However, in the case of eccentric
systems like UZ Tau E, pseudo-synchronization (in which the stellar angular velocity is
synchronized with the orbital angular velocity at periastron) occurs instead, since the tidal
interactions are strongest around periastron (Hut 1981). The predicted pseudo-synchronous
rotation period for UZ Tau E, using the weak friction formulation of Hut (1981) and the
orbital parameters in Table 3, is 11.4 ± 1.2 d, inconsistent with the observed variability
– 11 –
period.
We can also estimate the rotation period of the UZ Tau E primary directly if three
quantities are known: the inclination irot of the star’s rotation axis, the star’s projected
rotational velocity v sin irot, and the stellar radius. Of these, the inclination is typically
impossible to measure, except under special circumstances.
As noted in Section 3, the dynamical mass measurement of UZ Tau E allows us to
determine the binary orbital inclination iorbit. Based on studies of other binary systems, it
is reasonable to assume that this inclination is the same as that of the stellar rotation axis,
irot. The most detailed study comparing the orientations of these axes in binary systems
is that of Hale (1994). Considering spectral types of F5–K5, he finds that binaries with
separations less than 30–40 AU tend to exhibit co-planarity between rotational equators and
orbital planes, while wider binaries have random orientations. Using a similar method, Weis
(1974) found a tendency for the stellar rotational equators to align with the binary orbit
among primaries in F star binaries. Interestingly, Weis (1974) did not find a tendency toward
co-planarity between rotational and orbital planes among A stars, suggesting that caution
is necessary when comparing stars of different masses. Similarly, Guthrie (1985) found no
correlation between orbital inclination and v sin i among 23 A2–A9 binaries with semi-major
axes of 10–70 AU. The low mass and short period of UZ Tau E suggest, however, that the
conclusions of Hale (1994) are most applicable here.
Prato et al. (2002) find L = 0.63+0.19
−0.17
L⊙ and Teff = 3700 ± 150 K for the primary in
UZ Tau E. Combining these values yields R = 1.9 ± 0.2 R⊙. Hartmann & Stauffer (1989)
find v sin i = 15.9±4.0 km s−1 for UZ Tau E using optical spectra, consistent with the value
v sin i = 16 ± 2 km s−1, which we measure from our new spectra and adopt here. Since
absorption lines of the secondary of UZ Tau E have only been seen in near-infrared spectra
and are not evident in any of our optical spectra, we take this to be the projected rotation
velocity of the primary. Combining these measurements with sin iorbit = 0.81±0.05 (Section
3), and assuming iorbit = irot, we find Prot = 4.9± 0.8 d. If iorbit 6= irot, we find Prot ≤ 6± 1 d
since sin i ≤ 1. Thus, uncertainty on the inclination cannot reconcile the photometric period
with the inferred rotation period.
The most uncertain remaining quantity is v sin i, but since Hartmann & Stauffer (1989)
measured v sin i from 11 different spectra of UZ Tau E, with self-consistent results from two
different parts of the spectrum (including spectra near λ = 5200 Å) and consistency with our
new v sin i measurement, it is unlikely that line broadening from photospheric lines of the
faint, red secondary could lead to an overestimate of v sin i by a factor of four. Similarly, given
the uncertainties on L and Teff , it is difficult to see how the radius could be underestimated
by a factor of four. Thus, we conclude that the observed periodic variations are unlikely to
– 12 –
be due to stellar rotation.
5.2. Evidence for pulsed accretion
We have shown above that UZ Tau E exhibits periodic photometric variations that have
the same period as the binary orbit, and that these variations are unlikely to be caused by
stellar rotation. Here, we examine the predictions made by the pulsed accretion model of
AL96 and compare them to our observations.
5.2.1. What are the predictions?
Broadly speaking, Artymowicz & Lubow (1996) predict that a binary with an eccentric
orbit and a circumbinary disk will have an accretion flow from the circumbinary disk, and
thus onto the circumstellar disks or stellar surfaces, that varies periodically at the binary
orbital period.
The exact behavior of the accretion rate with orbital phase depends on the binary orbital
parameters. AL96 show the results of two simulations, one for mass ratio M2/M1 = 0.43 and
eccentricity e = 0.1, and another for M2/M1 = 0.79 and e = 0.5. The former shows accretion
that varies relatively smoothly over the orbital period, while the latter is strongly peaked
at periastron. As noted by AL96, the exact timing of the accretion variability depends
on the orbital parameters, most strongly on e. Some previous observational studies of T
Tauri spectroscopic binaries have focused specifically on looking for enhanced accretion near
periastron; however, we note here that the actual prediction of the model is more general
than that, and that the peak accretion rate need not come near periastron.
5.2.2. How well do the data match the predictions?
First, we note that our observations match the general predictions of AL96 quite well,
in that there are indeed periodic photometric variations at the binary orbital period, which
are readily interpretable as a variable accretion rate. The comparison with the spectroscopic
data is more ambiguous; if more intensive monitoring of the Hα line in UZ Tau E were to
show that there are no orbit-modulated Hα variations, it would present a problem for the
model.
For a more specific comparison with our data, Figure 9 shows the variations of accretion
– 13 –
with orbital phase predicted by AL96 for a binary with M2/M1 = 0.43, e = 0.1. UZ Tau E
has a more extreme mass ratio (M2/M1 = 0.30) and larger eccentricity (e = 0.33) than this,
but these parameters are closer to those of UZ Tau E than those of the other simulation in
AL96. AL96 do note that the timing of the maxima of the accretion depend largely on e
rather than M2/M1. Since e for UZ Tau E is intermediate between the two models calculated
by AL96, we might then expect the maximum accretion to come between the phase of ∼ 0.75
they calculate for the low-e case and the phase of ∼ 1 for the high-e case.
For comparison with our data, we have taken the logarithm of the variations of accretion
rate predicted by AL96 to shift them onto a “magnitude-like” scale, and added an arbitrary
offset and scale factor to match the mean of the data and amplitude of the variations. The
phase of the minimum predicted by this simulation does not match our data well; when the
model is given a shift of +0.2 in orbital phase, there is better agreement between the model
predictions and the data. This scaling and shifting to match the data is obviously ad hoc,
but it allows us to compare the phase width of the observed variations, which appear to
match the predictions relatively well. In addition, this shifted position of the maximum is
indeed between the two cases calculated by AL96, as expected if eccentricity is the dominant
factor in determining the timing of maximum accretion.
5.3. Evidence for periodic accretion in other T Tauri binaries
The discussion and data above show that looking for evidence of periodic accretion can
be complicated, with other sources of variability perhaps being important and masking the
effect in small datasets, and with the exact behavior expected to be a function of the specific
binary orbital parameters. That said, is evidence for pulsed accretion seen in other young
binary systems? In Table 4 we present characteristics of young binaries with periods of
less than one year and evidence of circumbinary material, in order of increasing eccentricity.
Below, we examine the observational data for some of these systems, attempting to relate
them to what we see in UZ Tau and exploring similarities and differences. Unfortunately,
the small number of systems and their somewhat heterogeneous properties means that it is
difficult to generalize, so we offer these comments in the spirit of attempting to pull together
the existing data, rather than arguing one way or the other for the validity of the AL96
model for the sample as a whole.
– 14 –
5.3.1. DQ Tau
DQTau was the first system to be scrutinized for evidence of pulsed accretion. Mathieu et al.
(1997) showed that the photometric variations are modulated at the binary orbital period,
and Basri et al. (1997) showed that the Hα line and spectral veiling are as well. Fortuitously,
the mass ratio and eccentricity of DQ Tau are quite similar to those of the high-eccentricity
case modeled by AL96, allowing for specific comparison with the theory. The timing and
phase width of the photometric and spectroscopic variations match the predictions well,
being sharply peaked near periastron. However, the DQ Tau observations did show consid-
erable orbit-to-orbit variation, with the periastron brightening being seen roughly 65% of
the time. This is reminiscent of the large scatter that we see in the UZ Tau light curves;
clearly the periodic accretion process is not exactly repeatable from orbit to the next, nor is
it the only source of variability.
5.3.2. AK Sco
The orbital eccentricity and binary mass ratio of AK Sco are quite similar to those of DQ
Tau, and indeed, simulations by Günther & Kley (2002) for a binary with AK Sco’s orbital
parameters predict pulsed accretion. Thus, it comes as some surprise that the system does
not show periodic photometric variability, despite extensive monitoring (Alencar et al. 2003).
The overall variability is large (up to 1.5 mags in y), but apparently random. There are
periodic variations in the Balmer lines, but they are not sharply peaked around periastron.
Examining Table 4, we note two properties of AK Sco that are quite different from those
of DQ Tau or UZ Tau. First, AK Sco is considerably hotter and more luminous. Thus,
accretion variations of the same luminosity as those occurring in UZ Tau and DQ Tau would
result in substantially smaller magnitude changes, which could be swamped by the large
random variability. Second, AK Sco has considerably lower millimeter flux than either of
the other two systems. If the systems are fit with similar disk models (in which the disk
is assumed to be optically thin at millimeter wavelengths in its outer regions), AK Sco’s
disk mass is an order of magnitude smaller than that of DQ Tau or UZ Tau E (Jensen et al.
1996a; Jensen & Mathieu 1997; Mathieu et al. 1997). Alencar et al. (2003) fit AK Sco with
an optically-thick disk model that has a comparable mass to the disk models fit to DQ Tau
and UZ Tau E. However, such disk models have not been fit to DQ Tau or UZ Tau E, and
would presumably result in even larger disk masses for those systems. In a direct comparison
of λ = 1.1 mm flux, DQ Tau and UZ Tau E are 3 and 5 times brighter than AK Sco at
roughly the same distance, presumably reflecting larger disk masses. It is possible that a
somewhat lower-mass disk has different dynamics, and that the accretion flow in the AK Sco
– 15 –
system is fundamentally different than that in the other systems with more massive disks.
5.3.3. GW Ori
This system has a near-circular orbit, and thus would not be expected to show pulsed
accretion under the model set forth by AL96. However, Stempels & Gahm (2004) quote
Artymowicz, private communication, as saying that pulsed accretion is possible for systems
with circular orbits as well, and indeed D’Angelo et al. (2006) show that this occurs for giant
planets embedded in disks. Thus, pulsed accretion appears to be possible for at least some
circular-orbit systems and thus may be for GW Ori as well, though the larger gap cleared by
a stellar companion (Artymowicz & Lubow 1994) will clear some of the disk resonances that
might contribute to disk eccentricity growth in a system with a planetary-mass companion.
Like AK Sco, GW Ori is very luminous, and shows significant random variability, though
no obvious periodic variability. It does have a much more massive disk than AK Sco, however,
and indeed than any of the systems considered here. Because of its much larger semimajor
axis, and to some extent its circular orbit, GW Ori is much more likely to have significant
circumstellar disks, as the stars do not approach each other very closely at periastron. Thus,
material flowing from the circumbinary disk may merge with the circumstellar disks and then
accrete more gradually onto the stars, rather than falling directly on (or near) the stellar
surfaces as is expected to happen in the shorter-period systems. If the infalling material does
not shock strongly as it merges with the circumstellar disk, and if any density enhancements
are smoothed out somewhat by the time the material reaches the stellar surface, then any
photometric signature of the periodic infall would be weakened. We note that UZ Tau E
likely has circumstellar disks as well (Jensen et al. 1996a), so a similar effect could be at work
in reducing the amplitude of the periodic variability relative to the stochastic variability.
5.3.4. V4046 Sgr
Like GW Ori, V4046 Sgr has a nearly circular orbit. However, V4046 Sgr has shown
periodic photometric variations at the binary orbital period (Quast et al. 2000; Mekkaden
2000). These variations persist over several years and are relatively sinusoidal (Walter,
unpublished data, 2003–2005). Unlike the other binaries discussed here, in this case stellar
rotation is a plausible explanation of the observed variations. It is common for stellar
rotational periods to become synchronized with the binary orbital period, particularly for
short-period binaries like V4046 Sgr. Given the short period (resulting in stronger tidal
– 16 –
interactions and a shorter synchronization time scale) and the somewhat older age of this
system (∼ 10 Myr), synchronization is plausible, and indeed is supported by detailed analysis
of the system (Stempels & Gahm 2004). However, rotation does not explain the periodic
Balmer line variations observed, which Stempels & Gahm (2004) attribute to accumulations
of gas co-rotating with the binary orbit.
5.3.5. ROXs 42 and ROXs 43B
These two spectroscopic binaries are both weak-lined T Tauri stars (Bouvier & Appenzeller
1992; Walter et al. 1994), indicating less-active accretion than some of the other systems dis-
cussed here. Neither has been detected at millimeter wavelengths, yielding only an upper
limit on the disk masses (Skinner et al. 1991; Jensen et al. 1996b). Both systems show
mid-infrared excesses, indicating the presence of circumbinary material, and a lack of near-
infrared excess, which can be modeled as a cleared central region in the disk (Jensen & Mathieu
1997). The fact that both are higher-order multiple systems complicates matters; ROXs 42
(NTTS 162814−2427) is a triple system with a separation of 0.′′15 (Lee 1992; Ghez et al.
1993), while ROXs 43B (NTTS 162819−2423S) has a wide companion at 4.′′8 which is itself a
close binary system (Walter et al. 1994; Simon et al. 1995b). Since the evidence for the pres-
ence of a substantial disk rests on the low-spatial-resolution IRAS detections, it is possible
that the excess is associated with the wider companions rather than arising from circumbi-
nary disks around the spectroscopic binaries. In any case, the lack of millimeter detections
indicates that there is less disk mass in these two systems than in the others discussed here.
Neither system has been intensively monitored over timespans that would be necessary to
detect periodic photometric variations at the relatively long orbital periods. ROXs 42 shows
evidence for some semi-regular variations over roughly 1.5 orbital periods (Zakirov et al.
1993), while the combined light of the ROX 43 system shows only a 0.1-magnitude varia-
tion, with evidence of a 1.5-day or 3-day periodicity, presumably attributable to rotation of
one or more of the stars (Shevchenko & Herbst 1998).
5.3.6. KH 15D
The unusual pre–main-sequence system KH 15D (V582 Mon) is a spectroscopic bi-
nary that undergoes deep (∆I ∼ 3.5 mag) eclipses, thought to arise due to occultation
from a circumbinary disk (Hamilton et al. 2001; Herbst et al. 2002b; Hamilton et al. 2005;
Winn et al. 2006 and references therein). While the system has an eccentricity and mass
ratio that would suggest that pulsed accretion might be present, the photometric variations
– 17 –
at the binary orbital period are dominated by the deep eclipses. Furthermore, the depth and
detailed shape of these eclipses are evolving with time (Winn et al. 2003; Johnson & Winn
2004; Maffei et al. 2005; Johnson et al. 2005; Winn et al. 2006), making it very difficult to
determine whether there might currently be an additional, smaller-amplitude component
with the same period that is related to accretion rather than occultation. Winn et al. (2003)
showed that the current deep eclipses did not occur during the first half of the twentieth
century, raising the possibility of searching for evidence of accretion-related variability at
earlier epochs. Their limit of one mag on the variability during that time does not preclude
accretion-related variations like those seen in UZ Tau E. The ∼ 0.9-mag periodic varia-
tions seen from the 1960’s through the 1980’s (Johnson & Winn 2004; Maffei et al. 2005;
Johnson et al. 2005), however, are relatively smooth and are well-fit by the eclipse model
(Winn et al. 2006), placing a limit on how much any accretion-related component was con-
tributing to the variability during that time. Since the inferred mass ratio and eccentricity
for KH 15D are similar to those of DQ Tau (Table 4), we might expect accretion-related vari-
ability to be strongly peaked around periastron, which is also when the current deep eclipses
occur. This might help explain several anomalously bright points seen during eclipses in the
late 1990’s that are not well-fit by the model of Winn et al. (2006).
The precessing circumbinary disk occultation model of Winn et al. (2006) is quite suc-
cessful in reproducing the shape and ongoing evolution of the light curve, and we do not
suggest that accretion explains most of the photometric variations. We note, however, the
possibility that such an additional component might be sporadically present (with the same
period) and that, if it is, this could complicate the modeling of the historical evolution of
the light curve, especially during earlier, more-sparsely-sampled epochs.
6. Conclusions
We have shown that the pre–main-sequence binary UZ Tau E shows clear photometric
variability at the binary orbital period of 19.13 days. This variability is consistent with a
model in which material in the circumbinary disk is periodically perturbed by the binary in
its eccentric orbit and falls from the outer disk, across the cleared central gap and onto the
stars or their circumstellar disks. There is significant scatter in the light curves, indicating
that this “pulsed accretion” may not occur during every binary orbit. Hα equivalent widths
show some suggestion of periodic variability, but it is not definitive.
The apparently intermittent behavior of the accretion, and the presence of other, random
sources of variability, suggest that searches for this sort of accretion signature require well-
sampled datasets with long time baselines in order to detect any periodic component. In
– 18 –
particular, simultaneous photometric and spectroscopic monitoring of UZ Tau E in the future
will help determine whether the Hα variations show a periodic component, as the photometric
variations do.
The good overall agreement between theory and observations suggests that resonant
interactions between stars (and, by extension, planets) and disks are indeed important in
determining disk structure and dynamics, while the random component of the observed
behavior shows that there is still work to be done in understanding the full complexity of
these interactions.
We gratefully acknowledge the support of the National Science Foundation through
grant AST-0307830. We thank the referee, Steve Lubow, for useful comments that improved
this paper. We are grateful to Michael Meyer, David Cohen, and Larry Marschall for useful
discussions; to Marcos Huerta and Pat Hartigan for use of their spectra of UZ Tau E; to
Matthew Richardson for assistance with data reduction; to Peter Collings for translating
early papers on UZ Tau from German to English; and to Thierry Forveille for use of his
ORBIT code. MS and FW are grateful for Stony Brook University’s partial support of
their participation in the SMARTS consortium. This research has made use of the SIMBAD
database, operated at CDS, Strasbourg, France, and of NASA’s Astrophysics Data System.
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Table 1. Observations of UZ Tau E
Telescope Exp. time (s) Filter(s) Season # of nights Timespan in days
SMARTS (1.3m) 5 BV RI 2003–2004 63 170
2005–2006 6 42
30 V RI 2005–2006 9 33
VVO (0.6m) 60 I 2004–2005 16 126
2005–2006 20 128
– 24 –
Table 2. Radial Velocities and Hα EW
Julian Date vhelio (km s
−1) Hα EW (Å)a
2450416.82 · · · 88.4
2450783.93 · · · 42
2450783.94 · · · 45.1
2450784.96 · · · 38
2450785.11 · · · 57
2450835.65 · · · 54:
2451060.99 · · · 45.6
2451061.00 · · · 39.8
2451077.14 · · · 63.8
2451120.92 · · · 101
2451137.94 · · · 51.3
2451138.88 · · · 62.6
2451162.86 · · · 69.3
2451163.82 · · · 57.8
2451164.79 · · · 57.5
2451165.78 · · · 58.3
2451166.79 · · · 61.7
2451169.84 · · · 48.9
2451504.73 · · · 35
2451507.68 · · · 25:
2451508.67 · · · 49.1
2451509.71 · · · 57
2451510.66 · · · 71
2451517.98 · · · 75
2451523.68 · · · 42.7
2451524.74 · · · 49.3
2451525.76 · · · 40.8
2451527.68 · · · 46.6
2451528.68 · · · 58.9
2451529.66 · · · 61.8
2451530.60 · · · 69.1
– 25 –
Table 2—Continued
Julian Date vhelio (km s
−1) Hα EW (Å)a
2452280.60 −4.3 ± 2.1 44
2452281.74 −3.1 ± 1.5 51
2452282.76 −5.8 ± 1.8 81
2452283.66 −4.9 ± 4.6 78
2452284.71 · · · b 77
2452286.66 6.2 ± 2.6 90
2452287.65 17.1 ± 1.7 87
2452288.68 27.2 ± 2.1 59
2452289.69 29.6 ± 3.0 65
2452290.72 37.8 ± 7.0 67
2452291.67 28.5 ± 2.1 58
2452292.64 29.1 ± 3.2 50
2452293.67 22.0 ± 3.6 45
2452579.11 · · · 50
aPositive values denote emission.
bThe spectrum on this date was too noisy
to allow measurement of an accurate radial
velocity.
– 26 –
Table 3. Binary orbital parameters for UZ Tau E
Period (days) 19.131± 0.003
e 0.33± 0.04
JD of periastron 2451328.3± 0.5
ω 239◦ ± 9◦
a sin i (AU) 0.124± 0.003
γ (km s−1) 13.9± 0.7
K1 (km s
−1) 17.3± 1.4
K2 (km s
−1) 57.4± 4.7
M sin3 i (M⊙) 0.69± 0.13
M2/M1 0.30± 0.03
Table 4. CTTS spectroscopic binaries
Binary Period e M2/M1 Spectral L Disk Mass
a Photometric Balmer line References
System (days) Type (L⊙) (M⊙) periodicity? periodicity?
V4046 Sgr 2.421 ≤ 0.01 0.94 K5 0.82 0.0085 Yes (∆B≈0.1) Yes 1, 2, 3
GW Ori 241.9 0.04 SB1 G0 26 0.3 ? (∆V≈0.7) ? 4, 5
UZ Tau E 19.131 0.33 0.30 M1 0.91 0.063 Yes (∆I≈0.8) Maybe 6, 7, 8
ROXs 43B 89.1 0.41 SB1 G0 0.4 < 0.00037 ? (∆V=0.1) ? 1, 9, 10, 11
AK Sco 13.609 0.47 0.99 F5 8.40 0.002 No (∆y≈1.5) Yes 1, 12, 13
ROXs 42 35.95 0.48 0.92 K4 0.4 < 0.00025 ? (∆V=0.4) ? 1, 10, 11, 14, 15
DQ Tau 15.804 0.56 0.97 K7–M1 0.95 0.020 Yes (∆V≈0.5) Yes 16, 17
KH 15D 48.38 0.57–0.65 0.83b K7c 0.4c · · · Eclipse (∆I≈3.5) ? 18, 19, 20
aDisk mass estimates were made assuming that the disk is at least partially optically thin at mm wavelengths. An optically-
thick model for AK Sco (Alencar et al. 2003) yields a disk mass of 0.02 M⊙.
bDerived from the stellar luminosity ratio that best fits the eclipse data (Winn et al. 2006).
cProperties of the secondary star, since the primary is never visible.
References. — 1. Jensen & Mathieu (1997). 2. Quast et al. (2000). 3. Mekkaden (2000). 4. Mathieu et al. (1991). 5.
Mathieu et al. (1995). 6. This work. 7. Prato et al. (2002). 8. Jensen et al. (1996a). 9. Shevchenko & Herbst (1998).
10. Manset & Bastien (2003). 11. Bouvier & Appenzeller (1992). 12. Alencar et al. (2003). 13. Manset et al. (2005). 14.
Lee (1992) 15. Walter et al. (1994). 16. Mathieu et al. (1997). 17. Basri et al. (1997). 18. Hamilton et al. (2001). 19.
Hamilton et al. (2005). 20. Winn et al. (2006).
– 28 –
Fig. 1.— The best-fit spectroscopic orbit for UZ Tau E. Crosses show velocities of the pri-
mary; those enclosed in boxes (in red in the on-line edition) show new radial velocity mea-
surements presented here. Open diamonds show velocities of the secondary from Prato et al.
(2002).
– 29 –
Fig. 2.— The Lomb-Scargle periodogram for the I-band data. The periodogram peaks at
a period of 19.20 days, with a false-alarm probability of 0.001. The dashed line shows the
binary orbital period of 19.131 days. The smaller peaks visible flanking the main peak (lower
panel) are near the alias periods expected for beat periods between one-year and two-year
periods (caused by the seasonal gaps in the data) and the binary period.
– 30 –
Fig. 3.— The Lomb-Scargle periodograms for the B, V , R, and I-band data. All show
significant power near the binary orbital period.
– 31 –
Fig. 4.— The I-band magnitude for UZ Tau E folded at the binary orbital period of 19.131
days and plotted against the binary orbital phase. Top to bottom, data from 2003–2004,
2004–2005, and 2005–2006.
– 32 –
Fig. 5.— The BV RI magnitudes for UZ Tau E folded at the binary orbital period and
plotted against the binary orbital phase. Left, 2003–2004; right, 2005-2006. The open circles
in the lower right plot show the VVO I-band data, which do not have corresponding B, V,
and R data.
– 33 –
Fig. 6.— The BV RI magnitudes for UZ Tau E folded at the binary orbital period and
plotted against the binary orbital phase, after removing a long-term linear trend from each
band.
– 34 –
Fig. 7.— V − I color vs. I magnitude and vs. orbital phase for 2003–2004 (closed circles)
and 2005–2006 (open circles). The system is redder when fainter and bluer when brighter,
the expected behavior either for changes in extinction or for brightening due to increased
accretion.
– 35 –
Fig. 8.— Top: Equivalent width of the Hα line as a function of binary orbital phase. Squares
are our new measurements; triangles are from Mart́ın et al. (2005). Bottom: For comparison,
the phased I-band data. There is some suggestion of reduced Hα equivalent width at phases
of 0.4–0.8 as seen in the photometric data, but the data are too sparse there for there to be
clear evidence for periodic variability of the Hα emission.
– 36 –
Fig. 9.— Left: The theoretical predictions of Artymowicz & Lubow (1996) for the depen-
dence of accretion rate on binary orbital phase in a binary with mass ratio M2/M1 = 0.43,
e = 0.1. The top curve shows the total accretion, while the lower curves show accretion onto
the secondary (higher dark curve) and primary (lower dark curve). Right: The same total
accretion curve, but placed onto a logarithmic scale and shifted vertically for comparison
with the phased I-band data. The model here has been given an ad hoc shift of +0.2 in
phase, roughly what is expected given the binary eccentricity (see Section 5.2.2).
	Introduction
	Observations
	Photometry
	Spectroscopy
	Binary orbital parameters
	Periodic variations
	Photometry
	Spectroscopy
	Discussion
	Could the variations be due to rotation?
	Evidence for pulsed accretion
	What are the predictions?
	How well do the data match the predictions?
	Evidence for periodic accretion in other T Tauri binaries
	DQ Tau
	AK Sco
	GW Ori
	V4046 Sgr
	ROXs 42 and ROXs 43B
	KH 15D
	Conclusions
