ReadingTimeMachine/rtm-sgt-ocr-v1 · Datasets at Hugging Face

source stringlengths 1 2.05k ⌀	target stringlengths 1 11.7k
Indeed. the cwarl galaxies of the Local Group. which experienced a variety of SELL. always show stars on the red. plume.	Indeed, the dwarf galaxies of the Local Group, which experienced a variety of SFHs, always show stars on the red plume.
The CALDs of Local Group cwarfs exhibit either the red. plume only (dSph). a weak blue plume as well as a red. plume (clLir/dSph). or a strong blue plume combined with a red plume (dli).	The CMDs of Local Group dwarfs exhibit either the red plume only (dSph), a weak blue plume as well as a red plume (dIrr/dSph), or a strong blue plume combined with a red plume (dIrr).
οσο morphologics serve as a guide to our exploration of ονο CMDs.	These morphologies serve as a guide to our exploration of CHVC CMDs.
We display in Figure 4 the V-L I] CMD of the transition αςαρ galaxy. Phoenix.	We display in Figure 4 the [V-I, I] CMD of the transition dIrr/dSph galaxy Phoenix.
The observations were obtained with the VEEP UT2. "IXOYIEN and FORS2 under good secing conditions.	The observations were obtained with the VLT UT2 “KOYEN" and FORS2 under good seeing conditions.
Phe exposure times were 3 300s each. in V and Lb. Phe data were reduced. and single-star photometry applied in the same wav. as described above for the CLIVCS.	The exposure times were 3 $\times$ 300s each, in V and I. The data were reduced, and single-star photometry applied in the same way, as described above for the CHVCs.
We measure a EWILM in V of 07554. and in L of 0777.	We measure a FWHM in V of 54, and in I of 7.
Ehe total number of stars detected in. V. and Eis S780. much higher iui in anv of the. CIIVC. fields.	The total number of stars detected in V and I is 8780, much higher than in any of the CHVC fields.
Phe DAOPLIOT photometric errors are. clisplavecl in Figure 5. in order to emphasise that the quality of these cata is similar to that of the CLIVC data sets.	The DAOPHOT photometric errors are displayed in Figure 5, in order to emphasise that the quality of these data is similar to that of the CHVC data sets.
Phoenix. is located at. high Galactic latitude. near 69°.	Phoenix is located at high Galactic latitude, near $^o$.
Phe foreground contamination and extinction are small.	The foreground contamination and extinction are small.
H has a distance modulus of about 23.25 (Mateo. 1998) or a distance of about 450 Ixpc.	It has a distance modulus of about 23.25 (Mateo 1998) or a distance of about 450 Kpc.
The CALD of Phoenix shows why it is classed as having a mixed morphology: it exhibits both a blue. plume of relatively voung stars. and a red plume which includes the ROB.	The CMD of Phoenix shows why it is classed as having a mixed morphology: it exhibits both a blue plume of relatively young stars, and a red plume which includes the RGB.
The stellar distributions on the CALDs of the CLVCs are quite dilferent from that of known Local Group carts.	The stellar distributions on the CMDs of the CHVCs are quite different from that of known Local Group dwarfs.
The morphology of the CLIVC CALDs does not reveal cither of the two plumes.	The morphology of the CHVC CMDs does not reveal either of the two plumes.
Instead. the CALDs of the CLIVCs are reminiscent of other. deep HIST. pointings into the clisk and halo of our Galaxy (e... Richer et al.	Instead, the CMDs of the CHVCs are reminiscent of other, deep HST pointings into the disk and halo of our Galaxy (e.g., Richer et al.
2002).	2002).
We overplotted on Fig.	We overplotted on Fig.
1 theoretical isochrones from the database of Girardi et al. (	1 theoretical isochrones from the database of Girardi et al. (
2000) for ages of LO and 100 Myr. and for 1 and 10 Civr.	2000) for ages of 10 and 100 Myr, and for 1 and 10 Gyr.
We adopted a metallicity of 1/5 of Solar. as cwarft galaxies tend to have sub-Solar metallicities.	We adopted a metallicity of 1/5 of Solar, as dwarf galaxies tend to have sub-Solar metallicities.
1n general. the elfect of varving metallicity at a given age is that more metal-poor isochrones shift to bluer V-I colors. while more metal-rich ones show redder colors.	In general, the effect of varying metallicity at a given age is that more metal-poor isochrones shift to bluer V-I colors, while more metal-rich ones show redder colors.
This shift in minimal for the voung" isochrones (e.g.. 10 LOO Myr). but noticable for the "old" isochrones (1 10 Cyr).	This shift in minimal for the “young" isochrones (e.g., 10 100 Myr), but noticable for the “old" isochrones (1 10 Gyr).
We show the set with a metallicity of 1/5 of Solar because the 1 10 Gyr isochrones have sullicicntly rec colors to encompass the data clistribution.	We show the set with a metallicity of 1/5 of Solar because the 1 10 Gyr isochrones have sufficiently red colors to encompass the data distribution.
When comparing the data with the “voung™ isochrones. we notice there are virtually no stars observed with V-l< 1. where we expect to find a voung population.	When comparing the data with the “young" isochrones, we notice there are virtually no stars observed with $<$ 1, where we expect to find a young population.
This emphasises the absence of young MS and blue supergiant stars in the CLIVCs.	This emphasises the absence of young MS and blue supergiant stars in the CHVCs.
The brightest’ evolutionary phase of stars with ages larger than about one Civr occurs on the red eiant. branch.	The brightest evolutionary phase of stars with ages larger than about one Gyr occurs on the red giant branch.
Furthermore. the PROB. at M; of 4. is a well calibrated distance indicator (Lee. Freedman Madore 1983). and at a distance of 1 Alpe. it would appear at Iz221.	Furthermore, the TRGB, at $_I$ of –4, is a well calibrated distance indicator (Lee, Freedman Madore 1983), and at a distance of 1 Mpc, it would appear at $\approx$ 21.
We overplotted on Fig.	We overplotted on Fig.
1 theoretical isochrones with ages of 1 ancl LO Cie.	1 theoretical isochrones with ages of 1 and 10 Gyr.
Note that the Cirardi et al.	Note that the Girardi et al.
isochrones extend. above the RGB into the thermally pulsing ACB phase.	isochrones extend above the RGB into the thermally pulsing AGB phase.
A PROB would have been detected for the entire metallicity range appropriate for cwarl galaxies. to L magnitudes of at least 22.5. or to a distance of up to 2 Alpe. in the four CIIVCS whose CMDs are shown in Fig.	A TRGB would have been detected for the entire metallicity range appropriate for dwarf galaxies, to I magnitudes of at least 22.5, or to a distance of up to 2 Mpc, in the four CHVCs whose CMDs are shown in Fig.
1.	1.
Notice that the absence of stars with V-E« 1 rules out extremely metal-poor ICDs. which can be observed in dwarf Spheroidal galaxies.	Notice that the absence of stars with $<$ 1 rules out extremely metal-poor RGBs, which can be observed in dwarf Spheroidal galaxies.
Since no ROB is obvious. a straightforward conclusion is that there is none. and that these objects are not faint galaxies within the Local Group.	Since no RGB is obvious, a straightforward conclusion is that there is none, and that these objects are not faint galaxies within the Local Group.
The absence of a PROB with Lz21 suggests the hypothesis hat LIIPASS J1712-64 is a faint cwarl ealaxy of the Local Group is similarly unliklev.	The absence of a TRGB with $\leq$ 21 suggests the hypothesis that HIPASS J1712-64 is a faint dwarf galaxy of the Local Group is similarly unlikley.
To illustrate the expected ocations of RGDs in our data if ΤΗΑΡ 1712-64 were at à distance of 3.2 Alpe as proposed by Ixilborn et al. (	To illustrate the expected locations of RGBs in our data if HIPASS J1712-64 were at a distance of 3.2 Mpc as proposed by Kilborn et al. (
2000). we overplotted in Fig.	2000), we overplotted in Fig.
2 the shifted Globular Cluster ridgelines rom da Costa ArcmandrolT (1990) of M 15. Fe/11]—-2.17. GO 6397. 21.91. NL 2. -1.58. NGC 6752. -1.54. NGC 1551. -1.29. and 47 Tuc. -0.71.	2 the shifted Globular Cluster ridgelines from da Costa Armandroff (1990) of M 15, [Fe/H]=-2.17, NGC 6397, -1.91, M 2, -1.58, NGC 6752, -1.54, NGC 1851, -1.29, and 47 Tuc, -0.71.
Even i£ HIDPASS J1712-64 were as ar as 3.2 Mpe. with a corresponding PROB I-magnitude of 23.5. we could. have detected. it if its metallicity were ow enough. although this is clearly at the very limit. of our data.	Even if HIPASS J1712-64 were as far as 3.2 Mpc, with a corresponding TRGB I-magnitude of 23.5, we could have detected it if its metallicity were low enough, although this is clearly at the very limit of our data.
Lewis et al. (	Lewis et al. (
2002) recently presented. single-star whotometry obtained under sub-aresee secing conditions in a 0755 [field of view centered at 17 12 35.7. 64 38 12 J2000) and reaching magnitudes fainter than r—24 and e=25.	2002) recently presented single-star photometry obtained under sub-arcsec seeing conditions in a $\times$ 5 field of view centered at 17 12 35.7, –64 38 12 (J2000) and reaching magnitudes fainter than r=24 and g=25.
They. too. Failed to detect an obvious stellar content associated with LIEPASS J1712-64.	They, too, failed to detect an obvious stellar content associated with HIPASS J1712-64.
Furthermore. they rule out any uncderlving unresolved dilfuse components within their large field of view to deep surface-brightness limits.	Furthermore, they rule out any underlying unresolved diffuse components within their large field of view to deep surface-brightness limits.
We can set upper limits on the non-detection of a possible οΝΟ stellar. population in the VET/FORS data. by constructing CALDs and luminosity functions (LE) with ποσο sub-populations.	We can set upper limits on the non-detection of a possible CHVC stellar population in the VLT/FORS data, by constructing CMDs and luminosity functions (LF) with added sub-populations.
This is accomplished by using the VLT FORS2 observation of Phoenix obtained. under very similar conditions and with similar photometric limits. as a template stellar population with which we 7contaminate" the observed. CLIVCG CMDs.	This is accomplished by using the VLT FORS2 observation of Phoenix obtained under very similar conditions and with similar photometric limits, as a template stellar population with which we “contaminate" the observed CHVC CMDs.
We added: Phoenix stars to the IVC οος at their location in the Phoenix €MD.	We added Phoenix stars to the HVC CMDs at their location in the Phoenix CMD.
ln other words. we added	In other words, we added
during the day. even on days when the general RFI level should be quite low (e.g.. the Easter holiday).	during the day, even on days when the general RFI level should be quite low (e.g., the Easter holiday).
Dipoles deployed at ground level should suppress RFI sources located near the horizon and should enable more sensitive RRL (Ellingson2009). (Kogan&	Dipoles deployed at ground level should suppress RFI sources located near the horizon and should enable more sensitive RRL \citep{ellingson09}. \citep{kc09},
Cohen.2009)... 4— 5c 107 The Η15>265(115MHz «v190MHz). (Payneetal..1959).	\ref{tab:lwadetect} $5\sigma$ $10^{-3}$ The $11.5> z > 6.5$ $115\,\mathrm{MHz} < \nu < 190\,\mathrm{MHz}$ \citep{1989ApJ...341..890P}.
1077 21" ?2))	$10^{-3}$ $21^{\mathrm{st}}$ \ref{sec:cosmo})
1077 21" ?2)).	$10^{-3}$ $21^{\mathrm{st}}$ \ref{sec:cosmo})
Since the advent. of high qualiy CCDs there has been a great deal of effort to obtain (inillimagnuitude. or less than 15€)) time-series. optical aud infrared photometry lor laree uumbers of stars.	Since the advent of high quality CCDs there has been a great deal of effort to obtain high-precision (millimagnitude, or less than ) time-series, optical and infrared photometry for large numbers of stars.
Photometry at te inillimagnitude leve of precision is considered a prerequisite in searches [or trausiting Jupiter-sized planets arouud sdlar-type stars.	Photometry at the millimagnitude level of precision is considered a prerequisite in searches for transiting Jupiter-sized planets around solar-type stars.
As a result. the number oL groups that have achieved. this level of precisiou is oo mauy to list here (see for example Horne2003 for a list of trausitit& planet searches).	As a result, the number of groups that have achieved this level of precision is too many to list here (see for example \citealt{horne03} for a list of transiting planet searches).
Tie search [or transiting planets has uot been in vain: at the time of writing there are 7 known ransitiug planets. 6 of which were first identiliel photometrically (see e.g. Udalskietal.2001 αιd Ixouackietal. 2004)).	The search for transiting planets has not been in vain; at the time of writing there are 7 known transiting planets, 6 of which were first identified photometrically (see e.g. \citealt{udalski04} and \citealt{konacki04}) ).
Milliniagnuitude photometry bas also contributed to the study of stars near the hydrogen fusion limit	Millimagnitude photometry has also contributed to the study of stars near the hydrogen fusion limit
especially near (he centre of the svstem.	especially near the centre of the system.
Of course the svstem may not be fully relaxed when the black hole is growing.	Of course the system may not be fully relaxed when the black hole is growing.
Results are shown in Figure 1. [or black hole masses from zero through 0.001. 0.01. 0.1 and 1.0 as labelled by increasing, effect.	Results are shown in Figure \ref{fig:iso} for black hole masses from zero through $0.001$ , $0.01$, $0.1$ and $1.0$ as labelled by increasing effect.
These dimensionless masses correspond to plivsical masses in the range of about 10* to LOM AL. if we use (he values of py and ry given in section yoQu.	These dimensionless masses correspond to physical masses in the range of about $10^7$ to $10^{10}$ $M_{\odot}$ if we use the values of $\rho_0$ and $r_0$ given in section \ref{sec:growth}.
We find the same characteristic density cusp of 2/7?> as obtained by Young(1980).. and the radial velocity profiles are also similar.	We find the same characteristic density cusp of $R^{-3/2}$ as obtained by \citet{you80}, and the radial velocity profiles are also similar.
Although Young(L980) did not show results for the anisotropy parameter. defined as 3=1—<Vg>/«2Vg9>. our results do agree with the prediction of Quinlanefaf(1995). lor distributions of this (wpe.	Although \citet{you80} did not show results for the anisotropy parameter, defined as $\beta = 1 - <V_T^2> / <2 V_R^2>$, our results do agree with the prediction of \citet{qui95} for distributions of this type.
We observe that (here is at most about a 106 anisotropy in favour of tangential motion. due to the increasing binding energv of a particle aud the resultant decrease in eccentricity al constant. angular momentum.	We observe that there is at most about a $10$ anisotropy in favour of tangential motion, due to the increasing binding energy of a particle and the resultant decrease in eccentricity at constant angular momentum.
Moreover (he perturbation extends as far as the core radius only when (he mass is comparable to that of the core. as is to be expected.	Moreover the perturbation extends as far as the core radius only when the mass is comparable to that of the core, as is to be expected.
These results are (he main (test of our progralni.	These results are the main test of our program.
We explore here the aclabatic growth of a central black hole in a collisionless dark matter halo that possesses a central cusp.	We explore here the adiabatic growth of a central black hole in a collisionless dark matter halo that possesses a central cusp.
For initial svstems. we choose (wo very dillerent starting points.	For initial systems, we choose two very different starting points.
The first is an isotropic sell-similar svstem. meant (o represent the final state of a halo undergoing self-similar relaxation (Ilenriksen&Widrow1995).	The first is an isotropic self-similar system, meant to represent the final state of a halo undergoing self-similar relaxation \citep{hen95}.
. The other svstem is the NEW system of Navarro.Frenk&White(1996).. which fits a wide range of dark matter halo sizes.	The other system is the NFW system of \citet{nav96}, which fits a wide range of dark matter halo sizes.
The self-similar svstem is described by the DF where 6 is a[ree parameter (2/3<901.0> 1) that essentially controls (he logarithmic slope of the initial density and potential. given by	The self-similar system is described by the DF where $\delta$ is afree parameter $2/3 < \delta < 1,\ \delta > 1$ ) that essentially controls the logarithmic slope of the initial density and potential, given by
keV band is 9.241012P cre 1 " ↸⊳⋯−↸⊳∪↥⋅↥⋅↸∖↴∖↴↻∪↕∐∐∐∶↴∙↑∪⋅⋝ a huninosty of L1 « 10/9 ere \|.	keV band is $\times10^{-12}$ erg $^{-1}$ $^{-2}$ corresponding to a luminosity of 1.1 $\times$ $^{40}$ erg $^{-1}$.
The contribulon of t1 power-law component. the thermal black body coll)OLnont. aud the thermal thin plasma comiponenu to the total O.3-8.0 keV flux are 77C. 16% aud respectiveY.	The contribution of the power-law component, the thermal black body component, and the thermal thin plasma component to the total 0.3-8.0 keV flux are 77, 16 and, respectively.
However. a longer EPIC-PN exposure Is nieed to discutangle a possible extended ticrinal thin plasnua ficun the »)ut-like power-law and thermal tick plasna couponenuts.	However, a longer EPIC-PN exposure is needed to disentangle a possible extended thermal thin plasma from the point-like power-law and thermal thick plasma components.
Tle spectral analysis of the combined NAAI data has coufirmed two spectral compoucuts in Toll N-1.	The spectral analysis of the combined XMM data has confirmed two spectral components in HoII X-1.
The hare COLL2011011 is best described by a powerlaw (D~ 2.6). aud the soft component is fitted by thermal nodels with relatively ow temperatures (AL~0.110.22 keV).	The hard component is best described by a powerlaw $\Gamma\sim2.6$ ), and the soft component is fitted by thermal models with relatively low temperatures $kT\sim0.14-0.22$ keV).
The intrinsic N-rav ΠΟ ΕΕΠΠ isotroplic emission) is about H9 oye tin the 0.3-8.0 keV band	The intrinsic X-ray luminosity (assuming isotrophic emission) is about $^{40}$ erg $^{-1}$ in the 0.3-8.0 keV band.
Before we discuss the nature of the ultraluminuots Νταν source ΠΟ N-1 we stuuarize our resits: There are three niodels for ULX extensively discussed in the literature: 1) ULX 1iuav be black holes of “normal” stellar masses (~ LOAL. ) in binaries. which accrete gas in a supercritical ποσο,	Before we discuss the nature of the ultraluminous X-ray source HoII X-1 we summarize our results: There are three models for ULX extensively discussed in the literature: i) ULX may be black holes of ”normal” stellar masses $\sim 10$ $_{\sun}$ ) in binaries, which accrete gas in a supercritical regime.
They are 33-like objects or mücroquasars m their transicut activity staCs. Wwrose hard radiation can be collimated (and heaued) along Jes ancl accretion cisks axes (Fabrika Moeschervakov. 012001.. Iàiug et al. 20013).	They are 433-like objects or microquasars in their transient activity states, whose hard radiation can be collimated (and beamed) along jets and accretion disks axes (Fabrika Mescheryakov, \cite{Fa01}, King et al., \cite{Ki01}) ).
i) ULX may be black holes wihi ew tens of soar Wass (~ LOAD: ) and that heir ταν enidesbons ds from the disk shining at super-Eddiieton Iuuinosities (Beecluan 20023).	ii) ULX may be black holes with few tens of solar mass $\sim 10$ $_{\sun}$ ) and that their X-ray emissions is from the disk shining at super-Eddington luminosities (Begelman \cite{Beg02}) ).
The slim disk luocl (A»yenmowiez ct al. 1988..	The slim disk model (Abramowicz et al. \cite{Abra88},
Ebisawa et al. 20033)	Ebisawa et al. \cite{Ebi03}) )
alows to explain the observed super-Eddineton hunuinosity. hard A-rav spectra and spectral variations.	allows to explain the observed super-Eddington luminosity, hard X-ray spectra and spectral variations.
ii) The ULN may be iuteriiediate mass black holes (IMBUs). ~ 107M; (€'olbert Mushotzky 19993). which were formed from the very first stars (Aladau Rees 20013) or in globular clusters (Miller ILuuiltou 2002)).	iii) The ULX may be intermediate mass black holes (IMBHs), $\sim 10^3$ $_{\sun}$ (Colbert Mushotzky \cite{Col99}) ), which were formed from the very first stars (Madau Rees \cite{Mad01}) ) or in globular clusters (Miller Hamilton \cite{Mil02}) ).
The IMDBII could. acciTo eas frou a close companion or even from the interstellar medium aud become bright XPav SOULCCS when they are in dense Bas environments.	The IMBH could accrete gas from a close companion or even from the interstellar medium and become bright X--ray sources when they are in dense gas environments.
Spectra of some ULX (Miller et al. 20033)	Spectra of some ULX (Miller et al. \cite{Mil03}) )
support the idea that they ave IMDITs.	support the idea that they are IMBHs.
We call sources “trausicuts” when the iutrinsic duration of UIIECR production at a source 07 is shorter than the characteristic time profile spread. τς) at the observed euergv of £.	We call sources "transients" when the intrinsic duration of UHECR production at a source $\delta T$ is shorter than the characteristic time profile spread $\tau(E)$ at the observed energy of $E$.
Itf the time dispersion is longer thaw the time scale of UITECT. observatious. we nüsperceive that ΤΕΠΟ bursts ave steady sources. aud therefore can define the “apparent” number density of UITECR sources DUE).	If the time dispersion is longer than the time scale of UHECR observations, we misperceive that UHECR bursts are steady sources, and therefore can define the "apparent" number density of UHECR sources $n_s(E)$.
The source muuber deusity is related to the rate of UIIECR bursts p, as The source nuuber deusitv generally depends on UIIECR cnereies. since the apparent duration is dependent on cucreics explicitly.	The source number density is related to the rate of UHECR bursts $\rho_s$ as The source number density generally depends on UHECR energies, since the apparent duration is dependent on energies explicitly.
\| UIIECTs observed at the Earth suffer from the GME and EGAIFs embedding their sources.	UHECRs observed at the Earth suffer from the GMF and EGMFs embedding their sources.
The GME typically as order of pO for a disk component. aud it also hasrandom and halo components.	The GMF typically has order of $\mu {\rm G}$ for a disk component, and it also hasrandom and halo components.
In principle. p, cau be estimated by equation (6)) if the time spread. by these Ποια» and the ECAIF in voids can be well estimated.	In principle, $\rho_s$ can be estimated by equation \ref{eq:estrate}) ) if the time spread by these fields and the EGMF in voids can be well estimated.
However. the ECGME in voids is lighly uncertain as discussed in the last subsection. but could contribute o the total time spread significantly because of large xopagation distance compared to the size of Galactic space and the magnetic structures around sources.	However, the EGMF in voids is highly uncertain as discussed in the last subsection, but could contribute to the total time spread significantly because of large propagation distance compared to the size of Galactic space and the magnetic structures around sources.
This uucertainty leads to a finite range of allowed values of py.	This uncertainty leads to a finite range of allowed values of $\rho_s$.
Given inevitable contributions of the CME aud ECGMES clubedding UITECT. sources to the apparent duration. TZuinCE). aud the allowed maximal time spread incliding the contribution from the poorly known EGALIF iu voids. ΤικE). the rate of UITECT bursts is linuted as Tere. Ες). cau be. i principle. estimated from. anisotropy iu the arrival distribution of VITECRs2009).. assuming that the time spread is loneer than the VITECR observation timescale.	Given inevitable contributions of the GMF and EGMFs embedding UHECR sources to the apparent duration, $\tau_{\rm min}(E)$, and the allowed maximal time spread including the contribution from the poorly known EGMF in voids, $\tau_{\rm max}(E)$, the rate of UHECR bursts is limited as Here, $n_s(E)$ can be, in principle, estimated from anisotropy in the arrival distribution of UHECRs, assuming that the time spread is longer than the UHECR observation timescale.
However, one should keep in munud that equation (6)) is valid when cach UITECR. burst. can be inclividually identified as a burst2009).	However, one should keep in mind that equation \ref{eq:estrate}) ) is valid when each UHECR burst can be individually identified as a burst.
. If more than one bursts or flares occurring in au angular patch contribute to UMECRs observed in the same time-window. ie. the time profiles of two independent CHECR bursts from the same direction (within the size of the augular patch) are overlapped at the Earth. equation (6)) cannot be used as it is.	If more than one bursts or flares occurring in an angular patch contribute to UHECRs observed in the same time-window, i.e., the time profiles of two independent UHECR bursts from the same direction (within the size of the angular patch) are overlapped at the Earth, equation \ref{eq:estrate}) ) cannot be used as it is.

Indeed. the cwarl galaxies of the Local Group. which experienced a variety of SELL. always show stars on the red. plume.

Indeed, the dwarf galaxies of the Local Group, which experienced a variety of SFHs, always show stars on the red plume.

The CALDs of Local Group cwarfs exhibit either the red. plume only (dSph). a weak blue plume as well as a red. plume (clLir/dSph). or a strong blue plume combined with a red plume (dli).

The CMDs of Local Group dwarfs exhibit either the red plume only (dSph), a weak blue plume as well as a red plume (dIrr/dSph), or a strong blue plume combined with a red plume (dIrr).

οσο morphologics serve as a guide to our exploration of ονο CMDs.

These morphologies serve as a guide to our exploration of CHVC CMDs.

We display in Figure 4 the V-L I] CMD of the transition αςαρ galaxy. Phoenix.

We display in Figure 4 the [V-I, I] CMD of the transition dIrr/dSph galaxy Phoenix.

The observations were obtained with the VEEP UT2. "IXOYIEN and FORS2 under good secing conditions.

The observations were obtained with the VLT UT2 “KOYEN" and FORS2 under good seeing conditions.

Phe exposure times were 3 300s each. in V and Lb. Phe data were reduced. and single-star photometry applied in the same wav. as described above for the CLIVCS.

The exposure times were 3 $\times$ 300s each, in V and I. The data were reduced, and single-star photometry applied in the same way, as described above for the CHVCs.

We measure a EWILM in V of 07554. and in L of 0777.

We measure a FWHM in V of 54, and in I of 7.

Ehe total number of stars detected in. V. and Eis S780. much higher iui in anv of the. CIIVC. fields.

The total number of stars detected in V and I is 8780, much higher than in any of the CHVC fields.

Phe DAOPLIOT photometric errors are. clisplavecl in Figure 5. in order to emphasise that the quality of these cata is similar to that of the CLIVC data sets.

The DAOPHOT photometric errors are displayed in Figure 5, in order to emphasise that the quality of these data is similar to that of the CHVC data sets.

Phoenix. is located at. high Galactic latitude. near 69°.

Phoenix is located at high Galactic latitude, near $^o$.

Phe foreground contamination and extinction are small.

The foreground contamination and extinction are small.

H has a distance modulus of about 23.25 (Mateo. 1998) or a distance of about 450 Ixpc.

It has a distance modulus of about 23.25 (Mateo 1998) or a distance of about 450 Kpc.

The CALD of Phoenix shows why it is classed as having a mixed morphology: it exhibits both a blue. plume of relatively voung stars. and a red plume which includes the ROB.

The CMD of Phoenix shows why it is classed as having a mixed morphology: it exhibits both a blue plume of relatively young stars, and a red plume which includes the RGB.

The stellar distributions on the CALDs of the CLVCs are quite dilferent from that of known Local Group carts.

The stellar distributions on the CMDs of the CHVCs are quite different from that of known Local Group dwarfs.

The morphology of the CLIVC CALDs does not reveal cither of the two plumes.

The morphology of the CHVC CMDs does not reveal either of the two plumes.

Instead. the CALDs of the CLIVCs are reminiscent of other. deep HIST. pointings into the clisk and halo of our Galaxy (e... Richer et al.

Instead, the CMDs of the CHVCs are reminiscent of other, deep HST pointings into the disk and halo of our Galaxy (e.g., Richer et al.

2002).

We overplotted on Fig.

1 theoretical isochrones from the database of Girardi et al. (

2000) for ages of LO and 100 Myr. and for 1 and 10 Civr.

2000) for ages of 10 and 100 Myr, and for 1 and 10 Gyr.

We adopted a metallicity of 1/5 of Solar. as cwarft galaxies tend to have sub-Solar metallicities.

We adopted a metallicity of 1/5 of Solar, as dwarf galaxies tend to have sub-Solar metallicities.

1n general. the elfect of varving metallicity at a given age is that more metal-poor isochrones shift to bluer V-I colors. while more metal-rich ones show redder colors.

In general, the effect of varying metallicity at a given age is that more metal-poor isochrones shift to bluer V-I colors, while more metal-rich ones show redder colors.

This shift in minimal for the voung" isochrones (e.g.. 10 LOO Myr). but noticable for the "old" isochrones (1 10 Cyr).

This shift in minimal for the “young" isochrones (e.g., 10 100 Myr), but noticable for the “old" isochrones (1 10 Gyr).

We show the set with a metallicity of 1/5 of Solar because the 1 10 Gyr isochrones have sullicicntly rec colors to encompass the data clistribution.

We show the set with a metallicity of 1/5 of Solar because the 1 10 Gyr isochrones have sufficiently red colors to encompass the data distribution.

When comparing the data with the “voung™ isochrones. we notice there are virtually no stars observed with V-l< 1. where we expect to find a voung population.

When comparing the data with the “young" isochrones, we notice there are virtually no stars observed with $<$ 1, where we expect to find a young population.

This emphasises the absence of young MS and blue supergiant stars in the CLIVCs.

This emphasises the absence of young MS and blue supergiant stars in the CHVCs.

The brightest’ evolutionary phase of stars with ages larger than about one Civr occurs on the red eiant. branch.

The brightest evolutionary phase of stars with ages larger than about one Gyr occurs on the red giant branch.

Furthermore. the PROB. at M; of 4. is a well calibrated distance indicator (Lee. Freedman Madore 1983). and at a distance of 1 Alpe. it would appear at Iz221.

Furthermore, the TRGB, at $_I$ of –4, is a well calibrated distance indicator (Lee, Freedman Madore 1983), and at a distance of 1 Mpc, it would appear at $\approx$ 21.

We overplotted on Fig.

1 theoretical isochrones with ages of 1 ancl LO Cie.

1 theoretical isochrones with ages of 1 and 10 Gyr.

Note that the Cirardi et al.

Note that the Girardi et al.

isochrones extend. above the RGB into the thermally pulsing ACB phase.

isochrones extend above the RGB into the thermally pulsing AGB phase.

A PROB would have been detected for the entire metallicity range appropriate for cwarl galaxies. to L magnitudes of at least 22.5. or to a distance of up to 2 Alpe. in the four CIIVCS whose CMDs are shown in Fig.

A TRGB would have been detected for the entire metallicity range appropriate for dwarf galaxies, to I magnitudes of at least 22.5, or to a distance of up to 2 Mpc, in the four CHVCs whose CMDs are shown in Fig.

1.

Notice that the absence of stars with V-E« 1 rules out extremely metal-poor ICDs. which can be observed in dwarf Spheroidal galaxies.

Notice that the absence of stars with $<$ 1 rules out extremely metal-poor RGBs, which can be observed in dwarf Spheroidal galaxies.

Since no ROB is obvious. a straightforward conclusion is that there is none. and that these objects are not faint galaxies within the Local Group.

Since no RGB is obvious, a straightforward conclusion is that there is none, and that these objects are not faint galaxies within the Local Group.

The absence of a PROB with Lz21 suggests the hypothesis hat LIIPASS J1712-64 is a faint cwarl ealaxy of the Local Group is similarly unliklev.

The absence of a TRGB with $\leq$ 21 suggests the hypothesis that HIPASS J1712-64 is a faint dwarf galaxy of the Local Group is similarly unlikley.

To illustrate the expected ocations of RGDs in our data if ΤΗΑΡ 1712-64 were at à distance of 3.2 Alpe as proposed by Ixilborn et al. (

To illustrate the expected locations of RGBs in our data if HIPASS J1712-64 were at a distance of 3.2 Mpc as proposed by Kilborn et al. (

2000). we overplotted in Fig.

2000), we overplotted in Fig.

2 the shifted Globular Cluster ridgelines rom da Costa ArcmandrolT (1990) of M 15. Fe/11]—-2.17. GO 6397. 21.91. NL 2. -1.58. NGC 6752. -1.54. NGC 1551. -1.29. and 47 Tuc. -0.71.

2 the shifted Globular Cluster ridgelines from da Costa Armandroff (1990) of M 15, [Fe/H]=-2.17, NGC 6397, -1.91, M 2, -1.58, NGC 6752, -1.54, NGC 1851, -1.29, and 47 Tuc, -0.71.

Even i£ HIDPASS J1712-64 were as ar as 3.2 Mpe. with a corresponding PROB I-magnitude of 23.5. we could. have detected. it if its metallicity were ow enough. although this is clearly at the very limit. of our data.

Even if HIPASS J1712-64 were as far as 3.2 Mpc, with a corresponding TRGB I-magnitude of 23.5, we could have detected it if its metallicity were low enough, although this is clearly at the very limit of our data.

Lewis et al. (

2002) recently presented. single-star whotometry obtained under sub-aresee secing conditions in a 0755 [field of view centered at 17 12 35.7. 64 38 12 J2000) and reaching magnitudes fainter than r—24 and e=25.

2002) recently presented single-star photometry obtained under sub-arcsec seeing conditions in a $\times$ 5 field of view centered at 17 12 35.7, –64 38 12 (J2000) and reaching magnitudes fainter than r=24 and g=25.

They. too. Failed to detect an obvious stellar content associated with LIEPASS J1712-64.

They, too, failed to detect an obvious stellar content associated with HIPASS J1712-64.

Furthermore. they rule out any uncderlving unresolved dilfuse components within their large field of view to deep surface-brightness limits.

Furthermore, they rule out any underlying unresolved diffuse components within their large field of view to deep surface-brightness limits.

We can set upper limits on the non-detection of a possible οΝΟ stellar. population in the VET/FORS data. by constructing CALDs and luminosity functions (LE) with ποσο sub-populations.

We can set upper limits on the non-detection of a possible CHVC stellar population in the VLT/FORS data, by constructing CMDs and luminosity functions (LF) with added sub-populations.

This is accomplished by using the VLT FORS2 observation of Phoenix obtained. under very similar conditions and with similar photometric limits. as a template stellar population with which we 7contaminate" the observed. CLIVCG CMDs.

This is accomplished by using the VLT FORS2 observation of Phoenix obtained under very similar conditions and with similar photometric limits, as a template stellar population with which we “contaminate" the observed CHVC CMDs.

We added: Phoenix stars to the IVC οος at their location in the Phoenix €MD.

We added Phoenix stars to the HVC CMDs at their location in the Phoenix CMD.

ln other words. we added

In other words, we added

during the day. even on days when the general RFI level should be quite low (e.g.. the Easter holiday).

during the day, even on days when the general RFI level should be quite low (e.g., the Easter holiday).

Dipoles deployed at ground level should suppress RFI sources located near the horizon and should enable more sensitive RRL (Ellingson2009). (Kogan&

Dipoles deployed at ground level should suppress RFI sources located near the horizon and should enable more sensitive RRL \citep{ellingson09}. \citep{kc09},

Cohen.2009)... 4— 5c 107 The Η15>265(115MHz «v190MHz). (Payneetal..1959).

\ref{tab:lwadetect} $5\sigma$ $10^{-3}$ The $11.5> z > 6.5$ $115\,\mathrm{MHz} < \nu < 190\,\mathrm{MHz}$ \citep{1989ApJ...341..890P}.

1077 21" ?2))

$10^{-3}$ $21^{\mathrm{st}}$ \ref{sec:cosmo})

1077 21" ?2)).

$10^{-3}$ $21^{\mathrm{st}}$ \ref{sec:cosmo})

Since the advent. of high qualiy CCDs there has been a great deal of effort to obtain (inillimagnuitude. or less than 15€)) time-series. optical aud infrared photometry lor laree uumbers of stars.

Since the advent of high quality CCDs there has been a great deal of effort to obtain high-precision (millimagnitude, or less than ) time-series, optical and infrared photometry for large numbers of stars.

Photometry at te inillimagnitude leve of precision is considered a prerequisite in searches [or trausiting Jupiter-sized planets arouud sdlar-type stars.

Photometry at the millimagnitude level of precision is considered a prerequisite in searches for transiting Jupiter-sized planets around solar-type stars.

As a result. the number oL groups that have achieved. this level of precisiou is oo mauy to list here (see for example Horne2003 for a list of trausitit& planet searches).

As a result, the number of groups that have achieved this level of precision is too many to list here (see for example \citealt{horne03} for a list of transiting planet searches).

Tie search [or transiting planets has uot been in vain: at the time of writing there are 7 known ransitiug planets. 6 of which were first identiliel photometrically (see e.g. Udalskietal.2001 αιd Ixouackietal. 2004)).

The search for transiting planets has not been in vain; at the time of writing there are 7 known transiting planets, 6 of which were first identified photometrically (see e.g. \citealt{udalski04} and \citealt{konacki04}) ).

Milliniagnuitude photometry bas also contributed to the study of stars near the hydrogen fusion limit

Millimagnitude photometry has also contributed to the study of stars near the hydrogen fusion limit

especially near (he centre of the svstem.

especially near the centre of the system.

Of course the svstem may not be fully relaxed when the black hole is growing.

Of course the system may not be fully relaxed when the black hole is growing.

Results are shown in Figure 1. [or black hole masses from zero through 0.001. 0.01. 0.1 and 1.0 as labelled by increasing, effect.

Results are shown in Figure \ref{fig:iso} for black hole masses from zero through $0.001$ , $0.01$, $0.1$ and $1.0$ as labelled by increasing effect.

These dimensionless masses correspond to plivsical masses in the range of about 10* to LOM AL. if we use (he values of py and ry given in section yoQu.

These dimensionless masses correspond to physical masses in the range of about $10^7$ to $10^{10}$ $M_{\odot}$ if we use the values of $\rho_0$ and $r_0$ given in section \ref{sec:growth}.

We find the same characteristic density cusp of 2/7?> as obtained by Young(1980).. and the radial velocity profiles are also similar.

We find the same characteristic density cusp of $R^{-3/2}$ as obtained by \citet{you80}, and the radial velocity profiles are also similar.

Although Young(L980) did not show results for the anisotropy parameter. defined as 3=1—<Vg>/«2Vg9>. our results do agree with the prediction of Quinlanefaf(1995). lor distributions of this (wpe.

Although \citet{you80} did not show results for the anisotropy parameter, defined as $\beta = 1 - <V_T^2> / <2 V_R^2>$, our results do agree with the prediction of \citet{qui95} for distributions of this type.

We observe that (here is at most about a 106 anisotropy in favour of tangential motion. due to the increasing binding energv of a particle aud the resultant decrease in eccentricity al constant. angular momentum.

We observe that there is at most about a $10$ anisotropy in favour of tangential motion, due to the increasing binding energy of a particle and the resultant decrease in eccentricity at constant angular momentum.

Moreover (he perturbation extends as far as the core radius only when (he mass is comparable to that of the core. as is to be expected.

Moreover the perturbation extends as far as the core radius only when the mass is comparable to that of the core, as is to be expected.

These results are (he main (test of our progralni.

These results are the main test of our program.

We explore here the aclabatic growth of a central black hole in a collisionless dark matter halo that possesses a central cusp.

We explore here the adiabatic growth of a central black hole in a collisionless dark matter halo that possesses a central cusp.

For initial svstems. we choose (wo very dillerent starting points.

For initial systems, we choose two very different starting points.

The first is an isotropic sell-similar svstem. meant (o represent the final state of a halo undergoing self-similar relaxation (Ilenriksen&Widrow1995).

The first is an isotropic self-similar system, meant to represent the final state of a halo undergoing self-similar relaxation \citep{hen95}.

. The other svstem is the NEW system of Navarro.Frenk&White(1996).. which fits a wide range of dark matter halo sizes.

The other system is the NFW system of \citet{nav96}, which fits a wide range of dark matter halo sizes.

The self-similar svstem is described by the DF where 6 is a[ree parameter (2/3<901.0> 1) that essentially controls (he logarithmic slope of the initial density and potential. given by

The self-similar system is described by the DF where $\delta$ is afree parameter $2/3 < \delta < 1,\ \delta > 1$ ) that essentially controls the logarithmic slope of the initial density and potential, given by

keV band is 9.241012P cre 1 " ↸⊳⋯−↸⊳∪↥⋅↥⋅↸∖↴∖↴↻∪↕∐∐∐∶↴∙↑∪⋅⋝ a huninosty of L1 « 10/9 ere |.

keV band is $\times10^{-12}$ erg $^{-1}$ $^{-2}$ corresponding to a luminosity of 1.1 $\times$ $^{40}$ erg $^{-1}$.

The contribulon of t1 power-law component. the thermal black body coll)OLnont. aud the thermal thin plasma comiponenu to the total O.3-8.0 keV flux are 77C. 16% aud respectiveY.

The contribution of the power-law component, the thermal black body component, and the thermal thin plasma component to the total 0.3-8.0 keV flux are 77, 16 and, respectively.

However. a longer EPIC-PN exposure Is nieed to discutangle a possible extended ticrinal thin plasnua ficun the »)ut-like power-law and thermal tick plasna couponenuts.

However, a longer EPIC-PN exposure is needed to disentangle a possible extended thermal thin plasma from the point-like power-law and thermal thick plasma components.

Tle spectral analysis of the combined NAAI data has coufirmed two spectral compoucuts in Toll N-1.

The spectral analysis of the combined XMM data has confirmed two spectral components in HoII X-1.

The hare COLL2011011 is best described by a powerlaw (D~ 2.6). aud the soft component is fitted by thermal nodels with relatively ow temperatures (AL~0.110.22 keV).

The hard component is best described by a powerlaw $\Gamma\sim2.6$ ), and the soft component is fitted by thermal models with relatively low temperatures $kT\sim0.14-0.22$ keV).

The intrinsic N-rav ΠΟ ΕΕΠΠ isotroplic emission) is about H9 oye tin the 0.3-8.0 keV band

The intrinsic X-ray luminosity (assuming isotrophic emission) is about $^{40}$ erg $^{-1}$ in the 0.3-8.0 keV band.

Before we discuss the nature of the ultraluminuots Νταν source ΠΟ N-1 we stuuarize our resits: There are three niodels for ULX extensively discussed in the literature: 1) ULX 1iuav be black holes of “normal” stellar masses (~ LOAL. ) in binaries. which accrete gas in a supercritical ποσο,

Before we discuss the nature of the ultraluminous X-ray source HoII X-1 we summarize our results: There are three models for ULX extensively discussed in the literature: i) ULX may be black holes of ”normal” stellar masses $\sim 10$ $_{\sun}$ ) in binaries, which accrete gas in a supercritical regime.

They are 33-like objects or mücroquasars m their transicut activity staCs. Wwrose hard radiation can be collimated (and heaued) along Jes ancl accretion cisks axes (Fabrika Moeschervakov. 012001.. Iàiug et al. 20013).

They are 433-like objects or microquasars in their transient activity states, whose hard radiation can be collimated (and beamed) along jets and accretion disks axes (Fabrika Mescheryakov, \cite{Fa01}, King et al., \cite{Ki01}) ).

i) ULX may be black holes wihi ew tens of soar Wass (~ LOAD: ) and that heir ταν enidesbons ds from the disk shining at super-Eddiieton Iuuinosities (Beecluan 20023).

ii) ULX may be black holes with few tens of solar mass $\sim 10$ $_{\sun}$ ) and that their X-ray emissions is from the disk shining at super-Eddington luminosities (Begelman \cite{Beg02}) ).

The slim disk luocl (A»yenmowiez ct al. 1988..

The slim disk model (Abramowicz et al. \cite{Abra88},

Ebisawa et al. 20033)

Ebisawa et al. \cite{Ebi03}) )

alows to explain the observed super-Eddineton hunuinosity. hard A-rav spectra and spectral variations.

allows to explain the observed super-Eddington luminosity, hard X-ray spectra and spectral variations.

ii) The ULN may be iuteriiediate mass black holes (IMBUs). ~ 107M; (€'olbert Mushotzky 19993). which were formed from the very first stars (Aladau Rees 20013) or in globular clusters (Miller ILuuiltou 2002)).

iii) The ULX may be intermediate mass black holes (IMBHs), $\sim 10^3$ $_{\sun}$ (Colbert Mushotzky \cite{Col99}) ), which were formed from the very first stars (Madau Rees \cite{Mad01}) ) or in globular clusters (Miller Hamilton \cite{Mil02}) ).

The IMDBII could. acciTo eas frou a close companion or even from the interstellar medium aud become bright XPav SOULCCS when they are in dense Bas environments.

The IMBH could accrete gas from a close companion or even from the interstellar medium and become bright X--ray sources when they are in dense gas environments.

Spectra of some ULX (Miller et al. 20033)

Spectra of some ULX (Miller et al. \cite{Mil03}) )

support the idea that they ave IMDITs.

support the idea that they are IMBHs.

We call sources “trausicuts” when the iutrinsic duration of UIIECR production at a source 07 is shorter than the characteristic time profile spread. τς) at the observed euergv of £.

We call sources "transients" when the intrinsic duration of UHECR production at a source $\delta T$ is shorter than the characteristic time profile spread $\tau(E)$ at the observed energy of $E$.

Itf the time dispersion is longer thaw the time scale of UITECT. observatious. we nüsperceive that ΤΕΠΟ bursts ave steady sources. aud therefore can define the “apparent” number density of UITECR sources DUE).

If the time dispersion is longer than the time scale of UHECR observations, we misperceive that UHECR bursts are steady sources, and therefore can define the "apparent" number density of UHECR sources $n_s(E)$.

The source muuber deusity is related to the rate of UIIECR bursts p, as The source nuuber deusitv generally depends on UIIECR cnereies. since the apparent duration is dependent on cucreics explicitly.

The source number density is related to the rate of UHECR bursts $\rho_s$ as The source number density generally depends on UHECR energies, since the apparent duration is dependent on energies explicitly.

| UIIECTs observed at the Earth suffer from the GME and EGAIFs embedding their sources.

UHECRs observed at the Earth suffer from the GMF and EGMFs embedding their sources.

The GME typically as order of pO for a disk component. aud it also hasrandom and halo components.

The GMF typically has order of $\mu {\rm G}$ for a disk component, and it also hasrandom and halo components.

In principle. p, cau be estimated by equation (6)) if the time spread. by these Ποια» and the ECAIF in voids can be well estimated.

In principle, $\rho_s$ can be estimated by equation \ref{eq:estrate}) ) if the time spread by these fields and the EGMF in voids can be well estimated.

However. the ECGME in voids is lighly uncertain as discussed in the last subsection. but could contribute o the total time spread significantly because of large xopagation distance compared to the size of Galactic space and the magnetic structures around sources.

However, the EGMF in voids is highly uncertain as discussed in the last subsection, but could contribute to the total time spread significantly because of large propagation distance compared to the size of Galactic space and the magnetic structures around sources.

This uucertainty leads to a finite range of allowed values of py.

This uncertainty leads to a finite range of allowed values of $\rho_s$.

Given inevitable contributions of the CME aud ECGMES clubedding UITECT. sources to the apparent duration. TZuinCE). aud the allowed maximal time spread incliding the contribution from the poorly known EGALIF iu voids. ΤικE). the rate of UITECT bursts is linuted as Tere. Ες). cau be. i principle. estimated from. anisotropy iu the arrival distribution of VITECRs2009).. assuming that the time spread is loneer than the VITECR observation timescale.

Given inevitable contributions of the GMF and EGMFs embedding UHECR sources to the apparent duration, $\tau_{\rm min}(E)$, and the allowed maximal time spread including the contribution from the poorly known EGMF in voids, $\tau_{\rm max}(E)$, the rate of UHECR bursts is limited as Here, $n_s(E)$ can be, in principle, estimated from anisotropy in the arrival distribution of UHECRs, assuming that the time spread is longer than the UHECR observation timescale.

However, one should keep in munud that equation (6)) is valid when cach UITECR. burst. can be inclividually identified as a burst2009).

However, one should keep in mind that equation \ref{eq:estrate}) ) is valid when each UHECR burst can be individually identified as a burst.

. If more than one bursts or flares occurring in au angular patch contribute to UMECRs observed in the same time-window. ie. the time profiles of two independent CHECR bursts from the same direction (within the size of the augular patch) are overlapped at the Earth. equation (6)) cannot be used as it is.

If more than one bursts or flares occurring in an angular patch contribute to UHECRs observed in the same time-window, i.e., the time profiles of two independent UHECR bursts from the same direction (within the size of the angular patch) are overlapped at the Earth, equation \ref{eq:estrate}) ) cannot be used as it is.