Homogeneous Photometry VI: Variable Stars in the Leo I Dwarf Spheroidal Galaxy
Peter B. Stetson, Giuliana Fiorentino, Giuseppe Bono, Edouard J. Bernard, Matteo Monelli, Giacinto Iannicola, Carme Gallart, Ivan Ferraro
aa r X i v : . [ a s t r o - ph . GA ] J un Homogeneous Photometry VI: Variable Stars in the Leo I DwarfSpheroidal Galaxy ∗ Peter B. Stetson , Giuliana Fiorentino , Giuseppe Bono , , Edouard J. Bernard ,Matteo Monelli , , Giacinto Iannicola , Carme Gallart , , and Ivan Ferraro drafted June 11, 2018 / Received / Accepted ABSTRACT
From archival ground-based images of the Leo I dwarf spheroidal galaxy wehave identified and characterized the pulsation properties of 164 candidate RRLyrae variables and 55 candidate Anomalous and/or short-period Cepheids. Wehave also identified nineteen candidate long-period variable stars and thirteenother candidate variables whose physical nature is unclear, but due to the limi-tations of our observational material we are unable to estimate reliable periodsfor them. On the basis of its RR Lyrae star population Leo I is confirmed tobe an Oosterhoff-intermediate type galaxy, like several other dwarf spheroidals.From the RR Lyrae stars we have derived a range of possible distance moduli forLeo I: 22 . ± . ∼ < µ ∼ < . ± .
07 mag depending on the metallicity assumedfor the old population ([Fe/H] from –1.43 to –2.15). This is in agreement withprevious independent estimates. We show that in their pulsation properties, theRR Lyrae stars—representing the oldest stellar population in the galaxy—are notsignificantly different from those of five other nearby, isolated dwarf spheroidalgalaxies. A similar result is obtained when comparing them to RR Lyrae stars inrecently discovered ultra-faint dwarf galaxies. We are able to compare the period Dominion Astrophysical Observatory, NRC-Herzberg, National Research Council, 5071 West SaanichRoad, Victoria, BC V9E 2E7, Canada INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy Dipartimento di Fisica, Universit`a di Roma Tor Vergata, via della Ricerca Scientifica 1, 00133 Rome,Italy INAF-Osservatorio Astronomico di Roma, via Frascati 33, Monte Porzio Catone, Rome, Italy SUPA, Institute for Astronomy, University of Edinburgh, Royal Observatory, Blackford Hill, EdinburghEH9 3HJ, UK Instituto de Astrof´ısica de Canarias, Calle Via Lactea, E38200 La Laguna, Tenerife, Spain Departamento de Astrof´ısica, Universidad de La Laguna, Tenerife, Spain ∼ ∼ ∼ d G ∼ <
14 kpc) to the outer ( d G ∼ >
14 kpc) halo regions. This suggeststhat the halo formed from (at least) two dissimilar progenitors or types of progen-itor. Considered together, the RR Lyrae stars in classical dwarf spheroidal andultra-faint dwarf galaxies—as observed today—do not appear to follow the welldefined pulsation properties shown by those in either the inner or the outer Galac-tic halo, nor do they have the same properties as RR Lyraes in globular clusters.In particular, the samples of fundamental-mode RR Lyrae stars in dwarf galaxiesseem to lack H igh A mplitudes and S hort P eriods (“HASP”: A V ≥ . ∼ < representative examples of the original building blocks ofthe Galactic halo. Subject headings: galaxies: dwarf — galaxies: individual (Leo I) — galaxies:stellar content — stars: variables ∗ Based in part on data obtained through the facilities of the Canadian Astronomy Data Centre operatedby the National Research Council of Canada with the support of the Canadian Space Agency; data obtainedfrom the ESO Science Archive Facility under muliple requests by the authors; data obtained from the IsaacNewton Group Archive, which is maintained as part of the CASU Astronomical Data Centre at the Instituteof Astronomy, Cambridge; and data distributed by the NOAO Science Archive. NOAO is operated bythe Association of Universities for Research in Astronomy (AURA) under cooperative agreement with theNational Science Foundation.
1. Introduction
The nearby dwarf spheroidal galaxies (dSphs, L ∼ > L ⊙ ) are important for under-standing the formation and evolution of the Milky Way Galaxy. Their very old ages( ∼ >
10 Gyr) and low mean metallicities ([Fe/H] ∼ < –1.3 dex) have long been used as evi-dence for the idea that these systems are fossils from the early Universe: surviving exam-ples of the pre-galactic units that formed the Galactic stellar halo, as predicted by hier-archical models of galaxy formation (see White & Frenk 1991; Salvadori & Ferrara 2009, 3 –and references therein). This view has been widely held for at least the last twentyyears, since detailed stellar population studies have been made possible by new deepcolor-magnitude diagrams (CMDs) and medium- and high-resolution spectroscopic anal-yses (Mateo 1998; Tolstoy et al. 2009). Today, we have at our disposal accurate star-formation histories (Tolstoy & Saha 1996; Gallart et al. 1999; Cole et al. 2007; Monelli et al.2010a,b; Hidalgo et al. 2011; de Boer et al. 2012a,b) and chemical abundance distribu-tions (Venn et al. 2004; Kirby et al. 2009, 2011; Fabrizio et al. 2012; Lemasle et al. 2012;Kirby et al. 2013) for a large number of nearby dSphs and dwarf irregulars (dIs). However,high-resolution spectroscopy has shown that the [ α /Fe] abundance patterns of stars in dSphsdo not seem to resemble those of Galactic halo stars (Shetrone et al. 2001; Venn et al. 2004;Pritzl et al. 2005). This could pose some problems for their identification as representativebuilding blocks of the halo. On the other hand, it could also be explained by the fact thatpresent-day dSphs, having not accreted at early epochs, have had a Hubble time to evolveas independent entities under conditions that are different in a variety of ways from thoseexperienced by field halo stars. In support of this interpretation, Font et al. (2006) haveargued that differences in the chemical abundances of the Galactic halo and its satellites canbe explained within their hierarchical formation model of the halo, which carefully tracksthe orbital evolution and tidal disruption of its satellites.A new source of puzzlement was raised by Helmi et al. (2006), who obtained largesamples of low- and intermediate-resolution spectra of red giant branch (RGB) stars ofthe Sculptor, Sextans, Fornax, and Carina dSph galaxies as part of a huge observationaleffort devoted to the characterization of the chemical abundances of those systems (seeTolstoy et al. 2006, and references therein). For the first time, Helmi et al. (2006) noticed asignificant lack of extremely metal-poor stars ([Fe/H] ∼ < –3 dex) in the galaxies’ metallicitydistributions when compared with that of the Galactic halo. The most metal-poor stars in agalaxy are likely to be also the oldest ones so—taken at face value—this evidence again seemsto argue against the idea that nearby dSphs are surviving examples of the main contributorsto the Milky Way halo. In order to further investigate this issue, Starkenburg et al. (2010)revisited the calibration of [Fe/H] as inferred from the observed strength of the infraredcalcium triplet, which has been widely used to estimate the metal content of large samples ofRGB stars in nearby dSph galaxies. The initial calibration of this relationship had presumedthat it was linear over the relevant range of abundance, but in the new study it turned outto be nonlinear for [Fe/H] ∼ < –2 dex, thus modifying the shape of the metallicity distributionin the metal-poor tail. Using the new calibration, Starkenburg et al. suggested that classicaldSphs are not as devoid of extremely low-metallicity stars as previously believed. Selectingthose stars with [Fe/H] ∼ < –3 dex based on this new calibration of the IR calcium triplet,Starkenburg et al. (2013) followed up with high-resolution spectroscopy that confirmed the 4 –stars’ extremely low metallicities. Today more than 30 extremely metal-poor stars havebeen identified in dSphs, and their chemical abundance patterns now seem to more closelyresemble those in the Galactic halo (e.g., Kirby et al. 2008; Koch et al. 2008; Kirby et al.2009; Frebel et al. 2010a,b; Starkenburg et al. 2013).Furthermore, the discovery of ultra-faint dwarf galaxies (UFDs, L ∼ -10 L ⊙ Willman et al. 2006; Zucker et al. 2006; Belokurov et al. 2006) seems to shed some morelight on this intriguing scenario. They are numerous and very old ( ∼ >
10 Gyr), and theycontain extremely metal-poor stars (Kirby et al. 2008). Thus to some investigators theyshow promise as ideal candidates for the original building blocks of the Galactic halo, thatmay at the same time help solve the missing-satellite problem (see White & Frenk 1991;Salvadori & Ferrara 2009, and references therein). Moreover, spectroscopic investigations ofseveral UFDs provide strong constraints on the dynamical interaction between their baryoniccontent and their dark matter halos (Mayer 2010; Gilmore et al. 2013)Tracing the oldest stellar component in a galaxy, RR Lyrae variable stars (RRLs) havethe potential to independently test this interesting global picture through their pulsationproperties. In particular, RRLs observed in both the Galactic halo and globular clus-ters (GCs) show a well defined dichotomy in their period distribution (see discussion inContreras Ramos et al. 2013, and references therein). This behavior was first observed byOosterhoff (1939) in Galactic GCs: he divided them into two groups according to the meanperiod of their RRab variable stars. Specifically, Oosterhoff I (OoI) GCs show mean periodsaround h Pab i ∼ .
55 d, whereas OoII GCs show longer periods ( h Pab i ∼ .
65 d). Furtherobservations disclosed that the mean periods of the RRc-type variables are also different—again shorter in OoI clusters—and also that the fraction of RRc variables among RR Lyraesof all types is systematically smaller in OoI ( ∼ ∼ ∼ –2 dex) halo objects, while OoI clusters cover a broader range of generally highermetallicities. At present—75 years later—there is still no full and satisfying explanation forthis phenomenology among the GCs.More recently, major observational surveys have shown that field RRLs in the Milky Wayhalo show the same period dichotomy as those in the globulars (Bono et al. 1997; ASAS,Pojmanski et al. 2005; LONEOS: Miceli et al. 2008; LINEAR: Sesar et al. 2013, etc.). Incontrast, those in nearby Local Group galaxies and their GCs are perversely characterizedby mean periods precisely in the range 0 .
58 d < h Pab i < .
62 d—the so-called “Oosterhoffgap”—that is avoided by the Milky Way halo cluster and field populations (see Bono et al.1994; Catelan 2009, for a detailed discussion). However this comparison, based on the 5 –statistical properties of very different samples of stars that may be subject to differentselection biases, is not truly decisive in characterizing the differences between GCs, theGalactic halo, and dwarf galaxies.Recently, extensive variability searches to characterize the RRL populations in UFDshave been carried out (Siegel 2006; Dall’Ora et al. 2006; Kuehn et al. 2008; Greco et al.2008; Moretti et al. 2009; Musella et al. 2009; Dall’Ora et al. 2012; Musella et al. 2012;Clementini et al. 2012; Garofalo et al. 2013; Boettcher et al. 2013). Given the very smallsample of RRL stars observed in most of the cases, it is very difficult even to provide aclear Oosterhoff classification for most UFDs. A few exceptions do exist: Bootes (11 RRLs,see Dall’Ora et al. 2006) is the sole example of an apparently well-established OoII type,whereas Ursa Major I (7 RRLs, see Garofalo et al. 2013) is classified as Oo-intermediate,resembling the classical dSph galaxies. Indeed, at the present date we might regard thedistinction between UFDs and dSphs as largely a matter of semantics: according to theavailable evidence, the two terms may merely refer to the lower- and higher-mass membersof a single family of objects (see, e.g., McConnachie 2012, especially Figs. 6, 7, and 12).This paper is part of the effort of one of us (PBS) to establish and maintain a database ofhomogeneous photometry for resolved star clusters and galaxies. This effort seeks the great-est completeness and longest time span possible by exploiting public-domain data currentlyavailable from astronomical archives. In this paper we analyse in detail the properties ofvariable stars that we have identified in the Leo I dwarf spheroidal galaxy. Part of these datahave already been presented in Fiorentino et al. (2012), where we focused our attention onunderstanding the intermediate-age central-helium-burning stellar population (Anomalousand/or short-period Cepheids, hereinafter AC/spCs). Here we present the full infrastructureunderlying that discussion: identifications and equatorial positions for the complete set ofvariable star candidates together with their periods and light curve parameters. In § § § §
2. Observations
The observational material for this study consists of 1,884 individual CCD images ob-tained on 48 nights during 32 observing runs. These data are contained within a muchlarger data collection ( ∼ ∼
500 observing runs) compiled and maintained bythe first author. Summary details of all 32 observing runs used here are given in Table 1.In fact, the observing runs designated “23 wfi4” and “24 wfi” were the same run, as werethose designated “27 suba” and “28 suba2.” However, in each case different subsets of theimages had been provided to us through different channels. To avoid any possible, unknowndifferences in the ways the images had been treated before we obtained them, we chose tokeep them separate. Considering all the Leo I images together, the median seeing for ourobservations was 1 ′′ .0 arcseconds; the 25th and 75th percentiles were 0 ′′ .8 and 1 ′′ .3 arcseconds;the 10th and 90th percentiles were 0 ′′ .7 and 1 ′′ .6 arcseconds.The photometric reductions were all carried out using standard DAOPHOT andALLFRAME procedures (Stetson 1987, 1994) to perform profile-fitting photometry, whichwas then referred to a system of synthetic-aperture photometry by the method of growth-curve analysis (Stetson 1990). There were insufficient U -band observations from photometricoccasions to establish a reliable U magnitude system in Leo I, and we have made no attemptto calibrate those data photometrically. The few U -band images that we do have, how-ever, were included in the ALLFRAME reductions to exploit whatever information they canprovide on the completeness of the star list and the precision of the astrometric positions.Calibration of these instrumental data to the photometric system of Landolt (1992) (seealso Landolt 1973, 1983) was carried out as described by Stetson (2000, 2005). If we define a“dataset” as the accumulated observations from one CCD chip on one night with photometricobserving conditions, or one chip on one or more nights from the same observing run withnon-photometric conditions, the data for Leo I were contained within 95 different datasets,each of which was individually calibrated to the Landolt system. Of these 95 datasets, 63were obtained and calibrated as “photometric,” meaning that photometric zero points, colortransformations, and extinction corrections were derived from all standard fields observedduring the course of each night, and applied to the Leo I observations. The other 32 datasetswere reduced as “non-photometric”; in this case, color transformations were derived fromall the standard fields observed, but the photometric zero point of each individual CCDimage was derived from local Leo I photometric standards contained within the image itself;these local standards were established by us from the images obtained during photometricconditions. Considering the 95 datasets employed here, the median values for the root-mean-square residuals of the observed results for the standard stars relative to their establishedstandard values were 0.029 mag in B (minimum 0.007, maximum 0.079), 0.026 in V (0.011, 7 –0.076), 0.024 in R (0.007, 0.059), and 0.024 in I (0.011, 0.090). The (minimum, median,maximum) numbers of standard-star observations contained in each dataset were (88, 1644,15084) in B , (79, 2059, 25664) in V , (85, 1440, 8742) in R , and (26, 1526, 14790) in I .The different cameras employed projected to different areas on the sky, and of coursethe telescope pointings differed among the various exposures. Furthermore, the Wide-FieldCamera on the Isaac Newton Telescope, the Wide-Field Imager on the 2.2m MPG/ESOtelescope on La Silla, Suprime-Cam on the Subaru Telescope, FORS on the VLT on Paranal,and LBC on the Large Binocular Telescope on Mount Graham consist of mosaics of non-overlapping CCD detectors. The last column of Table 1 gives the number of independentdetectors in each camera when this number is greater than one. Therefore, although wehave 1,884 images, clearly no individual star appears in all those images. In fact, no starappeared in more than 97 B -band images, 240 V images, 51 R images, or 41 I images. The median number of observations for one of our stars is 22 in B , 72 in V , 24 in R , and 16 in I . Since most pointings were centered on or near the galaxy, member stars typically havemore observations than field stars lying farther from the galaxy center.Our complete photometric catalog for Leo I and a stacked image of the field are availableat our web site
3. Variable star detection and characterization
We have used an updated version of the Welch & Stetson (1993) method to identify can-didate variable stars on the basis of our multi-band photometry of Leo I. We have identified251 variable star candidates in a luminosity range that goes from the Horizontal Branch (HB)to the tip of the red giant branch (TRGB). A first guess at the variability period of each starwas found using a simple string–length algorithm (Stetson et al. 1998b), then a robust least-squares fit of a Fourier series to the data refined the periods and computed the flux-weightedmean magnitudes and amplitudes of the light curves (Stetson et al. 1998b; Fiorentino et al.2010, and references therein). The resulting pulsation properties are listed in Tables 2, 3 and4. This list includes 164 stars that we have classified as RRLs, 55 AC/spCs and 19 LPVs.There are also thirteen more variable candidates whose physical nature is unclear (Table 5).Some of their light curves are shown in Fig. 1, and all are available in the on–line version ofthe paper and at our web site.On the basis of the photometric quality of the data and the completeness of the phase ∼ < appear to be occurring. We commend these stars to the attention of individuals having extensiveaccess to telescopes of moderate aperture.Our sample of 164 RRL candidates contains 95 good stars (see Table 2) including 24that apparently either show the Blazhko effect or pulsate in multiple modes. Among thegood RRLs we recognize 14 first overtone (FO) and 81 fundamental-mode (FU) pulsatorson the basis of their period/amplitude distribution in the Bailey diagram (see Fig. 3). Theremaining 69 RRL candidates with fair and poor light-curve fits are provisionally dividedinto 14 FO (12 “B” and 2 “C”) and 55 FU (35 “B” and 20 “C”) pulsators. The pulsation 9 –properties of the RRL sample will be discussed in more detail in the next section.The sample of AC/spCs is quite abundant, populating the whole theoretical IS, andconsists of 55 objects (see Table 3), including two with uncertain or very uncertain periods(V212 and V241). Our sample also includes one possible Blazhko or multiple-mode Anoma-lous Cepheid, V129. The evolutionary classification of these stars has been discussed in acompanion paper (Fiorentino et al. 2012) and will not be repeated here, but in Section 5 wewill place some constraints on their masses and mode classification.The spatial locations of all our variables are shown in Fig. 4. The ( x, y ) coordi-nates in this diagram are referred to an origin (0 ,
0) that we have arbitrarily placed at10 h m s .21 +12 ◦ ′ ′′ .0 (close to the center of an early subset of our images). The y -axis is the great circle x = 0 and is equal to the great circle α = 10 h m s .21, and the x -axis is tangent to δ = +12 ◦ ′ ′′ .0 at the origin. The photocenter of Leo I has been esti-mated to lie at ( α, δ ) = (152 . , . h m s .00 +12 ◦ ′ ′′ .0) according toMateo (1998); we ourselves estimate the center of Leo I to lie at 10 h m s .13 +12 ◦ ′ ′′ .2[( x, y ) = (+160 . , − . ′′ in each coordinate. This esti-mate has been derived by determining the median x and y coordinates of probable memberstars lying within a range of distances from the center, from 100 ′′ to 850 ′′ ; the consistency ofthese values provides our estimate of the uncertainty (see Stetson et al. 1998a, §
4. RR Lyrae stars
The Bailey diagrams are shown in Fig. 3 for the B (top) and V (bottom) photometricbands. As expected, the RRLs separate very cleanly into two different groups, the longer-period FU pulsators (log P ∼ > –0.35, P ∼ > ∼ < –0.35: RRc, or RR1). These diagrams arevery often used to give an indication of the Oosterhoff class. For this reason, in the V -band (bottom) panel we have also plotted the curves (Cacciari et al. 2005) representativeof the Oosterhoff dichotomy shown by RRab stars in Galactic GCs. We note that most ofthe RRab stars seem to follow the OoI line, reaching a maximum A V amplitude of about 10 –1.2 mag. However, as discussed in (Fiorentino et al. 2012), based on the mean period ofthe RRab sample we have classified Leo I as Oo-intermediate, since h Pab i = 0.596 ± σ =0.05), based on only those stars with good light-curve fits. This result does not changewhen we include the RRab variables with fair or poor periods or when we add the RRc starswith fundamentalized periods (Coppola et al. 2013).We have also calculated the ratio between the number of RRc and the total number ofRRLs, which is considered another indicator of Oosterhoff class, N c / ( N c + N ab ) = 28 / ≈ .
17. This value of the ratio includes the variable candidates with fair and poor light-curvefits; a slightly lower value, 14 / ≈ .
15, is obtained when it is based on only the good RRLs.We have a slight preference for the higher value, however, because the very classification“good,”, “fair,” and “poor” disproportionately assigns the lower quality classes to the c-typevariables with their smaller amplitudes. Although those same smaller amplitudes probablymean that the c-type variables are more affected by incompleteness even if we ignore thequality classes, these very small values of the c-to-total ratio still suggest that it would behard for us to have missed enough c-type variables to push Leo I out of the the OoI class.This conclusion is in line with what has been observed in other nearby dSph galaxies (seesection 6 for details).We take advantage of the good sampling of our light curves, in particular in the B and V bands, to estimate the mean amplitude ratios h A V /A B i = 0 . ± .
002 mag ( σ = 0.20)and h A R /A V i = 0 . ± .
004 mag ( σ = 0.24), based on 95 and 63 RRLs respectively (not allstars with “A”-quality data in B and V have good data in R ). This A V /A B value is in fullagreement with the value discussed in Di Criscienzo et al. (2011) based on 130 RRLs takenfrom nine GGCs. In this section we use the RRL pulsation properties to constrain the distance to Leo I. Inparticular, we use two independent methods: a linear formulation of the M V versus [Fe/H]relation (Bono et al. (2003); see also Cacciari & Clementini (2003)) and the First OvertoneBlue Edge (FOBE) method (Caputo et al. 2000, and references therein).For the M V versus [Fe/H] relation we consider only the 81 RRab-type variables withlight curves of quality class “A” (“good”), and we compare their measured properties to thefollowing theoretical formulations: M V ( RR ) = (0 . ± . . ± . , for [Fe/H] ≤ − . M V ( RR ) = (0 . ± . . ± . , for [Fe/H] > –1.6 dex(Bono et al. 2003). In order to use these formulae we must assume a metal abundancefor Leo I. Recent measurements of the galaxy’s metallicity based on a large ( ∼ ∼ –1.43 ( σ = 0.33 dex; Kirby et al. 2011). However, it isworth mentioning that Leo I seems to have a broad metallicity distribution, with iron abun-dances ranging from ∼ –2.15 to ∼ –1. RRLs being among the oldest and presumably themetal-poorest stars in a stellar system, we decided to derive two different distance modulicorresponding to both the average Leo I metallicity and the metal poor extreme. We found h V i = 22 . ± .
001 mag ( σ = 0 .
11) for the RRL sample, and assuming E(
B–V ) = 0.02mag , the two distance moduli turn out to be µ = 22 . ± .
08 and 22 . ± .
07 mag for[Fe/H] = –1.43 and –2.15, respectively. These values are in good agreement with the distancemodulus derived by Bellazzini et al. (2004) using the TRGB method : µ =22.02 ± +16 − kpc.Now we use the technique extensively discussed in Caputo et al. (2000), which is agraphical method based on the predicted period-luminosity (PL) relation for pulsators lo-cated along the FOBE. It seems quite robust for clusters with significant numbers of RRcvariables and is thus applicable to Leo I. The distance modulus is derived by matching theobserved distribution of RRc variables to the following theoretical relation: M V (FOBE) = − . ± . − .
255 log P (FOBE) − .
259 log(
M/M ⊙ ) + 0 .
058 log Z. Here, we assume the same two metallicity values for the RRLs in Leo I and adopt M =0.7M ⊙ from the evolutionary HB models for RRc variables, with an uncertainty of the orderof 4% (Bono et al. 2003). The comparison between the observed RRLs and the theoreticalrelation is shown in Fig. 5. Given the period-luminosity distribution of our RRc sampleand the high uncertainty in the FOBE determination, we decided to use two possible FOBEevaluations, the first defined by the “good” RRc V85 and the other by V80, V141 and V214.Thus, the FOBE method returns a range of possible distance moduli, µ =21.96 (V85) to22.14 (V80,V141,V214) assuming [Fe/H]=–2.15 and E( B–V ) = 0.02 mag; these decrease by The adopted value is taken from Burstein & Heiles (1984) and is equal to a mean of different estimatesused in the literature (Gallart et al. 1999; Bellazzini et al. 2004; Held et al. 2001). We note that the use ofthe recent and slightly higher reddening value provided by Schlafly & Finkbeiner (2011) (0.031 mag) wouldcause a decrease of ∼ In this context it is worth noting that the TRGB method has a minimal dependence on metallicity inthe relevant range, − . < [Fe/H] < − .
2; (see, for example, Tammann et al. 2008).
12 –0.04 mag when using the metal-rich stellar content (–1.43 dex). It is worth noting that thedistance evaluation based on the FOBE defined by V85 gives a distance determination thatis also in very good agreement with that given by the TRGB method (Bellazzini et al. 2004).
5. Anomalous and short-period Cepheids
In a previous paper (Fiorentino et al. 2012), we discussed the evolutionary classificationof the Cepheid sample we have detected in Leo I. On the basis of a comparison with a setof evolutionary tracks at different metallicities ([Fe/H] ranging from –1.8 to –1.0) we haveconcluded that we are dealing with an unprecedented mix of ACs and spCs. In this section,we discuss the pulsation properties of this unique sample of variable stars.In Fig. 6 we show their distribution in the amplitude versus period diagrams in the B (top) and V bands (bottom). It seems clear that the sample can be divided into high- andlow-amplitude subsamples separated at A B ∼ . A V ∼ .
8, with the low-amplitudestars dominating for periods less than one day and the high-amplitude stars dominating atlonger periods. Similar behavior seen in other classes of pulsating variables, e.g., RRLs,usually identifies a mode separation between the first overtone (lower amplitudes, shorterperiods) and fundamental (higher amplitudes, longer periods) modes of pulsation. However,for ACs the mode classification is far from trivial and cannot be easily linked to eitherthe period-amplitude distribution or the morphology of the light curves, as discussed inMarconi et al. (2004).We note that observational uncertainties probably do not contribute much to the scatterof points in the amplitude versus period diagrams. For the 53 out of 55 candidates to whichwe have assigned quality class “A,” we feel that the amplitude uncertainties are probably notworse than 0.1 mag and the period uncertainties are probably not larger than 1% ( σ (log P ) ∼ . V , B–V ) Wesenheit magnitude versus period plane. We have chosen this color combinationbecause the available B and V observations best sample the light curves, resulting in themost accurate mean magnitudes. Uncertainties of a few × B and V magnitudes for the stars lead to uncertainties ∼ < ⊙ (Marconi et al. 2004; 13 –Fiorentino et al. 2006) and metallicities Z from 0.0001 to 0.0004, as well as those for Pop-ulation II Cepheids (Di Criscienzo et al. 2004). We have assumed µ =22.11 ± B–V )=0.02 mag, values that were used by Fiorentino et al. (2012) to classifyour sample of variable stars and that agree quite well with what is found using RRLs (seeprevious section). Our sample occupies the general region where ACs are expected withoutdefining any tight sequence in this diagram. There are also no distinct subdistributions form-ing different sequences of AC/spCs such as are typically observed in other galaxies, whereabundant samples of ACs (Fiorentino & Monelli 2012) and/or spCs (Bernard et al. 2013)show that both FU and FO pulsators are detected.To investigate this behavior, we again use the ap-proach discussed in Fiorentino & Monelli (2012), one that allows us to identify the modeclassification within the sample of ACs in the LMC released by OGLE III (Soszy´nski et al.2008). This method, detailed in Caputo et al. (2004), returns simultaneously the mode andthe pulsation mass for each individual AC. It is based on the use of a theoretically predictedmass- and luminosity-dependent relation between period and V -band amplitude (MPLA)for FU mode pulsators, which is not followed by pulsators in higher modes. Coupling thisrelation with the mass-dependent period-luminosity-color (MPLC) relation that exists forboth modes, we will assign the FU pulsation mode only when the two masses agree within1 σ . We use the following relations:log P F + 0 . M V + 0 .
77 log M = 0 . − . A V , predicts the visual amplitude from the period, luminosity and mass for fundamental pul-sators, and M V = − . − .
85 log P + 3 . B–V ) − .
88 log M (FU) ,M V = − . − .
90 log P + 3 . B–V ) − .
82 log M (FO)predict the luminosity from the period, color and mass for FU and FO pulsators. These havebeen derived in (Marconi et al. 2004), and can be inverted to yieldlog[M/M ⊙ (MPLA-FU) ] = (cid:0) . − . A V − log P − . M V (cid:1) / . ⊙ (MPLC-FU) ] = − (cid:0) M V + 1 .
56 + 2 .
85 log P − . B–V ) (cid:1) / . ⊙ (MPLC-FO) ] = − (cid:0) M V + 1 .
92 + 2 .
90 log P − . B–V ) (cid:1) / . . Our application of these equations produces the masses and the classifications given in Ta-ble 3 and summarized in Fig. 8. In this figure we show the mass distribution (gray histogram)predicted for our sample based on the pulsation relations. This is to be compared with thevery peaked mass distribution shown by ACs observed in the LMC. This confirms our in-terpretation that a continuous spread in masses in Leo I causes the dispersion shown in the 14 –Wesenheit plane (see Fig. 7), whereas in the LMC you can perceive two distinct sequencesin the same plane (Soszy´nski et al. 2008; Fiorentino & Monelli 2012).Three AC/spCs in Leo I show unreasonably high masses, larger than 5 M ⊙ (indicatedwith crosses in Figs. 6 and 7); in Fig. 7 these stars are the ones deviating the most fromthe global distribution in the Wesenheit plane due to their very short periods comparedto other stars with comparable luminosities and colors. In particular, the star with thehighest indicated mass, V241, belongs to our “very uncertain” class (C) and shows a veryred color ( B–V ≥
6. Comparing Leo I with classical spheroidal and ultra-faint dwarf galaxies
In Figs. 9 and 10 we show the amplitude versus period diagrams (left) and theperiod histograms (right) for RRL and AC/spC stars in six dSphs ordered by increas-ing baryonic mass (McConnachie 2012), namely Draco (Kinemuchi et al. 2008), Carina(Coppola et al. 2013), Tucana (Bernard et al. 2009), Sculptor (Kaluzny et al. 1995), Cetus(Bernard et al. 2009; Monelli et al. 2012) and Leo I (this paper). In the top panel of eachfigure we have also shown all the RRLs and AC/spCs observed in the eleven UFDs surveyedfor variability (Dall’Ora et al. 2006; Kuehn et al. 2008; Moretti et al. 2009; Musella et al.2009; Dall’Ora et al. 2012; Musella et al. 2012; Clementini et al. 2012; Garofalo et al. 2013;Boettcher et al. 2013). For comparison, in Fig. 9 we also show the two curves that trace theOosterhoff dichotomy observed in Galactic GCs as defined by Cacciari et al. (2005). Lyingin general between the two Oosterhoff curves, the ab-type RRL distributions for the classicaldwarfs in the Bailey diagram are very similar to each other. The only (slight) exceptionsare Cetus and Carina, which show a different slope in the period versus amplitude planeas discussed in Bernard et al. (2009) and Monelli et al. (2012). For most dSphs, the perioddistributions of RRab stars (right panels of Fig. 9) are quite peaked around their mean valuewith two exceptions, viz. Sculptor and Tucana. In Table 6 we have listed the mean propertiesof the RRLs according to the catalogs we have adopted; the mean RRab and RRc periodsare the same within 1 σ . In the last column we have also listed the total masses and meanmetallicities of these galaxies according to the references used in McConnachie (2012). Themean periods of RRLs in UFD galaxies seem to generally follow the same behavior as inclassical dwarfs, occupying the same general location in the Bailey diagram with a slightly 15 –higher mean period for the RRab stars, very similar to that of Carina (Dall’Ora et al. 2003;Coppola et al. 2013). This suggests that, at least in terms of their RRL properties, there isnot a significant difference between classical and ultra-faint dwarf galaxies.On the basis of their similar general properties, we decided to build up a single largesample of well studied RRLs in those dwarf galaxies where the variability surveys do not seemto suffer from strong completeness problems, thus very likely describing the RRL propertiesof these stellar systems rather well. This initial sample contains 1,726 objects (with 1,299ab-type variables); we included Cetus and Carina because their inclusion does not bias themean properties of the sample nor change our final conclusions.As discussed in Fiorentino et al. (2012), different conclusions result from an inspectionof Fig. 10, where it is clear that the AC/spCs have different period distributions and Baileydiagrams. In particular, we note that the period distribution seems to move toward longerperiods when the baryonic total mass of the galaxy increases (see Table 6). This is easilyunderstood when one considers that the high masses typical of ACs (Fiorentino et al. 2012,1.2 ∼ < M/M ⊙ ∼ < ∼ <
7. Comparing dwarf galaxies with globular clusters and the halo field
Usually, the average properties of RRLs in individual dSphs and UFDs—such as themean periods of both RRab and RRc stars—are compared to those observed in GCs asrepresentative of the Galactic halo (Catelan 2009; Clementini 2010, and references therein).However, given their different total stellar masses, GCs typically host RRL populations oforder a factor 10 smaller than those observed in dwarf galaxies, making a proper statisticalcomparison frequently difficult. Moreover, in the last ten years or so, GCs have been demon-strated to be quite complex stellar systems despite what we previously believed in termsof their chemical enrichment histories (Gratton et al. 2004; Piotto et al. 2005; Milone et al.2012; Monelli et al. 2013), and they may not fairly represent the properties of halo field stars.Even though the net effect of their complex histories on their global average chemical abun-dances (and therefore on the pulsation properties of their RRLs) may be negligible, this hasnot yet been demonstrated. Finally, although it seems trivial that the statistical meaningof averaged properties is, by definition, more significant when a large sample of objects isconsidered, we must remember that small number statistics become quite relevant in those 16 –galaxies with very few confirmed member RRLs, as is the case for most UFDs (Moretti et al.2009; Musella et al. 2009; Dall’Ora et al. 2012; Clementini et al. 2012; Boettcher et al. 2013)where even the term “average” approaches meaninglessness.The justification for our assembling a large sample of RRLs in classical and ultra-faintdSphs is driven in large part by the opportunity to compare directly , for the first time, theirperiod and amplitude distributions with those of a huge RRL sample representing the Galac-tic halo. RRLs are ancient stars, older than 10 Gyr. Their younger selves were born duringthe early millennia of the Universe, and their testimony can provide important informationabout chemical evolution during the early stages of the assembly of the Galactic halo. More-over, they are their own robust distance indicators that might also reveal details of the halo’sspatial structure (Layden 1994; Kinemuchi et al. 2006; Drake et al. 2013; Zinn et al. 2014).With this goal in mind, we have collected several catalogs—mainly on the basis of theavailability of robust periods and reliable V -band amplitudes—that, taken together, provideabout 14,000 RRab stars spanning distances from ∼ ∼
80 kpc, namely the QUEST(Vivas et al. 2004; Zinn et al. 2014), NSVS (Wo´zniak et al. 2004), ASAS (Szczygie l et al.2009) and CATALINA surveys (Drake et al. 2013). In cases where a star appeared in multiplesurveys, we have retained only those data originating in the most recent and complete studyavailable.We have computed the three-dimensional Galactocentric position of each individualRRL, first converting (RA,Dec) to (l,b) coordinates and then assuming for the stars in thevery inner part of the Halo (d G ∼ < > –1.6 dex (Bono et al. 2003), assuming a distance to the Galactic Center of 7.94 kpc(Eisenhauer et al. 2003; Groenewegen et al. 2008; Matsunaga et al. 2011).In our final catalog, we have not included the samples from Miceli et al. (2008) andSesar et al. (2013) because amplitudes on their filterless magnitude system must be trans-formed into V (Landolt) amplitudes using a metallicity- and temperature-dependent scalefactor that could affect the stars’ apparent distribution in the Bailey diagram. On the otherhand, we do include the CATALINA sample, while acknowledging that it has a bias in am-plitudes; we retain only RRLs with amplitudes A V larger than 0.4 mag. However, this is thelargest, deepest, and most homogeneous catalog at our disposal, and its known bias turnsout to have a negligible effect on the following discussion. Finally, we use only ab-type RRLsbecause they are less affected by both time-sampling and completeness problems. Although 17 –not 100% complete, this final huge catalog will allow us to make the most comprehensiveanalysis of the Galactic halo using field RRLs possible so far.In order to draw the most complete global view of the globular clusters belonging tothe Galactic halo, we have decided to exclude all the RRLs listed in the updated catalog ofClement et al. (2001) that are observed in GCs belonging to the Galactic bulge, i.e., thosehaving both Galactocentric distance d G ≤ ≤ V , including amplitudes, is available. The full sample consists of 1,617 RRLs (1,054 ab-type)residing in 35 GCs (see Table 7).In Fig. 11 we show the distributions of RRab stars in period-amplitude diagrams (left)and period histograms (right) for the full set of catalogs we have collected. The first fourpanels are devoted to the Galactic halo divided into four non-overlapping regions: 1) theinner halo a, consisting of stars with d G ∼ < ∼ < d G ∼ <
14 kpc (panel b); 3) the outer halo a, stars with 14 ∼ < d G ∼ <
30 kpc (panelc); 4) the outer halo b, stars with d G ∼ >
30 kpc (panel d). The boundary between innerand outer halo (d G ∼
14 kpc) has been chosen according to Kinman et al. (2012), whereasthe other reference distances are arbitrary. In the last two panels, we have plotted the fulldistributions obtained for GCs (panel e) and UFDs+dSphs (panel f). As in Fig. 9, we haveoverplotted the Oosterhoff curves given by Cacciari et al. (2005). The average of their twocurves (blue dashed curve in Fig. 11) allows us to roughly separate the two populations ofOoI (large dots) and OoII (small dots) RRLs. We have computed the ratio between OoIand the total number of RRLs, as given in the last column of Table 7. These numbersare very similar to each other and they are compatible with a global OoI classification, inagreement with recent and independent results from the ROTSE (Miceli et al. 2008) andLINEAR (Sesar et al. 2013) surveys.We have also used these data to compute the mean periods of the RRLs withlog P ≥ –0.35 (see Table 7), and find that h Pab i is the same within the standard devia-tions of the distributions. Again, the mean properties of the different RRab samples seemvirtually indistinguishable from this statistical point of view.To compare these distributions in more rigorous detail, we have performed two moresophisticated statistical analyses. The first is a chi-squared test on the histograms smoothedwith a Gaussian filter (see blue solid curves in Fig. 11), and the second is a Kolmogorov- 18 –Smirnov test applied to the cumulative period distributions (see Fig. 12). For the purposesof these last two tests, only, we take advantage of the large sample of RRLs in the LargeMagellanic Cloud that has been provided by the OGLE project (17,693 stars; Soszy´nski et al.2009). Although the RRLs in the LMC have been surveyed in the I band by OGLE, wedecided to attempt to match the amplitude bias of the CATALINA sample by using onlythose stars with A I ≥ A V /A I = 1 .
58 as found byDi Criscienzo et al. (2011). These two statistical tests returned very similar results, whichare listed in Table 8 and can be summarized as follows: • Inner halo a and b have a likelihood of 30% of being drawn from the same parentpopulations. This likelihood increases dramatically to 90% when only the CATALINAsurvey is considered. As shown in Fig. 11, these two regions show broad period distri-butions including the longest-period RRLs. • Outer halo a and b do not share a common period distribution, with each other orwith the inner halo. Each component of the outer halo seems to have a more peakeddistribution than that of the inner halo. In particular, outer halo b seems to show arelative deficiency of RRab stars with high amplitudes and short periods (hereinafter,“HASP”). This can not be attributed to a completeness effect because large amplitudesare the easiest ones to detect. • The GC period distribution mildly resembles that of the RRLs in the inner halo (within ∼
14 kpc, a+b). This is reasonable when we consider that 23 out of the 35 GCsconsidered here belong to this region. We note that this result must be viewed aspreliminary because: 1) the GC period distribution strongly depends on the choiceof GCs included in the full catalog; 2) GCs have their own space motions that mayeventually tell us more about the halo region to which they belong, rather than wherethey happen to be currently located. • The UFD+dSph distribution is quite different from the others. It is very peaked aroundthe mean period (see Table 7). In particular, this distribution is quite lacking in HASPRRab stars compared to the others: none of the RRLs in dwarfs with high amplitude( A V ≥ .
0) reaches periods P ∼ < .
48 d (log P ∼ < –0.32). Fundamental-mode RR For completeness, we should mention that there are two RR Lyrae variables in the HASP region of theBailey diagram for the dSph galaxy Cetus: V11 and V173 in Bernard et al. (2009). However, both thesestars are peculiar: they are subluminous and have much shorter periods than any other RRab stars in thisgalaxy, and moreover the light curve for V173 is of very poor quality. For these reasons we have not includedthese two stars in the UFD+dSph histogram in Fig. 11, but they are included in the Cetus panel of Fig. 9.
19 –Lyraes with large amplitudes and periods less than 0.48 d are quite common both inglobular clusters and in the halo field. Even the relative deficiency of HASP RRabstars in outer halo b already noted above is not as complete as in the UFD+dSphsample. It is true that at the other extreme, UFDs are seen to help in filling out thelong-period tail ( cf.
Fig. 9a) but—within the limited statistics available—they do notappear to contribute to the short-period tail. • We find a very low formal likelihood that the LMC sample matches any of the others(see Table 8), as is also true for classical dSphs and UFDs. We note here, though, thatthis result may be affected by the very high temporal and photometric completenessof the OGLE sample for the LMC. In other words, when huge samples are available(as they are here for the LMC and the Milky Way halo), any minor difference inselection biases achieves very high statistical significance and can be mistaken for areal physical difference between the samples. A fair comparison with the LMC sampleneeds much more complete RRL surveys of the Galactic halo. The same considerationis relevant to the recent result discussed by Zinn et al. (2014) where, on the basis ofthe OGLE (Soszy´nski et al. 2009) and their own QUEST RRL samples, the authorsdeclined to exclude the possibility that the smooth Galactic halo may have formed froma combination of LMC- and Sagittarius-like galaxies. We believe—but are unable toprove—that because our comparisons between the halo, the GCs, and UFD+dSphs arebased on samples closer in size, and probably with selection effects more nearly similarto each other than to the OGLE LMC sample, our conclusions regarding the statisticalsignificance of their differences are probably more reliable than similar conclusionsrelating to the LMC.
8. Discussion: Building up the Galactic Halo with dwarf galaxies
Perhaps it is worthwhile to remind the reader that the dwarf galaxies in the Milky Wayhalo contain stellar populations with some range of ages and metallicities. The classicalglobular clusters as a class exhibit a broad range of metallicities ( − . < [Fe/H] < . ∼ < –1 or so, but some—including Leo I, Carina, and Fornax, for example—havevery broad ranges of age, much greater than the range found among the GCs.However, it must also be remembered that the range of age and metal abundance thatis conducive to the formation of RR Lyraes is more restricted: the most metal-rich globular 20 –clusters contain no RR Lyraes, and the most metal-poor contain few, while moderately metal-poor clusters contain the most. Populations younger than ∼
10 Gyr probably cannot produceRR Lyraes by the normal evolution of single stars. The range of metallicities where theglobular clusters and the field halo are proficient at producing RR Lyraes, − . ∼ < [Fe/H] ∼ < − .
0, is very well represented in the dwarf galaxies, and, as we have just mentioned, stars inthe dwarf galaxies have—if anything—a broader range of age than those in the globularclusters and the field halo. That being the case, it is remarkable that there exists a class ofRR Lyrae star that is present in the globular clusters and the field halo and absent in thedwarf galaxies. One would think that if there were some odd corner of age-metallicity spacethat is capable of making a particular class of RR Lyrae stars, but is not populated amongthe globular clusters, that niche would be more likely to be occupied in the dwarf galaxies: ifanything, the dwarf galaxies should have more varieties of RR Lyrae stars than the globularclusters, not fewer .One possible exception to this rule might be if the HASP RR Lyrae stars were uniquelythe progeny of stellar collisions or merged gravitational-capture binaries; such stars couldthen be present in the globular clusters and not in the much sparser dwarf galaxies. Butif this is the case, then why are these “special” RR Lyraes also found in the field halo,in roughly the same relative numbers as in the globular clusters? To be consistent, thisexplanation would suggest that a major fraction of the field halo stars came from globularclusters that were disrupted recently , i.e., only after the globular clusters had had enoughtime to produce the usual numbers of stellar collisions and gravitational-capture binaries.But even this seems highly unlikely, since halo field red giants do not show the light-element(C, N, O, Na, Mg, Al, etc.) correlations and anti-correlations found among the globularcluster giants.In conclusion, and for the sake of argument, let us assume that in the early days ofthe Galactic halo, star formation occurred in one fundamentally indistinguishable familyof structures, ranging in mass, that all underwent a common process of internal chemicalevolution. During this time the young main sequence stars that grew into today’s RR Lyraeswere born. During subsequent dynamical evolution most of these structures dissolved or weredisrupted, releasing the globular clusters and field stars of today’s Galactic halo. At the sametime a few of those structures—by random chance—survived, continued evolving internally,and became the classical dSph and UFD galaxies that we see today.It seems that on the basis of the available RRL samples we can say that this scenariois not supported. Neither the inner (d G ∼ <
14 kpc) nor the outer (d G ∼ >
14 kpc) regionsof the Galactic halo, nor any linear combination of the two, can be formed simply fromthe progenitors of today’s classical dSph or UFD galaxies. If present in sufficient numbers, 21 –structures like those that became today’s UFDs may have contributed to the long-period tailof the Galactic halo RRL period distribution, but—barring a major statistical fluke—neitherthey nor the dSph progenitors appear to have been capable of producing the kind of mainsequence stars that became the HASP RRab variables now found in significant numbersthroughout the Galactic halo and in the globular clusters.Unless it can be shown that a protracted residence within the environment of a dSphor UFD galaxy can somehow alter the internal structure and evolution of an isolated starbetween the main sequence and RR Lyrae phases of its life, it seems that the simple model—where the dSphs and UFDs are surviving examples of representative
9. APPENDIX: Previously identified stars
Leo I was previously surveyed for bright variable stars by Hodge & Wright (1978), whoidentified candidate Anomalous Cepheids and RR Lyraes, and by Menzies et al. (2002), whoidentified candidate long-period variables.Hodge & Wright did not publish coordinates for their stars, but they did provide afinding chart. We extracted a JPEG version of their finding chart from the on-line electronicedition of their article and converted it to FITS format. Using a slightly modernized versionof the Stetson (1979) software for astrometry and photometry from digitized photographicplates, we measured the positions of stars within the Hodge & Wright finding chart. Wethen used quadratic equations to transform those positions to the coordinate system of ourLeo I catalog. We were able to match 355 stars within a tolerance of 4 ′′ (the r.m.s. residualswere 0 ′′ .45 in right ascension and 0 ′′ .37 in declination). We were then able to geometricallytransform their finding chart and overlay it on our stacked digital image of Leo I. Overall,we feel we are able to recover positions from the Hodge & Wright finding chart with aprecision better than 0 ′′ .5. We considered stars within several arcseconds of each of thepositions indicated in the published chart, basing our identification on (in order of decreasing 23 –importance) (1) evidence of variability in our data, (2) relative proximity to the indicatedposition, and (3) similarity of the mean B -band magnitude. We summarize our resultingcross-identifications in Tables 2 (RR Lyrae stars), 3 (AC/spC stars) and 9 (stars whosevariability we are unable to confirm).Notes on individual stars: • The tick marks for HW1 in their chart enclose nothing but blank sky in our image.However, there are two of our variable candidates within 5 ′′ of this position: V66 is acandidate long-period variable with h B i ≈ . ′′ .2 from the predicted position,and V59 is an AC/spC with P = 1 . h B i ≈ . ′′ .8 from the predictedposition. Despite its greater offset, the latter star more closely resembles HW1 and weaccept it as the match. • HW2 has no detectable star in our image within the tick marks on the finding chart.There is a spot of fuzz most likely representing a galaxy about 1 ′′ southwest of theindicated position, but this is probably too faint to have been detected in Hodge &Wright’s photographic material. One of our AC/spC candidates does lie almost 6 ′′ ( > σ ) from the indicated position. Since it has approximately suitable period andmagnitude, we provisionally make this identification despite the large offset. • HW20 has at least eleven stars lying closer to the indicated position than our adoptedmatch, but the one we have chosen is clearly the most suitable on the basis of itsmean B -band magnitude, and furthermore its Welch/Stetson variability index is notfar below the detection threshold that we had imposed. This star may be variable, butwe are not confident enough to claim definite confirmation.Menzies et al. (2002) did not publish a finding chart for their five stars in Leo I. Theydid publish a table of equatorial coordinates for five luminous red stars in the galaxy field,but unfortunately we have not been able to find any relationship whatsoever between thosecoordinates and reality: we have marked their coordinates on our digital image of Leo I andwe have also marked all the stars that are both luminous and red in our catalog, and we arenot able to recognize any correspondence between the two patterns. For example, their starsA and B constitute a pair separated by 29 ′′ .7, slightly inclined with respect to the east-westdirection. We have examined our list of luminous red stars for pairs separated by ∼ ′′ ,slightly inclined to the east-west direction, and do find a few. But when we accept any ofthose provisional matches none of the other three stars from Menzies et al. coincides withanything interesting. 24 –Azzopardi et al. (1986) published a finding chart identifing several luminous, red starsin Leo I, and we give cross-identifications with our catalog in Table 10. REFERENCES
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This preprint was prepared with the AAS L A TEX macros v5.2.
30 –Table 1. Leo I Observations
Run ID Dates Telescope Camera/Detector
U B V R I – – 4 – –15 jka 1995 03 24 INT 2.5m TEK3 – – 1 – –16 apr95 1995 04 29-05 01 Steward Bok 2.3m 12x8bok – – 26 – –17 mdm(b) 1995 05 05 Hiltner 2.4m Charlotte Tek 1024 – – 4 – –18 bond2 1996 03 13 KPNO 4.0m t2kb 1 1 1 – 119 int 1998 06 24 INT 2.5m WFC EEV42 – 1 1 – 1 ×
420 wfi3 1999 03 17 ESO 2.2m WFI – – 13 – 13 ×
821 emmi 1999 04 14-15 ESO NTT 3.6m EMMI – 6 6 – 122 fors9912 1999 12 02 ESO VLT 8.0m FORS1 – – 3 – 323 wfi4 2000 04 21-26 ESO/MP 2.2m WFI – 13 31 – 2 ×
824 wfi 2000 04 22 ESO/MP 2.2m WFI – 4 10 6 1 ×
825 tng 2001 01 24-26 TNG 3.6m LRS Loral – 30 49 – –26 bellaz 2001 03 19-22 TNG 3.6m LRS Loral – – 8 – 827 suba 2001 03 20 Subaru 8.2m SuprimeCam – – – 7 – ×
828 suba2 2001 03 20-21 Subaru 8.2m SuprimeCam – – 28 16 – ×
829 fors0204 2002 04 13 ESO VLT 8.0m FORS2 MIT/LL mosaic – – 4 – 4 ×
230 abi36 2002 11 13-15 KPNO 0.9m s2kb – 3 3 3 331 lbt(a) 2006 11 21-25 LBT 8.4m LBC 43 20 – – – ×
432 lbt(b) 2006 12 16 LBT 8.4m LBC – – 5 – – ×
31 –Table 2. RR Lyrae candidates: periods, photometric parameters and positions.
ID Var period h B i h V i h R i h I i A B A V A R A I Q RA Dectype d mag mag mag mag mag mag mag mag hh mm ss dd mm ssV1* RRab 0.515750 23.72 22.79 21.60 — — 1.62 — — C 10 07 13.92 +12 21 50.8V2* RRab 0.627407 22.99 22.53 22.14 — 1.14 1.17 0.96 — A 10 07 18.89 +12 13 10.2V3* RRab 0.554130 23.04 22.61 22.31 — 0.84 1.03 — — B 10 07 28.27 +12 14 40.4V4* RRab 0.638565 22.88 22.53 22.50 — 1.01 0.93 — — B 10 07 33.35 +12 21 56.1V5* RRab 0.593564 23.25 22.75 22.46 22.16 1.10 0.95 0.60 — B 10 07 36.68 +12 18 49.1V6* RRab 0.698670 23.13 22.61 22.28 21.74 1.00 0.88 0.75 — C 10 07 38.33 +12 12 40.3V7* RRab 0.6024018 23.04 22.75 22.57 22.25 0.96 0.79 0.23 — A 10 07 40.62 +12 18 55.6V8* RRab 0.537169 23.15 22.59 22.20 22.15 1.14 0.94 — — B 10 07 44.43 +12 13 07.5V9 RRab 0.574817 22.81 22.69 22.65 22.34 1.17 1.22 0.36 — A 10 07 44.44 +12 17 17.4V10 RRab 0.681396 23.00 22.61 22.41 — 0.45 0.37 0.32 — B 10 07 46.57 +12 20 15.0V11 RRab 0.515261 23.05 22.64 22.44 22.45 1.24 1.18 — — B 10 07 48.43 +12 18 56.8V13 RRab 0.636654 23.13 22.64 22.47 — 0.58 0.46 0.40 — A 10 07 52.04 +12 15 34.6V14 RRab 0.708036 23.15 22.90 22.70 22.44 0.84 0.60 0.18 — C 10 07 55.61 +12 20 53.9V15* RRab 0.630248 23.10 22.72 22.56 — 0.79 0.68 0.38 — A 10 07 55.99 +12 14 13.6V16 RRab 0.601348 23.04 22.78 22.35 — 1.00 0.90 0.68 — A 10 07 56.41 +12 19 47.1V17 RRc 0.3131235 23.00 22.74 22.36 — 0.72 0.49 — — A 10 07 56.48 +12 13 54.4V18 RRab 0.628029 23.02 22.70 22.41 — 1.05 0.79 0.45 — A 10 07 56.55 +12 18 16.1V19 RRab 0.607947 23.08 22.66 22.39 — 0.90 0.83 0.34 — A 10 07 58.54 +12 18 00.0V21 RRab 0.659662 22.88 22.67 22.56 22.04 0.98 0.85 0.37 — B 10 08 00.69 +12 24 54.3V22 RRab 0.557603 23.00 22.70 22.50 22.19 1.03 0.97 0.42 — B 10 08 01.08 +12 17 54.0V23 RRc 0.2737468 23.05 22.69 22.54 22.16 0.45 0.37 0.32 — B 10 08 01.44 +12 20 09.2V24* RRab 0.607729 23.20 22.78 22.57 22.25 0.66 0.76 0.30 — A 10 08 02.45 +12 18 07.4V25 RRab 0.5615871 23.11 22.68 22.54 — 0.84 1.03 — — A 10 08 03.48 +12 19 58.0V26 RRab 0.541359 23.02 22.64 — 9.999 1.16 1.13 — — B 10 08 04.08 +12 15 57.6V28 RRc 0.385380 22.94 22.67 22.49 22.12 0.72 0.44 0.38 — B 10 08 05.74 +12 19 24.9V31 RRab 0.6026533 23.07 22.77 22.49 — 1.19 0.96 0.70 — A 10 08 07.07 +12 21 32.5V32 RRab 0.580355 22.97 22.66 22.44 22.27 1.05 0.88 0.51 — A 10 08 07.68 +12 18 50.1V35 RRab 0.626002 23.08 22.73 22.57 — 0.57 0.64 0.22 — A 10 08 08.25 +12 13 48.6V36 RRab 0.558055 23.07 22.70 22.75 22.23 0.72 0.85 0.14 — A 10 08 08.27 +12 15 13.0V37* RRab 0.5865854 22.96 22.74 22.66 22.10 1.04 1.08 0.78 — B 10 08 08.52 +12 13 15.9V38 RRab 0.5694029 22.94 22.66 22.40 21.92 1.05 0.90 0.38 — A 10 08 08.82 +12 21 23.8V42 RRab 0.590415 22.89 22.59 22.44 21.91 0.93 0.84 0.26 — A 10 08 10.01 +12 16 16.7V43 RRab 0.518913 23.11 22.65 22.01 21.98 0.70 0.60 0.50 — B 10 08 10.30 +12 15 14.9V44 RRab 0.5368825 22.82 22.53 22.50 21.84 1.32 0.96 — — A 10 08 10.75 +12 18 28.9V47 RRc 0.3741098 22.89 22.59 22.42 22.22 0.43 0.40 0.31 — A 10 08 11.42 +12 17 55.8V49 RRab 0.6106710 23.11 22.79 22.48 22.05 0.85 0.63 0.48 — A 10 08 12.23 +12 20 50.1V51 RRab 0.631211 23.09 22.62 22.46 21.98 0.38 0.43 0.27 — B 10 08 12.36 +12 18 53.0V52 RRab 0.5576776 22.99 22.66 22.57 22.14 1.30 1.03 0.40 — A 10 08 12.45 +12 17 09.2V53 RRab 0.5541654 22.87 22.54 22.09 22.01 1.23 0.94 0.54 — A 10 08 12.53 +12 17 54.7V56 RRc 0.3591109 22.82 22.56 22.44 22.08 0.58 0.48 0.44 — A 10 08 13.29 +12 16 41.7V57=HW14 RRab 0.5644922 22.91 22.42 22.06 21.74 0.78 0.54 0.44 — A 10 08 14.05 +12 15 06.2V58* RRc 0.368497 22.86 22.58 22.59 22.15 0.71 0.58 0.24 — A 10 08 14.18 +12 15 30.2V59 RRab 0.635755 22.87 22.41 22.04 21.71 0.62 0.40 0.32 0.25 A 10 08 14.20 +12 17 38.2V60 RRc 0.3755769 22.95 22.69 22.54 22.24 0.55 0.43 0.38 — A 10 08 14.30 +12 15 17.5V63 RRab 0.517857 22.90 22.73 22.77 22.09 1.52 1.20 — — B 10 08 15.37 +12 17 32.4V64 RRab 0.5419958 23.00 22.77 22.41 22.19 1.14 1.06 0.67 — A 10 08 15.38 +12 15 09.6V65 RRab 0.5252377 22.89 22.55 22.49 22.36 1.38 1.09 0.40 — A 10 08 15.52 +12 16 27.9V66 RRab 0.5413377 23.05 22.79 22.30 21.94 1.32 1.35 1.28 — A 10 08 15.70 +12 16 51.8V67 RRab 0.5505050 23.08 22.76 22.44 22.11 1.19 0.95 0.57 — A 10 08 16.12 +12 17 52.5V68 RRab 0.5348781 22.82 22.63 22.67 22.23 1.76 1.22 0.90 — A 10 08 16.32 +12 18 24.3V70 RRab 0.5559762 23.02 22.78 22.41 22.43 1.16 0.73 0.53 — A 10 08 16.44 +12 18 12.1V72 RRab 0.601951 23.10 22.78 22.47 22.27 0.57 0.54 0.29 — A 10 08 16.79 +12 22 03.7V73 RRab 0.547411 22.95 22.63 22.75 — 1.46 1.18 0.80 — A 10 08 17.01 +12 13 37.0V74 RRab 0.5097438 22.98 22.71 22.35 22.33 2.17 1.27 0.54 — A 10 08 17.44 +12 22 04.6V76 RRab 0.548736 23.01 22.71 22.45 22.37 1.16 1.11 0.42 — A 10 08 17.59 +12 15 25.1V77 RRab 0.5883625 22.99 22.67 22.62 21.86 1.15 1.08 0.45 — A 10 08 17.93 +12 21 51.1V79 RRc 0.417855 22.90 22.69 22.52 22.15 0.83 0.47 — — C 10 08 18.55 +12 11 07.5V80 RRc 0.3866383 22.66 22.44 22.33 22.03 0.50 0.47 0.31 0.19 A 10 08 18.99 +12 16 23.3V81 RRab 0.596522 23.15 22.82 22.50 22.23 0.85 0.62 0.59 — A 10 08 19.12 +12 15 31.5V82 RRab 0.5711366 23.12 22.77 22.72 22.14 1.34 1.02 — — A 10 08 19.43 +12 21 59.7V83 RRab 0.6778582 22.36 21.82 21.62 21.17 0.53 0.38 — — C 10 08 19.84 +12 18 27.4V84 RRab 0.615537 22.90 22.52 22.49 21.84 1.71 1.08 0.31 — A 10 08 19.88 +12 14 47.4V85 RRc 0.3175981 22.67 22.45 22.13 22.11 0.61 0.50 0.25 — A 10 08 19.90 +12 18 58.7V86 RRc 0.2851411 23.01 22.55 22.41 — 0.92 0.84 0.29 — B 10 08 19.91 +12 26 28.2V87 RRab 0.651948 23.02 22.64 22.35 21.91 0.60 0.46 0.34 — A 10 08 19.93 +12 15 55.3
32 –Table 2—Continued
ID Var period h B i h V i h R i h I i A B A V A R A I Q RA Dectype d mag mag mag mag mag mag mag mag hh mm ss dd mm ssV88 RRc 0.3173750 22.96 22.72 22.57 22.45 0.73 0.75 — — B 10 08 19.99 +12 23 03.1V89 RRab 0.5300555 22.85 22.64 22.49 22.27 1.40 0.92 0.83 — A 10 08 20.11 +12 16 12.9V91 RRab 0.602603 23.14 22.72 22.58 — 0.69 0.72 0.64 — A 10 08 20.58 +12 24 18.3V92* RRab 0.6304080 22.44 22.23 21.96 21.75 0.99 0.80 0.67 0.12 B 10 08 20.77 +12 19 26.9V93 RRab 0.5911859 22.89 22.47 22.26 21.86 1.00 0.70 0.50 0.25 A 10 08 20.94 +12 16 32.6V96 RRab 0.6429762 22.68 22.43 22.22 22.04 1.08 0.76 0.70 — A 10 08 21.47 +12 18 54.8V97 RRab 0.591565 23.14 22.77 22.62 22.21 0.58 0.35 0.25 — C 10 08 21.52 +12 23 45.1V98* RRab 0.626239 22.96 22.63 22.47 22.24 1.27 0.89 0.70 — A 10 08 21.53 +12 14 27.1V99 RRab 0.574072 23.28 22.92 22.66 22.45 1.19 0.73 — — B 10 08 21.68 +12 15 46.3V101 RRab 0.7071830 22.92 22.62 22.29 22.00 0.86 0.78 0.78 — B 10 08 21.86 +12 21 54.3V102 RRab 0.952880 22.88 22.39 22.57 22.34 0.73 1.18 — — C 10 08 22.21 +12 15 54.6V105 RRab 0.5541252 23.10 22.74 22.67 22.01 1.41 1.03 0.30 — A 10 08 22.92 +12 14 36.6V106 RRab 0.6021605 22.90 22.51 22.21 22.02 1.31 0.97 0.57 — A 10 08 23.12 +12 16 16.8V107 RRab 0.5573474 22.94 22.63 22.47 22.37 1.11 1.04 0.75 — B 10 08 23.20 +12 16 35.9V108 RRab 0.5467398 22.59 22.15 21.84 21.49 0.96 0.52 0.43 0.40 C 10 08 23.36 +12 19 07.1V110 RRab 0.6124265 22.85 22.57 22.50 22.04 0.84 0.49 0.21 — C 10 08 23.41 +12 20 30.7V111 RRab 0.6183992 22.82 22.56 22.22 22.11 1.26 0.93 0.50 — A 10 08 23.80 +12 17 26.0V112 RRab 0.5880014 22.31 21.81 21.47 21.12 0.69 0.44 — — C 10 08 24.26 +12 17 27.5V116 RRab 0.5593915 22.92 22.60 22.26 22.18 1.34 1.19 0.66 — B 10 08 24.72 +12 17 17.1V117 RRab 0.4948878 22.43 22.19 22.13 21.60 0.52 0.50 — — B 10 08 24.76 +12 19 13.5V118 RRab 0.5966387 22.89 22.61 22.35 22.19 1.06 0.75 0.70 — A 10 08 25.36 +12 17 31.4V119 RRc 0.3769470 22.95 22.66 22.52 — 0.51 0.47 0.34 — A 10 08 25.55 +12 23 28.1V121* RRab 0.5797797 22.94 22.69 22.33 21.92 1.24 0.85 0.47 — A 10 08 25.61 +12 15 35.5V124 RRab 0.5866781 23.08 22.72 19.95 22.16 1.21 0.82 0.73 — B 10 08 26.01 +12 18 49.9V125 RRab 0.5942164 22.96 22.76 22.58 22.28 1.04 0.78 0.58 — B 10 08 26.22 +12 21 53.4V127 RRab 0.5511360 22.18 21.64 21.22 20.88 0.55 0.29 0.20 0.10 B 10 08 26.42 +12 21 42.4V128 RRab 0.5398457 23.00 22.76 22.94 — 1.35 1.00 0.31 — A 10 08 26.61 +12 24 11.1V132 RRab 0.5881914 23.09 22.76 22.46 22.18 1.26 1.00 0.82 — A 10 08 27.00 +12 15 11.7V134 RRab 0.8547496 22.94 22.57 22.48 21.98 0.80 0.64 — — A 10 08 27.26 +12 15 44.5V135 RRab 0.5886723 23.03 22.53 22.25 21.99 1.35 0.77 0.51 — A 10 08 27.35 +12 20 42.7V137 RRab 0.5978996 23.19 22.81 22.74 22.02 1.08 0.83 — — A 10 08 27.61 +12 20 12.7V138 RRab 0.5881765 23.13 22.79 22.75 22.32 0.84 0.69 0.42 — A 10 08 27.87 +12 16 14.6V141 RRc 0.3667905 22.80 22.49 22.28 22.10 0.54 0.43 0.46 — A 10 08 28.13 +12 15 29.3V142* RRab 0.5871330 22.93 22.51 22.21 21.87 1.12 0.82 0.43 — A 10 08 28.23 +12 17 13.3V148 RRab 0.5599570 22.35 21.86 21.51 21.07 0.58 0.37 0.31 0.28 B 10 08 28.84 +12 20 12.2V149 RRc 0.3703344 22.99 22.67 22.45 22.11 0.65 0.47 0.36 — A 10 08 28.94 +12 15 01.2V151 RRab 0.6250526 22.80 22.46 22.15 21.91 0.69 0.59 0.56 0.44 C 10 08 29.40 +12 18 21.3V152 RRab 0.5552238 23.07 22.76 22.56 22.21 1.20 0.77 1.64 — A 10 08 29.40 +12 16 33.3V153 RRab 0.5612972 23.08 22.70 22.82 22.20 1.12 1.07 0.19 — A 10 08 29.41 +12 23 48.3V156 RRab 0.646932 23.35 22.92 22.58 22.23 0.64 0.57 0.55 0.36 C 10 08 30.41 +12 17 16.3V157 RRab 0.4903282 22.94 22.80 22.44 22.47 1.45 0.93 — — B 10 08 30.55 +12 18 37.8V158 RRc 0.3446077 22.87 22.58 22.38 22.12 0.61 0.45 — — B 10 08 30.60 +12 15 11.0V165 RRc 0.3231045 23.01 22.78 22.76 22.34 0.80 0.49 0.69 0.21 A 10 08 31.56 +12 21 48.1V166 RRc 0.3872139 22.83 22.58 22.41 21.94 0.58 0.42 0.36 0.82 B 10 08 31.55 +12 15 59.4V168* RRab 0.579400 23.00 22.64 22.57 21.96 1.12 1.02 0.52 — A 10 08 31.78 +12 23 35.3V169 RRab 0.5494407 23.25 22.85 22.68 22.19 1.16 0.91 0.35 0.30 A 10 08 31.94 +12 17 01.9V170 RRc 0.3890384 22.61 22.39 22.15 21.98 0.68 0.51 0.53 — B 10 08 31.94 +12 16 56.2V171* RRab 0.5608915 22.98 22.76 22.12 21.81 1.13 1.09 0.70 — A 10 08 32.06 +12 13 49.0V172 RRab 0.615115 23.33 22.87 22.55 22.26 0.85 0.52 0.42 — C 10 08 32.13 +12 16 35.3V174 RRab 0.5950424 23.02 22.67 22.42 22.14 0.73 0.65 0.32 — B 10 08 32.24 +12 16 33.5V175 RRab 0.6380987 22.88 22.60 22.38 22.02 0.43 0.39 0.46 — C 10 08 32.51 +12 16 04.9V176 RRab 0.573147 22.38 21.85 21.51 21.11 0.74 0.54 — — C 10 08 33.07 +12 17 15.5V180* RRab 0.619147 22.96 22.72 22.43 22.19 1.24 0.85 0.49 — A 10 08 34.12 +12 23 43.2V185 RRab 0.738784 22.88 22.49 22.16 21.80 1.15 0.86 0.69 — A 10 08 35.94 +12 22 49.2V186 RRab 0.6128591 22.89 22.58 22.27 22.03 1.26 0.89 0.56 — A 10 08 36.05 +12 21 27.8V187 RRab 0.631086 22.80 22.43 22.18 21.81 0.70 0.58 0.36 — A 10 08 36.09 +12 19 33.3V188 RRab 0.622437 23.20 22.74 22.65 22.07 0.42 0.50 0.21 — B 10 08 36.31 +12 27 15.9V189 RRab 0.6313365 23.01 22.57 22.17 22.01 1.22 0.93 0.73 — A 10 08 36.88 +12 19 45.0V192 RRab 0.667686 22.93 22.59 22.38 22.10 0.60 0.66 0.60 — A 10 08 37.44 +12 20 44.6V196 RRc 0.352561 22.86 22.65 22.55 22.16 0.80 0.61 0.37 — C 10 08 38.87 +12 21 52.8V198 RRc 0.364983 22.95 22.68 22.52 22.17 0.56 0.41 0.38 — A 10 08 39.35 +12 16 09.2V201* RRab 0.5663093 23.05 22.76 22.58 22.04 1.26 1.02 0.42 — A 10 08 40.00 +12 19 18.4V202* RRab 0.581068 23.05 22.69 22.58 21.95 1.23 0.67 0.21 — A 10 08 40.10 +12 14 28.5V204 RRab 0.748624 22.86 22.54 22.18 21.98 1.55 1.01 0.79 0.65 A 10 08 40.36 +12 18 53.7V205* RRab 0.6199640 22.92 22.60 22.29 22.16 1.08 0.86 0.48 — A 10 08 40.63 +12 19 18.7
33 –Table 2—Continued
ID Var period h B i h V i h R i h I i A B A V A R A I Q RA Dectype d mag mag mag mag mag mag mag mag hh mm ss dd mm ssV207 RRab 0.568369 22.98 22.66 22.56 22.29 1.25 0.86 0.43 — B 10 08 41.00 +12 19 57.9V213 RRab 0.602289 22.75 22.53 22.48 21.96 1.22 1.00 0.80 — A 10 08 42.91 +12 20 38.9V214* RRc 0.3190927 22.83 22.60 22.65 22.07 0.72 0.63 0.25 — A 10 08 44.37 +12 21 30.6V216 RRab 0.723219 22.98 22.64 22.26 21.99 0.79 0.71 0.30 0.69 B 10 08 44.45 +12 21 33.4V217 RRab 0.541640 23.04 22.72 22.49 22.25 0.88 1.01 0.52 0.66 C 10 08 45.15 +12 20 00.5V218 RRc 0.371302 22.98 22.68 22.49 22.09 0.58 0.40 0.35 — B 10 08 47.40 +12 21 08.0V219* RRab 0.587327 23.12 22.77 22.47 22.09 1.53 0.90 0.75 — A 10 08 48.15 +12 16 28.4V220 RRc 0.3124520 22.94 22.69 22.73 22.29 0.66 0.58 0.31 — B 10 08 48.32 +12 19 24.6V221* RRab 0.665057 22.76 22.50 22.10 21.72 1.20 1.02 0.64 — C 10 08 48.38 +12 26 51.8V222* RRab 0.566679 23.06 22.77 22.46 22.27 1.05 0.95 0.60 — A 10 08 48.43 +12 17 58.0V223 RRab 0.562598 22.91 22.72 22.55 22.35 0.86 0.74 0.38 — B 10 08 48.54 +12 17 52.5V224 RRab 0.647379 23.07 22.68 22.48 21.97 0.52 0.33 0.38 — B 10 08 48.68 +12 18 47.6V225 RRc 0.379586 23.01 22.70 22.40 22.15 0.54 0.43 0.40 — A 10 08 48.95 +12 21 41.2V226 RRc 0.2637345 22.76 22.50 22.55 22.15 0.58 0.50 0.24 — B 10 08 49.17 +12 21 26.3V227 RRab 0.644914 23.07 22.64 22.42 22.04 0.64 0.44 0.50 — C 10 08 49.36 +12 14 52.0V228* RRab 0.600145 23.10 22.74 22.46 22.08 1.00 0.83 0.34 — A 10 08 50.71 +12 18 00.0V229* RRab 0.5929263 22.92 22.71 22.74 22.31 1.49 1.02 0.41 — A 10 08 51.38 +12 21 54.3V230* RRab 0.565908 22.98 22.67 22.42 22.21 1.06 0.81 — — B 10 08 51.67 +12 15 56.6V231* RRab 0.623239 23.04 22.59 22.46 21.90 0.65 0.65 0.34 — B 10 08 51.73 +12 29 06.5V232* RRab 0.581134 22.98 22.66 22.33 21.90 1.12 0.90 0.43 — A 10 08 52.03 +12 21 00.5V234* RRab 0.719480 22.96 22.56 22.46 21.82 1.08 0.97 0.52 — A 10 08 52.99 +12 21 57.9V235* RRab 0.6603587 22.84 22.48 22.16 21.80 1.98 1.07 0.76 — A 10 08 53.20 +12 15 39.3V236* RRab 0.5605958 22.99 22.62 22.77 22.02 1.03 1.04 0.35 — A 10 08 54.82 +12 22 16.3V237* RRab 0.592160 22.86 22.51 22.38 21.94 1.22 1.23 0.42 — A 10 08 55.31 +12 17 41.1V239* RRab 0.6390675 22.85 22.60 22.23 22.28 1.01 0.77 0.83 — A 10 08 55.60 +12 21 25.1V240* RRab 0.601778 23.03 22.74 22.49 22.37 0.99 1.02 — — C 10 08 55.82 +12 14 50.7V243* RRc 0.358956 23.02 22.64 22.49 22.07 0.47 0.46 0.45 — B 10 08 57.04 +12 13 02.4V244* RRc 0.3881747 22.97 22.67 22.53 22.20 0.59 0.50 0.43 — B 10 08 58.23 +12 24 54.3V245* RRab 0.541998 23.16 22.65 22.44 22.18 0.36 0.79 0.54 — A 10 08 59.99 +12 19 49.4V246* RRab 0.5475991 23.07 22.66 22.27 22.20 1.06 0.90 — — B 10 09 01.43 +12 12 07.4V247* RRab 0.68277 22.86 22.46 22.51 21.91 1.79 1.14 — — B 10 09 04.52 +12 20 41.4V249* RRab 0.59676 22.96 22.67 22.62 22.07 0.75 0.69 — — B 10 09 06.66 +12 21 17.7V250* RRab 0.58273 22.94 22.70 22.77 22.28 0.64 0.59 — — C 10 09 08.17 +12 24 32.1V251* RRab 0.574 23.22 22.72 22.20 22.22 1.42 1.07 — — C 10 09 09.47 +12 19 26.3Note. — * indicates suspected Blazhko or double-mode pulsator.
34 –Table 3. AC/spC variable candidates: periods, photometric parameters, and positions.
ID Period h B i h V i h R i h I i A B A V A R A I Q Mode Mass RA Decd mag mag mag mag mag mag mag mag M ⊙ hh mm ss dd mm ssV20 1.332426 21.48 20.99 20.85 20.68 1.42 1.26 0.91 0.80 A FO 1.9 10 07 59.30 +12 17 45.5V27 1.6104244 20.88 20.51 20.35 19.87 1.84 1.46 0.46 0.60 A FO 1.5 10 08 05.42 +12 18 22.9V34 1.660596 20.78 20.54 20.18 19.93 1.74 1.30 1.00 0.83 A FO 0.8 10 08 08.16 +12 16 57.0V39 1.339882 21.12 20.75 20.32 20.02 0.87 0.68 0.37 0.20 A FU 2.6 10 08 08.89 +12 19 47.3V40 2.08521 20.91 20.52 20.36 19.96 0.84 0.62 0.60 0.53 A FO 1.1 10 08 09.06 +12 16 04.0V41 0.6001349 21.57 21.25 21.08 20.67 0.80 0.66 0.64 0.52 A FO 2.3 10 08 09.33 +12 15 10.5V45 0.8042814 21.27 20.92 20.77 20.49 0.77 0.68 0.70 0.46 A FU 4.1 10 08 11.07 +12 19 53.9V62 0.8356451 21.15 20.93 20.74 20.48 0.92 0.71 0.56 0.54 A FO 1.3 10 08 15.11 +12 20 41.4V69 1.656557 21.12 20.80 20.62 20.22 1.64 1.27 0.70 0.85 A FU 1.4 10 08 16.41 +12 16 51.1V71 0.7634400 21.72 21.38 21.25 20.83 0.84 0.66 0.56 0.47 A FO 1.5 10 08 16.75 +12 19 19.3V75 0.5443509 21.32 21.07 20.88 20.70 0.81 0.63 0.53 0.28 A FO 2.5 10 08 17.47 +12 18 04.4V78 1.304405 21.18 20.75 20.47 20.21 1.30 1.05 0.79 0.23 A FO 2.1 10 08 18.15 +12 18 08.0V94 0.7377763 20.90 20.62 20.46 20.05 0.91 0.68 0.26 0.26 A FO 3.1 10 08 21.12 +12 17 47.4V103 1.4382564 20.82 20.48 20.23 19.87 1.76 1.36 9.99 1.00 A FO 1.7 10 08 22.29 +12 18 01.9V113 0.7554243 21.18 20.89 20.67 20.39 0.89 0.60 0.37 0.38 A FO 2.2 10 08 24.27 +12 20 05.5V114=HW19 1.1232940 21.21 20.96 20.51 20.43 1.67 1.37 0.83 0.61 A FO 0.9 10 08 24.49 +12 15 04.8V115 1.0984368 21.40 21.01 20.76 20.44 1.47 1.20 1.12 0.87 A FO 1.6 10 08 24.53 +12 18 23.7V120 0.984594 21.13 21.00 20.82 20.46 1.46 1.21 0.67 0.53 A FO 0.6 10 08 25.58 +12 16 51.5V122 1.306483 21.28 20.86 20.67 20.31 1.24 1.00 0.55 0.46 A FO 1.7 10 08 25.92 +12 18 03.3V123 0.754381 21.42 21.03 20.74 20.44 0.44 0.31 0.31 0.31 A FU 4.8 10 08 25.99 +12 19 46.8V129=HW23 ∗
35 –Table 4. Positions and photometry for the candidate LPV stars.
ID “period” h B i h V i h R i h I i A B A V A R A I RA Decd mag mag mag mag mag mag mag mag hh mm ss dd mm ssV29 137.33 20.73 19.17 18.25 17.39 0.87 0.78 0.78 0.65 10 08 05.93 +12 15 22.7V30 22.966 21.06 19.59 18.77 18.07 0.16 0.12 0.12 0.11 10 08 06.43 +12 16 26.9V33 13.515 20.98 19.28 18.36 17.56 0.13 0.20 0.13 0.13 10 08 08.03 +12 15 18.1V50 52.523 20.83 19.31 18.51 17.79 0.38 0.26 0.33 0.21 10 08 12.29 +12 23 19.7V61 13.5399 20.91 19.11 18.10 17.34 0.26 0.20 0.17 0.11 10 08 14.59 +12 18 01.9V90 24.226 20.93 19.49 18.68 17.93 0.22 0.16 0.10 0.03 10 08 20.47 +12 21 04.0V95 29.027 20.95 19.11 18.11 17.33 0.53 0.23 0.18 0.11 10 08 21.38 +12 14 42.9V100 746 21.75 19.63 18.32 17.46 0.60 1.02 0.54 0.64 10 08 21.77 +12 17 25.1V126 451.9 21.55 19.45 18.40 17.46 0.47 0.31 0.43 0.17 10 08 26.36 +12 14 02.8V136 35.828 21.61 19.98 19.25 17.75 0.88 1.39 1.10 1.01 10 08 27.52 +12 16 54.0V139 86.00 20.90 19.53 18.87 18.05 0.29 0.16 0.40 0.13 10 08 27.95 +12 17 56.5V143 20.170 20.77 19.43 18.48 17.60 0.48 0.26 0.24 0.22 10 08 28.51 +12 19 48.8V162 65.075 20.72 18.80 17.84 17.03 0.25 0.21 0.21 0.03 10 08 31.04 +12 17 07.8V167 63.334 20.97 19.70 18.84 17.66 0.65 0.69 0.92 0.27 10 08 31.76 +12 18 18.1V178 23.708 21.27 19.98 19.18 18.58 0.16 0.14 0.02 0.06 10 08 34.06 +12 23 16.7V183 73.73 21.43 19.26 18.14 17.20 0.44 0.43 0.40 0.24 10 08 34.72 +12 18 37.7V190 16.6496 20.70 19.16 18.35 17.57 0.23 0.17 0.15 0.12 10 08 36.92 +12 20 11.5V200 30.594 21.93 19.80 18.22 17.35 0.38 0.76 0.63 0.38 10 08 39.90 +12 22 14.8V238 36.707 21.06 19.70 18.82 18.08 0.38 0.38 0.34 0.40 10 08 55.53 +12 14 01.7
Table 5. Suspected variables of unknown type ID h B i h V i h R i h I i RA Decmag mag mag mag hh mm ss dd mm ssV12 23.23 22.46 22.02 21.88 10 07 49.83 +12 34 06.3V46 22.84 22.77 22.48 22.31 10 08 11.32 +12 16 26.8V48 23.02 22.52 21.85 23.02 10 08 12.13 +12 17 19.0V54 23.10 22.72 22.68 22.28 10 08 12.54 +12 15 55.9V55 22.87 22.75 22.45 22.41 10 08 13.06 +12 10 48.4V104 23.10 22.84 23.34 22.81 10 08 22.34 +12 16 26.6V109 22.71 22.73 22.44 22.10 10 08 23.38 +12 12 50.6V161 24.06 23.48 23.00 22.93 10 08 30.88 +12 17 17.1V191 23.46 22.93 22.29 22.32 10 08 37.37 +12 13 01.2V206 22.07 21.91 21.79 21.37 10 08 40.90 +12 20 58.4V210 22.91 22.81 22.17 22.17 10 08 42.23 +12 16 10.8V233 22.74 21.68 21.54 21.15 10 08 52.16 +12 23 49.2V242 21.59 21.37 21.47 20.74 10 08 56.21 +12 20 02.7
Table 6. Mean RRL properties of selected dwarf Spheroidal galaxies.
Galaxy’s name h Pab i h Pc i N RRab N TOT N RRc /N TOT
M/M ⊙ < [Fe/H] > d d dexUFDs 0.640 ± σ = 0.06) 0.377 ± σ = 0.03) 43 61 0.29 ≤ . × ≤ –2.0Draco 0.615 ± σ = 0.04) 0.389 ± σ = 0.03) 214 270 0.21 0 . × –1.93Carina 0.635 ± σ = 0.05) 0.370 ± σ = 0.05) 63 82 0.23 0 . × –1.72Tucana 0.604 ± σ = 0.06) 0.361 ± σ = 0.05) 216 298 0.27 0 . × –1.95Sculptor 0.590 ± σ = 0.08) 0.337 ± σ = 0.04) 131 221 0.41 2 . × –1.68Cetus 0.614 ± σ = 0.04) 0.383 ± σ = 0.04) 503 630 0.20 2 . × –1.90Leo I 0.596 ± σ = 0.05) 0.357 ± σ = 0.03) 81 95 0.15 5 . × –1.43Note. — N TOT does not include d–type RRLs. Mass of the galaxy is given assuming a stellar mass-to-light ratio of 1 (McConnachie 2012).
36 –Table 7. Mean properties for RR Lyraes in the Galactic halo, globular clusters, classicaland UFD dwarf galaxies.
Galaxy’s name h Pab i N RRab N RROoI /N RRab dInner halo a 0.586 ± σ = 0.08) 1898 0.64Inner halo b 0.583 ± σ = 0.08) 4064 0.67Outer halo a 0.577 ± σ = 0.07) 5659 0.73Outer halo b 0.580 ± σ = 0.07) 2268 0.69GCs a ± σ = 0.12) 1054 0.67dSphs & UFDs 0.610 ± σ = 0.05) 1299 0.78Note. — a The GCs included here are NGC 1851, NGC 1904, NGC 3201,NGC 362, NGC 4147, NGC 4590, NGC 4833, NGC 5024, NGC 5139, NGC 5272,NGC 5286, NGC 5466, NGC 5897, NGC 5904, NGC 5986, NGC 6101, NGC 6121,NGC 6171, NGC 6205, NGC 6229, NGC 6266, NGC 6333, NGC 6341, NGC 6362,NGC 6626, NGC 6715 NGC 6809, NGC 6864, NGC 6934, NGC 6981, NGC 7006,NGC 7089, NGC 078, NGC 7099 and IC 4499. The collected data belong tothe updated 2013 catalog of (Clement et al. 2001) with few exceptions, namelyNGC 6121 (Stetson et al. 2014, submitted), NGC 6934 (Kaluzny et al. 2001) andNGC 7006 (Wehlau et al. 1999).
Table 8. Chi-squared likelihood that RRab period distributions come from the sameparent population (expressed in per cent).
Sample Inner halo a Inner halo b Outer halo a Outer halo b GCs dSphs & UFDs LMCInner halo a 100 25(36) a a −
100 0 0 0(2) a − −
100 0 0 0 0Outer halo b − − −
100 0 0 0GCs − − − −
100 0 0dSphs & UFDs − − − − −
100 0LMC − − − − − − a In parentheses we report the likelihood coming out of the KS-test.
37 –Table 9. Candidates from Hodge & Wright (1978) that do not appear variable in our data ID h B i h V i h R i h I i RA Decmag mag mag mag hh mm ss dd mm ssHW4 23.72 22.95 22.52 22.02 10 08 32.79 +12 18 37.6HW5 23.82 23.28 22.90 22.39 10 08 32.13 +12 18 03.7HW6 21.55 20.64 20.09 19.55 10 08 37.97 +12 17 45.5HW7 22.06 21.21 20.67 20.14 10 08 31.88 +12 16 40.4HW9 21.53 20.29 19.58 18.95 10 08 28.36 +12 16 04.5HW11 21.28 19.86 19.16 18.57 10 08 36.39 +12 16 02.1HW12 22.46 21.04 20.43 19.91 10 08 40.51 +12 15 04.5HW13 22.98 21.71 20.79 20.08 10 08 29.62 +12 15 07.8HW15 21.81 20.88 20.32 19.79 10 08 27.98 +12 16 53.5HW18 21.76 20.85 20.30 19.78 10 08 28.77 +12 15 06.5HW20 20.81 19.21 18.27 17.54 10 08 15.86 +12 20 32.6HW21 21.31 20.39 19.69 19.12 10 08 19.45 +12 16 34.1HW22 20.90 19.75 19.05 18.41 10 08 26.48 +12 18 23.9
Table 10. Cross-identifications for stars from Azzopardi et al. (1986)
ALW Our
B V R I
RA DecID ID a hh mm ss dd mm ss9 91206 21.73 19.82 19.28 17.73 10 08 27.52 +12 16 54.010 91627 21.21 19.45 18.42 17.62 10 08 27.69 +12 17 31.912 97647 21.13 19.60 18.78 18.18 10 08 30.09 +12 18 34.314 99852 20.70 18.83 17.83 17.03 10 08 31.04 +12 17 07.815 67567 = V79 21.34 19.28 18.15 17.16 10 08 34.72 +12 18 37.718 109932 21.18 19.28 18.33 17.59 10 08 35.30 +12 17 24.619 62888 21.12 19.31 18.41 17.72 10 08 16.61 +12 20 07.220 66463 21.64 19.49 18.40 17.50 10 08 18.09 +12 20 09.4 a Sequential identification number in our catalog for Leo I, available on-line.
38 –Fig. 1.— Examples of our light curves to illustrate the quality obtained in the B and V bands. We have selected one RRc (V60, left panel), one RRab (V92, possibly a Blazhko ordouble-mode pulsator candidate, middle) and one AC (V120, right). This last star has theuncommon 1 d period, but with our twenty-year baseline aliasing is minimal. 39 –Fig. 2.— CMD and topology of the instability strip for the variable stars of Leo I. Red circlesand blue triangles represent RR Lyrae stars pulsating in the fundamental and first-overtonemode, respectively. Squares and stars represent Cepheids and LPV variables. Variablecandidates with uncertain periods are indicated by yellow symbols. 40 –Fig. 3.— Bailey diagram for RRL stars in the B (top panel) and V bands (bottom panel).Symbols are as in Fig. 2. We have also shown, as solid curves, the empirical relations forOosterhoff types I (left) and II (right) as derived by Cacciari et al. (2005). 41 –Fig. 4.— Distribution of our variable candidates in the field of view of our photometricstudy. The symbols for variable stars are color-coded as in Fig. 2. Gray dots representnon-variable stars and are expanded according to their brightness in the V band. Units arearcseconds relative to an arbitrary origin (stated in text); X increases east (to the left) andY increases north (up). An ellipse represents an isopleth of the stellar distribution with asemimajor axis of 12 ′ as derived by Mateo et al. (2008). 42 –Fig. 5.— The period-luminosity distribution of RRLs in Leo I. The symbols are color-codedas in Fig. 2. The solid line corresponds to the FOBE position as constrained by V85, whereasthe dashed one is drawn to match V80, V141, and V214. 43 –Fig. 6.— The same as in Fig. 3 but for Cepheids. Both amplitude distributions, A V (bottompanel) and A B (top panel), clearly show that our Cepheid sample forms two clusters of starsdistinguished by small amplitudes for log P ∼ <
0, and larger amplitudes for log P ∼ >
0. Starsare color-coded according to their classification as derived in Section 5. Filled green andblack squares correspond to FU and FO pulsators respectively. Crosses indicate those starswith uncertain mode classification (see the text for details). 44 –Fig. 7.— The Wesenheit distribution for RRLs and Cepheids. The same color code asin Fig. 6 has been used. The theoretical Wesenheit relations for stars with masses of 1.8and 3.2 M ⊙ pulsating in the fundamental mode (solid and dotted lines, respectively) areshown, as predicted by Fiorentino et al. (2006). Arrows indicate the offset of the 1.8 M ⊙ first-overtone (FO) locus from the FU locus. For comparison, the location of Population IICepheids is also indicated by a dashed line (Di Criscienzo et al. 2007). 45 –Fig. 8.— The mass distribution for Leo I Anomalous and short-period Cepheids (black solidlines and gray fill) has been computed using the approach discussed in (Fiorentino & Monelli2012). The mean mass and the standard deviation reported in the legend exclude six outlierswith derived masses larger than 5 M ⊙ . For comparison, the dashed lines show the massdistribution of ACs in the OGLE sample for the LMC. The mean mass of the LMC Cepheidshas also been specified in the legend. 46 –Fig. 9.— left – Period versus V -band amplitude diagram for RRLs observed in eleven UFDs(top panel) and six dSphs (remaining six panels) for which mostly complete samples ofvariable stars have recently been published. The galaxies have been ordered by increasingbaryonic mass from top to bottom. The UFDs have been considered together to improve thestatistical significance of the plot. right – Period-frequency distributions of the same RRLs;the mean period of the RRab variables is indicated by a vertical red line in each panel. Thevertical blue dashed lines in this plot indicate log P = − .
35, which we have taken as thenominal boundary between RRc (FO) and RRab (FU) variables. 47 –Fig. 10.— Same as Fig. 9 for the AC/spC sample. 48 –Fig. 11.— left –
Period versus V -band amplitude diagram for RRab observed in the Galactichalo (panels a, b, c, d)), compared to GCs (panel e) and dSphs plus UFDs (panel f). Thetwo Oosterhoff curves given in Cacciari et al. (2005) are also indicated together with a meancurve used to notionally separate the OoI from the OoII components (blue dashed line). right –right –