ground-based detection of terrestrial extrasolar planets by photometry: the case for cm draconis

GROUND-BASED DETECTION OF TERRESTRIAL EXTRASOLAR

PLANETS BY PHOTOMETRY: THE CASE FOR CM DRACONIS

JEAN SCHNEIDER CNRS-Observatoire de Paris, 92195 Meudon, Fr ance

and

LAURANCE R. DOYLE SETI Institute, 245-3, NASA Ames Research Center, Moffett Field, CA 94035

(Received June 5, 1995)

Abstract. The detection ofextrasolar planets by measuring a photometric drop in the stellar brightness due to a planetary transit can be statistically improved by observing eclipsing binary systems and photometrically improved by observing small component systems. In particular the system CM Draconis, with two dM4 components, would allow the detection of extrasolar planets in the size range of Earth-to-Neptune requiring a ground-based photometric precision of about 0.08% to l.1% (photometric precision of about 0.3% is routinely achievable with 1-meter class telescopes at the magnitude of CM Draconis, 11.07 in R-filter). In addition, the transit of extrasolar planets in a binary star system provides a unique, quasi-periodic signal that can be cross-correlated with the observational data to detect sub-noise signals. We examine the importance of making such observations to an understanding of the formation and evolution ofterrestrial-type planets in main-sequence star systems. Terrestrial planets could have formed with substancially shorter periods in this lower luminosity system, for example, and might be expected to have accreted essentially in the binary orbital plane (however, non-coplanar planets may also eventually be detectable due to precession). We also report on a network of medium-sized telescopes at varying longitudes that have been organized to provide such constraints on terrestrial-planet formation processes and discuss the extention of near-term observations to other possible binary systems, as well. Finally, we discuss a more speculative, future observation that could be pefformed on the CM Draconis system that would be of exobiological as well as astrophysical interest.

1. Introduction

The detection of extrasolar planets remains one of the most important challenges in astronomy today. There are two quite heterogeneous motivations for making such detections. First of all we would like to investigate the standard astrophysical questions such as whether or not planetary systems are frequent in the galaxy, what their characteristrics are (how they depend on the central star, etc.), and what characterizes their formation processes. Secondly, there is a more qualita- tive aspect: Are there planets in what is callëd the 'habitable zone' around stars where the temperatures and luminosities allow for the development of biochemical activity? Habitable planets are most likely to be the inner, smaller planets. These are unfortunately, - at least for the next couple of decades, - going to likely be impossible to detect with present standard imaging, astrometric, or spectroscopic (Schneider 1995 for a review). The present standard planetary detection scenario

Earth, Moon and Planets 71: 153-173, 1995. @ 1995 Kluwer Academic Publishers. Printed in the Netherlands.

154 JEAN SCHNEIDER AND LAURANCE R. DOYLE

is to therefore start with the detection of giant, exterior planets ('Jupiters') leaving the detection of inner, Earth-like planets for the next generation. However, even one such detection would provide essential information on the formation of "solar" systems and could supply a strong calibration for the uniqueness (of not) of planetary habitats in the galaxy. Thus, aside from the atypical case of pulsars (Wolszczan 1994) and candidiates around Proxima Centauri (Benedict et al. 1995) and/3 Pictoris (Lecavelier des Etangs et al. 1995), no other extrasolar planets have been found. This is because virtually all of the ground-based detection methods are presently at the limit of their performances (Schneider 1995) requiring special techniques to be developed, - (interferometry with adaptive optics, for example) or orbital observations to be made. Outer planets may be expected to be gaseous (and likely giant; Safronov and Ruskol 1994) as they taust be located at temperatures compatible with the avoidance of hydrogen evaporation (~200 °K). The question as to whether giant planets are frequent or not has recently emerged, as none has been found in a spectroscopic survey of 21 nearby solar-type stars (in the range 1.5 to 3.0 Jupiter masses; Walker et al. 1995), and their formation may be inhibited by early stellar winds (Zuckerman et al. 1995). This enforces the importance of the search for inner, terrestrial planets. Inner planets may be expected to be more dense although not necessarily as small in mass as the terrestrial planets of the solar system (the inner pulsar planet around PSR 1257 + 12 at 0.3 AU has a mass of 3M e, for example). Only inner planets will, in general, be compatible with the possibility of surface liquid water (temperatures of ~300 °K) and therefore relevant to estimates of exobiological habitats.

There are various methods presently being investigated for the detection of extrasolar planets (Schneider 1995). For example, the astrometric detection method relies on detecting the angular displacement of a star around the star-planet barycenter, and ground-based techniques may only possibly detect jovian-sized outer planets as the moment arm must be large. Direct imaging methods are also constrained to large outer planets since the brightness ratio of the star to planet is 1arge (six to ten orders of magnitude from IR to visible, respectively) and therefore the resolution required generally consigns this method to orbital observations. The radial velocity techniques, which measure periodic Doppler shifts in the star indicative of displacement about the star-planet barycenter in the line of sight, is optimal for large, inner planets as lhe velocity displacement of the star would therefore be maximal. In practice, the ground- based application of this method may eventually be able to detect Neptune-mass planets in inner orbits. Ground-based photometric measurements fall into two catagories. One is the measurement of momentary brightness surges due to gravitational lens amplification of stellar signals caused by the passage of a planetary body along the star-to-Earth null geodesic. While sufficient applification is available, this technique has the difficulty that observations can not be confirmed, as the gravitational lense event will not be expected to repeat. Nevertheless it could provide an interesting general survey of the distribution of the lower end of the mass function (planets) in the galaxy.

TERRESTRIAL EXTRASOLAR PLANETS AND CM DRACONIS • 5 5

TABLE I

Characteristics of terrestrial and jovian planetary transits

Characteristic Terrestrial planet Jovian planet

Luminocity drop 10 .4 10-2

Transit duration (hours) 13 28 Transit probability 0.5% 0.1%

The second photometric technique and last of the major detection methods (several other variations on those above are reviewed in Schneider 1995) is the occultation technique which relies upon the detection of differential brighness changes in a star due to the transit of an extrasolar planet(s) across the line of sight. While this method is expected to require extremely precise photometry (the photometric drop is proportional to the square of the planet-to-star radii), it is the statistical improbability of line-of-sight transits that have, in general, constrained the application of this method. In this work we outline ways in which both the statistical as well as the photometric limitations to the application of this method may be overcome to the point where the detection of terrestrial-sized planets with medium-sized telescopes from the ground may be achieved (for the system CM Draconis, for example, and possibly several other systems, as well).

2. The Photometric Method and CM Draconis

2.1. THE OCCULTATION METHOD IN GENERAL

The detection of extrasolar planets by observing a differential photometric drop in the brighness of a star caused by the transit of an extrasolar planet across the line of sight was first proposed by Struve (1952), developed by Rosenblatt (1971) and quantified by Borucki and Summers (1984). The luminosity drop during a planetary transit is given by the ratio of the planet area to the star area, (Re/R~)2. The duration of the transit is given by (P/~r)(R./a) where P and a are the period and the radius of the planet's orbit. The probability of a planetary transit can be estimated by taking the ratio of the area of detectablity to that of the area on the celestial sphere, that is, the product of the required line-of-sight inclination of the planetary orbital plane, 2~r × o × 2R, to 4~ra 2, which gives an estimated random alignment probability of R./a. In Table I we give the main characteristics of transits for a terrestrial and a jovian transit for a solar-type star.

Depending on the period of the planet to be detected, then, one would generally have to continously observe thousands of stars to obtain even one transit in the line of sight. (However, even planets with inclined orbital planes can be expected to


also precess across the line of sight, and this somewhat improves the statistics of detection as we mention below; see simulations in Schneider 1994b.) Constrains on the planetary orbital plane - expecting it to form very near the stellar equatorial plane - can also be placed by determining the stellar rotational inclination, providing some helpful indication of the planetary orbital inclination to the line-of-sight (Doyle et al. 1984; Doyle 1988; Hale and Doyle 1994). However, the sine function is more than an order of magnitude more sensitive around 0 ° inclination (stellar rotation axis near pole-on) than around 90 ° (edge-on) where the photometric method requires it.

In addition, - taking the Earth in transit around the Sun as an example, - the differential photometric drop in brightness due to such a transit would be about one in 10 -4, while the scintillation limit on ground-based photometeric precision is usually taken to be about 0.1% (Young et al. 1991; but Gilliland et al. 1994 indicate that a limit of 0.025% with 4-meter telescopes may be achievable). The major limitations come, however, from the atmospheric transparency fluctuations which have a 1 / f power spectrum independent of the telescope size and have, at the frequency of f = 0.05 min-I characteristic of the luminosity drop due to a planetary transit, an amplitude of about 1-2 × 10 -4 (Harvey 1987). Both the statistical probability of transit and the photometric limits, however, may fortunately be significantly improved for a number of special systems (of which CM Draconis may be the best example), namely eclipsing binary pairs of dM stars.

2.2. OCCULTATIONS BY A PLANET IN AN ECLIPSING BINARY SYSTEM

We can start to address these concems by beginning with already known "edge-on" systems. Eclipsing binaries, in general, have two advantages: First the probability that the planetary orbital plane is correctly oriented is close to 1. Also, even non- planar planetary planes will precess with a period:

3 ' (1)

where .s is the binary separation (Schneider 1994a). This precession, therefore, necessarily forces at least one planetary transit every half precession period. If it is a close binary, the possible planetary precession periods will be nonnegligible (Section 3 below) so that the probabilty that a transit finally occurs is exactly 100% (Schneider 1994a). A second point is that since the star is double, there are likley two transits per transit event (Schneider and Chevreton 1990). Furthermore, the two components of the binary have a (mean) transverse velocity of (GMx/2s) 1/2 (for an equal mass binary like CM Draconis) which is larger than the transverse velocity (GM, /a ) 1/2 of the planet. Therefore, the duration of a single occultation can take any value between a maximum of:

Rs

(aMx/a)~/2(,/U~- 1)

TERRESTRIAL EXTRASOLAR PLANETS AND CM DRACONIS 157

(for the star's motion parallel to the planets motion), and a minimum of:

R,

( G M , / a ) l / 2 ( x / ä 7 ~ + 1)

depending on the relative phase between the binary and the planetary motions. Thus the number of transits per transit event (planet passing in from of the two stars) depends on the configuration of the system (Schneider and Chevreton 1990, Brandmeier and Doyle t995, Jenkins et al. 1995). In addition, the two major advantages of dM stars (and possbily eclipsing binaries with white dwarf companions) are that since these stars am small in radius, the depth of a transit for a given planetary radius is larger. We also point out below that the M dwarf stars, being cooler, a solid (terrestrial-type) planet may form in orbit much closer to the stellar binary pair (Schneider et al. 1990).

Also, although the precession of the planetary orbital plane already mentioned forces an accultation to occur even if this plane makes a non-zero angle with the binary orbital plane, one can be reasonably sure that these two planes are, in fact, parallel. This is due to the formation scenario of a planet around a binary. Suppose, indeed, that the protoplanetary disk has some thickness and that it is initiatly not parallel to the binary plane. Then the pmcession rate as given in Equation (1) will be different for dust rings with different radii, leading to a disk symmetrical with respect to the binary plane. Furthermore, this precession will induce dissipative collisions between dust grains giving a dissipative force perpendicular to the disk. The component of the dust velocity perpendicular to the disk will be damped, finally leaving the disk totally flat, a scenario similar to the flattening of Saturn's rings. This process will last a few precession periods, taking place well before the beginning of planetestimal accretion. Therefore, the planets will form in a flat protoplanetary disk parallel to the binary orbital plane, and therefore the planetary orbits themselves have to be parallel to this plane.

The suggestion of photometric detection of extrasolar planets around eclipsing binaries was first developed by Schneider and Chevreton (1990). CM Draconis (with the exception of eclipsing binary systems with white dwarf components) is the smallest know eclipsing binary system with component dM3e and dM4 stars (Lacy 1977). Its characteristics relevant extrasolar planetary detection are listed in Table II (solar luminosities and limb darkening coefficients listed are for a wavelength of A = 0.82 #m). Being an eclipsing binary one may expect that any planets in the system will transit in the line-of-sight. Since the area of the two stars is only about 12% that of the solar disc, the transit of an Earth- to-Neptune-sized planet would produce a relative drop in the flux from 0.08% to 1.1%. Another important advantage of looking for planetary transits in such an eclipsing binary system is that a suspect transit can be almost immediately confirmed by a transit across the second star. Because of the displacement of the binary components these transits occur at only quasi-periodic times and thereby provide a unique signature for cross-correlation with models of each transit events


TABLE II

Summary of the CM draconis system parameters relevant to planetary detection (data from Lacy 1977)

Primary Secondary Relevance

1.26838965 Number of eclipse transits

0.534 -4- 0.001 0.466 4- 0.002 Planet formation locale

0.252 4- 0.008 0.235 4- 0.007 Magnitude of transit drop 0.237 4- 0.009 0.207 -I- 0.008 Duration of transits

0.50 4- 0.03 0.50 4- 0.03 Modification of transit depth

3150 4- 100 3150 4- 100 Observing filter (R-band)

0.00 4- 0.02 Simplify modelin correlation

89.82 -I- 0.05 Definitive transit events 1.75 4- 0.03 2.01 4- 0.07 Stable planetary orbit

locale 11.07 ~ 0.005 Meter-class telescopes OK

0.0105 4- 0.005 Known periodic noise source

Binary orbital period (days 4- O.OOO0001) Solar luminosity (solar luminosities) Radius (solar radii) Mass (solar masses) Primary limb darkening coefficient Primary effective temperature (K) Binary orbital eccentricity

Orbital inclination (degrees) Binary semimajor axis (a sini in solar radii) System R-band apparent magnitude Out-of-Eclipse sine wave magnitude variation (possibly starspots) Stellar flare rate (per hour)

System space velocity (km/sec), and eccentricity of galactic orbit Age of system implied by space velocity

0.02-0.05 Lower flare noise confusion

163 -I- 110.7 Population II star, enough "metals" to form planets?

6-10 Gyrs Planetary orbits remain stabel for this amount of time?

(Schneider and Chevre ton 1990, Brandmeie r and Doy le 1995, Jenkins et al.. 1995). Thus intrinsic stellar var iabi l i ty will not, in general, be a serious problem. Since the

stars are also t idally locked, pho tomet r i c variabi l i ty due to rotat ional modu la t ion

o f starspot act ivi ty on the stellar disk is also easi ly separable. By pho tomet r i ca l ly mon i to r ing the C M Draconis system, therefore, constraints on theories ofterrestr ial-

type planet fo rma t ion and evolu t ion m a y be obtained. Specific ques t ions that could

be addressed are out l ined in the next section.


3. Observational Constraints Possible on Planet Formation Proeesses

3.1. EMPHASIS ON TERRESTRIAL PLANET CONSIDERATIONS

A number of questions regarding the process of planetary system formation may be constrained by photometric observations of the CM Draconis system. Some important questions regarding terrestrial planet formation processes that could be formulated are, for example: (1) Do planets form in binary star systems? (2) Are planetary orbits stable in binary star systems? (3) Will planets form in the same plane as the binary orbital plane? (4) Can any noncoplanar planets be expected to precess across the line-of-sight in a reasonable amounts of time? (5) Will there be enough nebular material in the protoplanetary disk for planets to form in low- stellar mass systems? (6) Will planets forming around low mass stars consequently be smaller due to, for example, less protoplanetary material being available? (7) Will population II stars have enough "metals" (elements heavier than Hydrogen or Helium) to form planets? (8) Will planets form closer to lower luminosity stars or at the same distances, in general, as the scale in the solar system?

To address question (1), we note that infrared excesses of certain stars, as detected initially from the IRAS surveys, have been taken as evidence of circumstellar (perhaps protoplanetary) material around the star. This material has certainly been clearly evidenced in infrared images of the/3 Pictoris system and other consequent data. Statistical studies of infrared excesses around single compared to double stars (Backman and Gillett 1988, Beckwith and Sargent 1993, as examples) indicate that binary stars generally have infrared excesses as often as single stars do, implying that they are as likely to be surrounded by dust and gas as single stars are for extended time periods. If dust grains can remain stable for long periods of time in binary systems, then to expect the formation ofplanets in these systems may not be unreasonable. A dismissial of the possibility of planet formation in binary systems would therefore be, at best, premature, and such observations as outlined here could make a beginning at providing useful constraints on such planet formation processes.

(2) According to 3-Body calculations (Pendleton and Black 1983; Donnison and Mikulskis 1992, Kubala 1993, and references therein as examples) planetary orbits in binary systems should be stable for axis ratios of greater than about 3.3 (this applies to both outer planet as well as inner planet configurations, that is star-star to planet as well as star-planet to star axes ratios). The separation of CM Draconis A and B is about 3.75 solar radii (about 15.4 CM Draconis radii), which allows the nearest possible stable planet to orbit at about 12.4 solar radii (which would receive a stellar insolation equivalent to a planet in orbit about half way between the orbit of Mercury and Venus in the solar system). Thus most terrestrial planets, if formed near the same thermal regime as the solar system, could have stable orbits in this system (but see point 5 below).


(3) Models of the formation of protoplanetary nebulae universally expect angular momentum to be conserved by forming a disk essentially in the plane of the stellar rotation axis (Safronov and Ruskol 1994, as one recent example of many reviews). This conservation of angular momentum is offen modeled so definitively in the plane of grain accretion (with the width of the protoplanetary disk being supplied by turbulence, for example) that a disturbance of the solar system may have even been required to have perturbed the mean planetary plane out of the equatorial plane of the Sun (Laughlin and Bodenheimer 1994, for an example of non-axisymmetric evloution in protoplanetary disks). Massive polar jets may also have been a natural process of the rinal stages of star formation that could tilt the stellar equator to the plane of the protoplanetary disc Mundt 1989, for example). Recent initial work has indicated, however, that in the case of binary stars, the equatorial plane of stellar members would seem to lie very close to or in the binary orbital plane for separations less than about 30 A.U. (Hale 1994, Hale and Doyle 1994; however, the orbital planes of triple systems seem generally to be not so constrained). In addition however, as has been mentioned in Section 2.2, the precession of dust grain orbits in the protoplanetary disk induces a flattening mechanism which pull the disks into the plane of the binary. Thus, while to a first approximation one may reasonably expect planetary formation to, in general, be very close (+0.5 °) to the binary orbital plane. But this too, is a hypothesis testable by the observations suggested here as even planets out of the line-of-sight may be expected to eventually precess into the line-of-sight of the observer (point 4).

(4) The precession rate for a planet in a binary star system as given above depends on the separation of the binary s, weakly on the stellar masses M., and strongly on the planetary distance a, and can be re-written for the CM Draconis system as:

0 . 6 1 9 a 7/2 _ 0 .063a7 /2 , B,,- 82M1/2

(2)

where a and s are in solar radii, M, is in solar masses, and the precession period P,. is in days (remembering that one can expect two planetary transit events of two transits each every precession period). A terrestrial planet distance as close as such a stable orbit would allow (i.e. 3.3 times the binary semi-major axis; point 2 above) in the CM Draconis system would have an orbital period of about 7 days (and an orbital semi-major axis about 12.2 solar radii, as previously mentioned). In the CM Draconis system the precession period for the nearest stable planet will be about 58 times the planetary orbital period so that a non-planar planet would transit the system about every 6 and a half months (however, the probability of the planet being at the node when it precesses across the line-of-sight is inclination dependent and drops rapidly for inclinations > 10°; Schneider 1994a). Thus the hypothesis of planar versus non-planar close terrestrial planet formation could also be tested, albeit with significantly longer duration obselwations. This result introduces the

TERRESTRIAL EXTRASOLAR PLANETS AND CM DRACONIS .~ 61

possibility of including very close, yet non-eclipsing, binary systems as important candidates for terrestrial planet detection, and in particular of importance when large numbers of stars are to be monitored (Borucki et al. 1994, Koch and Borucki 1995, for example). In particular close, yet large mass systems such as binaries including white dwarfs could be attractive observational candidates (subject to the temperature constraints as we outline in Section 4 below in a discussion of possible future observations).

(5) Would there have been enough protostellar material for planets to have form in the small-mass CM Draconis system? In the solar system, although the present total planetary mass is only about 0.0015 solar masses, models indicate that the minimum mass for the protoplanetary disk to have formed planets would have to have been about 0.01 to 0.07 solar masses (Safronov and Ruskol 1994, Boss 1995). Higher mass disks would have "swallowed" all planetesimals, while lower mass disk would not have been expected to have formed gas giant planets. Therefore, even if we were to assume that the nebular mass scales as the stellar mass (CM Draconis components A and B have a total mass of 0.487 solar masses) a typical protostellar mass available around CM Draconis would still be in the range one could expect to allow for planet formation. In addition, the formation of the binary pair would also not have been expected to have used all initial angular momentum even if the binary pair had formed somewhat farther apart (shortening their semi-major axis by convective equilibrium tidal dissipation; Zahn 1975)

(6) Given that planets could form in the CM Draconis system, the question as to whether such planets would be of terrestrial-size or whether, if formed, they would scale as the stellar mass, is also of interest. Wetherill (1995, see also Wetherill and Stewart 1993), in work for single stars, indicates that Earth-mass or even larger terrestrial-type planets should form around single stars of 0.5 solar mass to 1.5 solar masses with approximately equal likelihood since "feeding zones" for accretional material would come from numerous distances around the protoplanetary nebula. It should be pointed out, however, that most of the planets in this model, form around 1 AU, also irrespective of the stellar mass. (We discuss this in point 7 below.) As a possibly relevant observation, and as mentioned in the introduction, terrestrial-sized and larger planetary masses have been observed around pulsars (Wolszczan 1994). In the case of PSR 1257 + 12, as mentioned, a 3M® (earth mass) body was detected at about 0.3 AU. This body would likely have to be of terrestrial density given that, although the central pulsar is small, they are modeled to have high surface temperatures (over 106 degrees) which may require close planetary mass bodies to be of terrestrial planet densities, (i.e. hydrogen would be evaporated). (This temperature constraint is again discussed in Section 4 with regard to red-dwarf/white-dwarf eclipsing binaries.)

(7) We now consider the question as to whether there might have been enough "metals" (elements that are not hydrogen or helium) available in the protoplanetary nebulae of CM Draconis, a Population II star, to have formed planets at all. The helium and metals content (Rucin'ski 1978, Paczyn'ski and Sienkiewicz 1984,


Vilhu et al. 1989) and, more persuasively, the galactic orbit (3-4 times the expected maximum Population I velocity; Rucin'ski 1978) of CM Draconis indicate that it is a Population II old halo object. Population II stars will generally have < 1% metals while Population I stars like the Sun will have about a 2-3% metal content (Allen 1976.) Thus, for example, although the TiO spectral lines of CM Draconis are comparable in depth to those of several Population I M stars, - (specifically Wolf 46, Krüger 60, and EV Lac, for example; Rodonò, 1986) - in this case they are likely due to the increase in gas pressure broadening rather than abundance for this very convective M star. CM Draconis is generally, then, considered to be a very metal poor star. As question 5 addressed, the mass range for planet formation in a protostellar nebula is about 0.01 to 0.07 solar masses. If the proto-CM Draconis nebula did indeed contain enough initial mass to form planets, the initial metal content, - perhaps 20% that of the Sun, - may still have been within the range of grain masses required. This hypothesis too, can be tested by long-term photometric observations of the system.

(8) But if planets did form around CM Draconis, where will they have formed? The inner boundary is constrained by the stability criterion for 3-body orbits mentioned in point 2 above (the accretion grain evaporation boundary is inside of this locale). However, recent work by Boss (1995) argues that due to the smaller thermal gradient of the protosolar nebula compared to the free space temperature profile, jovian planets should form in the region 4.5 to 6.0 AU for stellar masses of 0.1 to 1.0 solar masses rather than rauch closer as the fiee space location of the ice condensation temperature (about 160 ° K at the expected density of the protosolar nebula) would indicate. But the questions is, where can the terrestrial planets be expected to form? Gas giant planets taust form while the protostellar nebula is still present in quatities great enough to allow their runaway gas accretion; that is, gas giant planets should form in about 107 years and thereafter may assist the formation of inner terrestrial-like planets (Safronov and Ruskol 1994, and references therein). According to Safronov and Ruskol (1994) a robust result of modeling terrestrial planet formation is that they should form in about 108 years, - essentially after the accretion of the protostellar nebula responsible for the lower thermal gradient. It may be expected, then, that the formation of terrestrial inner planets would be rauch more dependent on the stellar mass (i.e. luminosity) than the formation of gas giant planets. For the purposes of beginning an observational program to detect inner terrestrial planets then, a free space thermal gradient should be a valid beginning assumption, that is:

Tp = T~ , (3)

where a is the condensation distance of proto-terrestrial planetary grains condens- ing at temperature Tp around a star of temperature T, and radius R,. If we start with the nearest stable planetary orbit possible in the CM Draconis system (point 2 above), and let the simple luminosity of the binary determine the formation of the


other terrestrial planets by analogy to the solar system, we obtain tenestrial planet periods of from about 4 days to about 27 days (the Earth-equivalent insolation forming a planet of period about 18 days, as mentioned). Continuous photometric observations of this system over periods as short as several months (and with medium-sized telescopes) could begin, then, to constrain hypotheses of terrestrial planet formation. For such continuous observations, a network of telescopes at varying longitudes are needed.

3.2. PHOTOMETRIC DETECTION OF JOVIAN-SIZED PLANETS WITHOUT A

TRANSIT

The detection of jovian-sized bodies that might exist in this system would also be possible without having to wait for the transit of such a planet. Periodic changes in the binary barycenter would be caused by the existance of one or more giant planets in the system, and therefore the expected time of eclipse miniuma, for example, would periodically change by an amount:

~- M--~ \~~~.] ' (4)

where lJlp is the planet's mass (in jovian masses, M j), M~, is the total stellar mass(es), as before (in solar masses, Mo), and e is the speed of light. For a jovian-sized planet in orbit around the CM Draconis system at a distance of 5 AU, the change in eclipse minima would be about 10 seconds, a quantity easily within the reach of standard photometric observations of eclipse minima at these magnitudes.

In this section then, we find that there are presently no intrinsic theoritical rea- sons to conclude that terrestrial-type inner planets have not formed in the CM Dra- conis system and observations looking for such extrasolar terrestrial-sized planets in this system could give significant clues to understanding inner planet formation regardless of whether a detection is made or not. The existance of jovian-sized planets in the system will also be possible using the same photometric data. In the next section we outline preliminary efforts of such an observational program.

4. Present Observing Program: The TEP (Transit of Extrasolar Planets) Network

In order to assure that non-detections also provide a valuable result, the CM Dra- conis system should be observed continuously over the expected orbital period of any terrestrial-type planets, and with as much precision as possible. One observing advantage of this system is that it is circumpolar for most observatories in the northern hemisphere (Dec = +57 ° 15', RA = 16h 33.5m). In addition it has an R-magnitude, as mentioned, of 11.07 (12.90 in the V-filter; Lacy 1977) and


TABLE III Presently collaborating observatories in the detection of transits of extrasolar planets (TEP) network

Observatory Longitude Latitude telescope (meters)

Haute Provence Observatory (France) 3 E +43 1.2 m Canary Islands Observatory (Spain) 15 W +28 0.8 m Rochester Institute of Technology Observatory (New York) 78 W +41 0.6 m University of California, Lick Observatory (California) 122 W +38 0.9 m Korean Astronomy Observatory (South Korea) 127 E +36 0.6 m University of Crete Observatory (Greece) 25 E +35 1.3 m

Summary of coverage (westward from Haute Provence in a 24 hour circle): 0 hours 1 h 5 h 8 h 16 h 22 h 0 hours Provence ~ Canaries -+ RIT ---+ Lick ~ Korea ---+ Crete ---+

therefore meter-class telescopes can reasonable collect enough photons to resolve transit events sufficiently. We have therefore set up a network of small telescopes called TEP (Transit of Extrasolar Planets), which consists of a series of 0.6 to 1.3 -meter telescopes at well-spaced longitudes using CCD arrays which may theoreti- cally provide ifferential photometric precision to within one part in 10-1 (Gilliland

and Brown 1992; Borucki et al. 1988). A list of the present TEP collaborating observatories is given in Table III. It is estimated that each observatory will be able to observe CM Draconis for an average of about 8 -9 hours per night, so that continuous coverage may be achieved by the present network over 1-2 month-long

observing runs. Characteristics of the CM Draconis system are weil known (Table II) so that

simulations of planetary transits in the system can be well modeled. The light curves of such planetary transits will vary in a quasi-periodic way since a planetary transit event could occur for numerous different configurations of the two stars across the line of sight. The typical transit pass (about 85% of the transits) of a habitable-zone (terrestrial insolation) planet in the CM Draconis system will be a drop in the brightness for a bit more than one hour (planet traveling with the first star) fol lowed by a transit drop slightly shorter than one hour (second star crossing behind the transiting planet; Brandmeier and Doyle 1995). For a smaller, but significant, number of transits (perhaps as much as 14% of the time)


such a planet will transit the CM Draconis stars while they are in eclipse, and such configurations can provide transit events that can be ten hours in duration or longer, providing consequent easier detection by providing more transit points in one pass (Brandmeier and Doyle 1995, Jenkins et al. 1995; this result is applied in the calculation in Section 6). Triple events are also possible for such a system, and even higher multiples of events are possible for planets farther out (and, of course, the transit duration increases outward, as do certain transit characteristics).

Such quasi-periodic transit events are unusual, and as such, provide a way as to avoid general confusion with the white Gaussian noise expected from the ground-based observations proposed here. The signal can be cross-correlated (as a masked filter) with the simulated planetary transitfoinary light curve to pull out sub-noise signals in an additive way - that is, longer periods will be able to reach consequently smaller planet sizes, as pointed out in Jenkins et al. (1995), instead of reaching an absolute atmospheric scintillation limit as was previously thought. Since the CM Draconis system stars are tidally locked, variations due to starspot modulation will also not be confused with a planetary transit event as it might be the case for single star transit events (Koch and Borucki 1995). The net result is that approximately 1-meter class telescopes may be utilized over several months to detect terrestrial-sized planets from the ground for this system. A series of well spaced (in longitude) robotic telescopes would, of course, be ideal for this type of observation, as well.

This masked filter method could also allow the detection of extrasolar planets in other small eclipsing binary systems with larger ground-based telescopes with observations over a more extended time period, as well. For example, in addition to CM Draconis, there area number of smaller eclipsing binary systems that might be successfully observed. Listed in Table IV are the ten smallest component eclipsing binary star systems presently catalogued from Brancewicz and Dworak (t980), with the star name, orbital period, spectral types of the two components (1 is the primary), respective stellar radii (in units of solar radii), respective luminosities (in solar luminosities), the distance to the binary system in parsecs, and the equivalent stellar area, (RA) 2 + ( R B ) 2 with R in solar radii. In cases where the secondary component spectral type has not been directly determined (marked "ukn.") it is nevertheless known to be less massive than the primary star. They are arranged in order of smallest total stellar surface areas. We see that while the CM Draconis system has a significant advantage over even the next smallest system in terms of both area and luminosity (the latter as it realtes to the possible close location of planet formation), most of the other systems have the advantage of smaller binary semi-major axes (closer stable planetary orbits) and shorter periods (more binary star transits "under" the planet per planetary transit pass).

There are also other possibly interesting targets for this method that could be of significant astrophysical interest with regard to understanding the formation of planetary bodies around more massive evolved stars. These are the eclipsing binary systems containing white dwarf components. The transit of even very small


TABLE IV Sample of observable eclipsing binaries with both components either lx" and/or M dwarfs

Star name P (days) SpA/SpB RA/I~B LA/LB D (pc) Area*

CM Dra 1.268400 dM3e/dM4: 0.25/0.24 0.01/0..01 16.9 0.12 YY Gem 0.814283 Mle/Mle 0.60/0.60 0.06/0.06 14.7 0.72 XY Leo 0.284110 K 0 / K 0 0.73/0.46 0.43/0.14 62.5 0.75 RW Dor 0.285464 K5/MO: 0.64/0.64 0.10/0.06 28.6 0.82 BI Vul 0.251828 K0:/K0: 0.75/0.75 0.33/0.33 333.3 1.14 RT Hyi 0.284038 K0:/K0: 0.78/0.74 0.30/0.27 250.0 1.15 TZBoo 0.297160 K0:/K2: 1.16/0.43 0.71/0.08 114.9 1.53 V947 Oph 0.371860 K0:/ukn. 0.90/0.88 0.54/0.41 384.6 1.59 MR Cas 0.352936 K2:/ukn. 1.08/0.84 0.63/0.40 434.8 1.89 TZ Lyr 0.528823 K0/(K9) 0.99/1.25 0.65/0.28 76.9 2.56

*Equivalent area is (RA) 2 + (RB) 2 in solar units.

planetary bodies would produce a significant drop in the brightness of the system in the blue region of the white dwarf contribution, and the transit of such a terrestrial- sized planet in the line-of-sight could be expected to completely occult the white dwarf component producing a momentary total eclipse. The distribution of the semi-major axes of binary white dwarf systems appears to be bimodal (Iben and Tutukov 1993, de Kool and Ritter 1993, for example) and it is implied that the process of energy loss via the common planetary nebula envelope is responsible for the decrese in binary angular momentum seen in a shortening of the mutual semi-major axis. It would be of interest to investigate in detail what the effect of such an evolutionary event would have on the semi-major äxis of any orbiting planetary bodies ofiginally in the system. Mass loss in a single star system - unless it was dense and persistent enough to drag the planetary bodies inward, may be expected to have cause any planetary bodies to increase their orbital semi-major axes according to Whitmire et al. (1995). However, a competing effect is that as the mutual binary semi-major axis is shortened, each component star may be expected to "bring" any original planets along with it toward the other star, providing a more easily detectable (both in time and possibly alignment), scaled-down version of the original "solar" system configuration for study. (According to the data of Iben and Tutukov 1993, the semi-major axes of such systems could be scaled down by a mean factor of about 50). Of course, such a system could not be considered to define a habitable zone around the stars due to the extremely high temperature of the white dwarf. But detection of remnant planets in such a system would nevertheless be of astrophysical interest as versions of an original planetary configuration in an evolved form. If the planets survived, or were not ejected from the system altogether, they could either be captured into orbit around both stellar components


(the outer planets scenario that we have previously discussed), remain in orbit around each star (the inner planets scenario where planets orbit each single star member of a binary system). The modeling of these latter types of configurations would produce a whole new family of planetary transit shapes and configurations that have not, as yet, been explored. If we were to assume the outer planets scenario for such systems and that planetary orbital planes were close but not exactly coplanar with the line-of-sight, we see (referring to Equation (2)) that such systems contain the large stellar masses and extremely short binary orbital periods needed to optimise the rate of precession of any non-planar inner planets. However, the existence of planetary bodies so close to white dwarf stars is problematic since their temperatures are so high. Rewriting Equation (3), (and including the planetary albedo term A = 0.39 for a terrestrial-type planet) we have:

a = ( 1 - A)~, (T , /Tp) 2, (5)

for the distance a at which a terrestrial-type planet could form. Here Rs, is the stellar radius (0.02 solar radii for the white dwart), T~ is the stellar temperature (50,000 °K for the white dwarf), and Tp is the planetary temperature. According to Boss et al. (1989) terrestrial planets would likely not form at temperatures higher than about 800 °K, while terrestrial planets might survive at temperatures closer to 1200 °K.

In Table V are four sample observational candidates of such eclipsing binary systems containing an M-dwarf main-sequence and a white dwarf component along with their spectral type, binary oi:bital period (in days), apparent magnitude, stellar masses (in solar units; a dash means not yet determined), equivalent area (as defined in Table III, estimated from spectral type) and orbital period (in days) of the closest possible terrestrial planet based on planetary temperatures of 800 °K and 1200 °K, (giving planetary distances of about 48 and 21 solar radii), respectively. In the case of RR Cae we assume a total system mass of 1.1 solar masses based on the spectral type. For the system GK Vir, whose binary inclination to the line of sight is only known to be >75 °, we assume that a planetary orbit would cross the line-of-sight, but it may have to precess across the line-of-sight (see Equation (2); Schneider 1994a) if the system is not very close to being eclipsing. In summary, continuous photometric observations of these systems may also therefore provide observational constraints on the existance of any post-nova planets that would either represent post-stellar-evolution formation (suggested as the case for pulsar planets) or an interesting modified version of the original plantary system.

5. Speculative Possible Future Observations of CM Draconis

In addition to being of astrophysical interest (as the candidate observational stars listed in Table IV may be), the detection of terrestrial-sized extrasolar planets at the "liquid water distance" around any main-sequence stars could be of exobiological


TABLE V

Planetary orbital periods (formation temperature dependent) around a sample of short-period M-dwaff/white-dwarf eclipsing binaries

Star name •pA/Spt3 P (binary) mag~ mA~mB Eq. area P (800) P (1200)

V471 Tau K2/DA2 0.521183 9.2 0.71/0.73 0.66 32.7 9.4

RR Cae Me/DA 0.30371 14.4 - / - 0.16 37 11 NN Ser M5/DO 0.13008 16.6 0.54/0.12 0.10 47.8 13.8

GK Vir* M3/D 0.344331 17.0 0.7/0.4 0.16 37.0 10.7

* Inclination not yet fully determined hut known to be _ 75 °.

interest, as well. Returning for consideration of this point to the CM Draconis system again, if terrestrial-sized planets were detected within the habitable zone of this system, (the habitable zone being defined as in Kasting et al. 1993, Whitmire and Reynolds 1995), then future observations might entail an effort to detect possible atmospheric characteristics of such planets. What would be required to make such a detection? To obtain the spectrum of an extrasolar planet is extremely difficult, but not as difficult in absorption (transiting planet) as in reflection (non- transiting planet; Schneider 1994b). Many atmospheric absorption lines would be of great interest, but those of free oxygen, 02 or 03, are so reactive that the detection of such a line in a planetary atmosphere, ~ although possibly more difficult than some lines, would be highly indicative of biological process such as photosynthesis, constantly producing free oxygen (Owen 1980, Leger et al. 1994, 1995). (We note that a planet in the process of a runaway greenhouse would also show massive oxygen lines, hut this process would be expected to be a rapid event on the planetary timescale; Kastirlg 1995.) The time for detection of, for example, the 7600 Ä 02 line in a transiting terrestrial-type planet can be parameterized as follows (after Schneider 1994b):

~ ~ ~~107~,004~~~~(~ )~(~~;) ~

( ~ ) ( 5 ) ' w (6)

where T is the time for detection (in seconds), m is the stellar magnitude, f = 1 (saturated spectral band; although M-star spectra can be quite complicated in this region of the spectrum), c is the detection efficiency (we will take 0.3 hefe), H is the planetary scale height (= 30 km here), Rp ,R®, R., Rsun, are the planetary radius, Earth's radius, the stars' radius, and the solar radius, respectively, A is the telescope diameter in meters, and W is the width of the spectral line (here 30 Ä).


TABLE VI

Field star V-magnitude contributions to arcminute area around CM draconis

Circular aperature (arcminutes) Additional stellar brightness fraction (excluding sky)

< 1' 0.16

< 2' 0.46

< 3' 0.59

< 4' 0.79

< 5' 1.12

< 6' 2.63

We will also take Rp = R e here. For CM Draconis, rn = 11.07 at 7600 Ä, and /~. = (0.252 + 0.235)Rsun. For a planet on the inner boundary of the habitable zone of CM Draconis (Kasting et al. 1993), two planetary transits may be expected to occur once every ten days for typically one hour each, except that about 1/6th of the transits could occur during stellar eclipse and last for significantly longer each (Brandmeier and Doyle 1995, Jenkins et al. 1995). Assuming that about 100 transits (about half) can be observed during the observing season, and that 1/6th are about ten hours each in duration, we obtain T ~ 8 × 105 seconds for the total spectral integration time possible over the observing season available during one year for CM Draconis. Solving for A we obtain the total collecting surface area required for detection of about 2.55 × 104 square meters, or an equivalent circular aperature of about 90-meter diameter. Since simply a spectrum is desired, precision focusing of such a collecting surface would only be required to avoid gathering light from other stars in the field. As can be seen in Table VI, from the faifly sparse field of view of CM Draconis, flux from nearby stars would not contribute much brightness to the spectra of the field even for a very low resolution instrument.

Therefore, the resolution of such a light-gathering surface for the spectrometer should be better than 4 minutes of arc and, if possible, better than 1 minute of arc (that is, most large light gathering aperatures available such as ordinary solar collectors, would not, in general, have sufficient resolution.) The radial velocity of CM Draconis is -118 .6 ± 0.8 km/sec which gives a Doppler shift in the 7600 2k 02 line of 3 Ä but it may still be greatly diminished by the Earth's absorption bands (Bézard et al. 1995; the 9.6 #m 03 line would be expected to be even more diminished by terrestrial absorption bands). Thus a 7600 Ä 02 absorption line (with a typical terrestrial atmospheric width of about 250 ~) from a terrestrial- type planet around CM Draconis would still be very much on the wings of the terrestrial 02 absorption feature. Nevertheless, if terrestrial-sized planets within the habitable zone were detected around the CM Draconis, the detection of biological


processes would remain, for this system, an interesting and not impossible near- future prospect.

Although planets within the habitable zone of an M-star will, in general, be tidally locked in their rotation (Kasting et al. 1993) so that previous notions were that such planets would not have unfrozen atmospheres, recent work (Haberle et al. 1995) has shown that with very little additional CO2 (100 millibars) the thermal gradient can be reduced to the point where liquid water can be maintained on the surface (this calculation was done for a 1 bar atmosphere of H20). Stable habitable zones have also been shown to exist for binary stars (see Hale 1995 for a recent review). Consequent work on the adaptability of complex biological systems (Heath et al. 1995) have shown also that essential biological processes (e.g. photosynthesis) could take place under such circumstances. Consequently, CM Draconis, in addition to being a good place to look for extrasolar planets, could also become an optimal place to begin the search for exobiological systems, as well, if suitable planetary environments are detected. Once the observational constraints can be mitigated, then given the duration of 02 on out own planet (at least 50% of the history of our star) the prospects for actual detection of biological processes, if they are taking place around the CM Draconis system, could be quite reasonable.

6. Conclusions

Terrestrial-sized (Earth-to-Neptune-radii) extrasolar planets may be detectable in the CM Draconis system and several others using a round-based network of 1- meter-class telescopes performing CCD photometry over several months. The small size of the system, and the orbital plane being edge-on, enhance both the probability as well as the sensitivity of detection. Consequent cross-correlations of modeled transits with observational data should allow the detection of sub-noise signals equivalent to those produced by the transit in this system of less than two-Earth-radii planet in a matter of a few months. This technique is extendable to the determination of the existance of suviving planetary bodies around close M-dwarf/white dwarf binaries even for planets out of the line-of-sight. If no such detections are forthcoming after these series of observations, we should nevertheless have bonif ide constraints on the existance of terrestrial-sized planets in other star systems. Such non-detections would place interesting constraints on models of the size and formation locale of terrestrial planet formation around low mass binary systems (the CM Draconis stellar configuration, M-star binary, could be representative of perhaps as many as one-third of the star systems in the galaxy; Allen 1976). In the case of positive planetary detections, CM Draconis may be singled out to also provide the most excessible first possibility for the study of extrasolar planetary atmospheres, as well, and speculatively, the possible presence of any exobiological influences on that atmosphere.


Acknowledgements

The authors would like to acknowledge the other members of the TEP network: M. Chevreton and R Freire of Observatoire de Meudon (France); H.-J. Deeg and E. Martin of Instituto d'Astrofisica d'Canarias (Spain); J. Jenkins and N. Heather of SETI Institute and T. Dunham of the Space Science Division, at NASA Ames Research Center (Califomia); Z. Ninkov of Rochester Institute of Technology (New York); E. Paleologo and N. Kylafis of University of Crete (Greece); W.- B. Lee of the Korean Astronomy Observatory in Taejon (Korea); D. Toublanc of University of Bordeaux (France); S. Brandmeier of Polytechnic University of Konstanz (Germany); S. Jeschke of University of Berlin (Germany); and C. Sterken of Brussels Vrije Universität. LRD was supported for this work by a NASA Ames Research Center Director's Discretionary Fund and a one-month visiting professorship to the Observatoire de Meudon, France.

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ground-based detection of terrestrial extrasolar planets by photometry: the case for cm draconis

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