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2004. 2006. 2009a: Malbon et al. | 2004, 2006, 2009a; Malbon et al. |
2007: Monaco et al. | 2007; Monaco et al. |
2007: Croton 2009: Cook et al. | 2007; Croton 2009; Cook et al. |
2009). | 2009). |
Quasar clustering provides additional. independent. constraints on the DII population. helping to discriminate among otherwise Viable models. | Quasar clustering provides additional, independent constraints on the BH population, helping to discriminate among otherwise viable models. |
As outlined by Martini Weinberg (2001) and Haiman Hui (2001: sce also Wyithe Loeb 2005: Lidz et al. | As outlined by Martini Weinberg (2001) and Haiman Hui (2001; see also Wyithe Loeb 2005; Lidz et al. |
2006: Hopkins et al. | 2006; Hopkins et al. |
20072: White et al. | 2007a; White et al. |
2008: Shen et al. | 2008; Shen et al. |
2009a.b: Shankar ct al. | 2009a,b; Shankar et al. |
20090: νο Loeb 2009: Bonoli ct al. | 2009c; Wyithe Loeb 2009; Bonoli et al. |
2009). the clustering is an indirect. measure of the masses. and therefore number densities. of the halos hosting the quasars. | 2009), the clustering is an indirect measure of the masses, and therefore number densities, of the halos hosting the quasars. |
In turn. the ratio between the quasar luminosity function and the halo mass function. provides information on the duty evele. ie. the fraction. of halos that host active quasars at a given time. | In turn, the ratio between the quasar luminosity function and the halo mass function provides information on the duty cycle, i.e., the fraction of halos that host active quasars at a given time. |
In general terms. stronger clustering implies that quasars reside in rarer. more | In general terms, stronger clustering implies that quasars reside in rarer, more |
confirm that the intrinsic Lyinan limit drop in the spectra of blue galaxies is al least 3 (0 4 magnitudes. | confirm that the intrinsic Lyman limit drop in the spectra of blue galaxies is at least 3 to 4 magnitudes. |
Even if the Lvman break occupies only à small part of the total filler bandwidth (e.g.. objects at z 2 observed in the U filter). it will be detectable in twpical Lyman break photometric color searches. | Even if the Lyman break occupies only a small part of the total filter bandwidth (e.g., objects at z$\sim$ 2 observed in the U filter), it will be detectable in typical Lyman break photometric color searches. |
We used our observed (lux upper limits and the Bruzual Charlot 107 !| * 700.4. ~ 10?!1072 ! ! *. llaardiandMadau(1996) 10°? ! ! * e «& « zo3-d. faycoe > [44. IHleckmanetal.(2001).. + | We used our observed flux upper limits and the Bruzual Charlot $^{25}$ $^{-1}$$^{-1}$ $^{-3}$ $\AA$ $\sim$ $^{24}$$^{23}$ $^{-1}$ $^{-1}$ $^{-3}$ \citet{haa96} $^{25}$ $^{-1}$ $^{-1}$ $^{-3}$ $\sim$ $<$ $<$ $\sim$ $_{rel~esc}$ $\ge$ $_{rel~esc}$ $\le$ \citet{hec01}, $^{-1}$ |
hampers an accurate treatment of rotational line broadening. | hampers an accurate treatment of rotational line broadening. |
laree distance of NANI 152 above the Galactic Plane. z;2.0 k&kpeford;? kkpe. | large distance of MAXI $-$ 152 above the Galactic Plane, $z>2.0$ kpc for $d>7$ kpc. |
aand ibave discovered a πο DBIIB. MANI 152. | and have discovered a new BHB, MAXI $-$ 152. |
Broadband high-cadeuce unmonitoriug of the outburst revealed the key spectral and timine signatures of DIIDs. iucludiug the presence of time variable QPOs and spectral evolution of the ransieut through canonical states associated with DIIBs. | Broadband high-cadence monitoring of the outburst revealed the key spectral and timing signatures of BHBs, including the presence of time variable QPOs and spectral evolution of the transient through canonical states associated with BHBs. |
ALANI Που 152 has similarities to other ΡΟΠΗΣ. or example Swift 0127. which has a similar outburst lieht-ciurve (Solerietal.2010) and the xesence of a low-temperature disc component. with AT,c0.2 keV. duriug the Hard State (Milleretal. 2006b). | MAXI $-$ 152 has similarities to other BHBs, for example Swift $-$ 0127, which has a similar outburst light-curve \citep{Soleri10} and the presence of a low-temperature disc component, with $kT_\mathrm{in} \simeq 0.2$ keV, during the Hard State \citep{Miller06b}. |
. The increasing flux contribution of the hermal disk component and rise of AZ, to OLS lkkeV sugeests evolution to the Thermal State from he Steep Power-Law State. although no full state ransition was secu before observations euded. | The increasing flux contribution of the thermal disk component and rise of $kT_\mathrm{in}$ to $0.8-1$ keV suggests evolution to the Thermal State from the Steep Power-Law State, although no full state transition was seen before observations ended. |
As the disk compoucut becomes more prominent aud hotter. the radius of the iuner disk edee of the accretion disk shrinks. approximately following the expected Adi,«XΠο: dependence predicted by Cabanacetal. (2009). | As the disk component becomes more prominent and hotter, the radius of the inner disk edge of the accretion disk shrinks, approximately following the expected $kT_\mathrm{in} \propto R_\mathrm{in}^{-3/4}$ dependence predicted by \cite{Cabanac09}. . |
. During the first dav of observations. we recorded a rapid increase m the fitted Αμ. likely caused by localized absorption from an evolving disk wind. | During the first day of observations, we recorded a rapid increase in the fitted $N_\mathrm{H}$, likely caused by localized absorption from an evolving disk wind. |
Optically. MANI 152 shows a correlated rise with N-rav diving the Ward State with briehtness peaking at the transition between the Tard aud Steep Power-Law states. followed by a decrease in the optical brghtuess while iu an intermediate state. | Optically, MAXI $-$ 152 shows a correlated rise with X-ray during the Hard State with brightness peaking at the transition between the Hard and Steep Power-Law states, followed by a decrease in the optical brightness while in an intermediate state. |
We detect a ~2.thhow periodicity in the N-rav heht-curve. | We detect a $\sim2.4$ hour periodicity in the X-ray light-curve. |
The structure of the associated dips is hiehlv variable frou orbit to orbit. sugeesting periodic obscuration from the accretion disk or stream. | The structure of the associated dips is highly variable from orbit to orbit, suggesting periodic obscuration from the accretion disk or stream. |
If we assune this is the orbital period. it would make ALANI J1659 152 the shortest orbital period BIB vet known. the previous beiug Swift 0127 at hhours (Zuritaetal.2008).. | If we assume this is the orbital period, it would make MAXI $-$ 152 the shortest orbital period BHB yet known, the previous being Swift $-$ 0127 at hours \citep{Zurita08}. . |
Caven the likely distance of d>6.1. kkpe. aud the high ealactic latitude. it is likely that MAXI J1659. 152. similar to Swift 0127 and other short period BIBs. is a Galactic Talo DIID. | Given the likely distance of $d>6.1$ kpc, and the high galactic latitude, it is likely that MAXI $-$ 152, similar to Swift $-$ 0127 and other short period BHBs, is a Galactic Halo BHB. |
This work is support by NASA eraut NNNIOAISloc. through the GCuest Iuvestigator Program. | This work is support by NASA grant NNX10AK40G, through the Guest Investigator Program. |
PR and VAL acknowledec financial contribution from the agrecment ASLINAF I/000/10/0. | PR and VM acknowledge financial contribution from the agreement ASI-INAF I/009/10/0. |
APB and PAE acknowledge the support of the Uk Space Agency. | APB and PAE acknowledge the support of the UK Space Agency. |
This work mace use of data supplied by the UI Swift Science Data Centre at theUniversity of Leicester. | This work made use of data supplied by the UK Swift Science Data Centre at theUniversity of Leicester. |
We acknowledge the useof public data from the ddata archive. | We acknowledge the useof public data from the data archive. |
stars with RMS < πας aud 7732 with RMS « 10 miuag. | stars with RMS $<$ 1 mmag and 7732 with RMS $<$ 10 mmag. |
The drastic reduction iu the uuuber ol stars with RMS < 1 uunae is uot unexpected as the differiug exposure times between the two datasets results in a different saturation level. | The drastic reduction in the number of stars with RMS $<$ 1 mmag is not unexpected as the differing exposure times between the two datasets results in a different saturation level. |
From Fig. | From Fig. |
3. it is clear that the observed IMS values are consistent. with photon statistics for all but the brightest magnitudes. | \ref{lcstat_separate} it is clear that the observed RMS values are consistent with photon statistics for all but the brightest magnitudes. |
For stars brighter than 2=16 mae there appears to be au additional source of error contributing to the RAIS. | For stars brighter than $R=16$ mag there appears to be an additional source of error contributing to the RMS. |
To determine this coustaut error for each clip ou the second aud third nights we estimated the RMS in maguituces of the jth light curve as: where i i5 the number of images. £5;=1026;7/75 js the flux in ADU of the jth light curve on the ‘th image. z; is the zero-point of the ith image. ij; is the maguitucde of the jth light curve on the 7th image. s; is the effective sky flux of the 7th image. g.ry is the effective gain of the chip. aud c is the coustaut error term for the chip. | To determine this constant error for each chip on the second and third nights we estimated the RMS in magnitudes of the jth light curve as: where $n$ is the number of images, $F_{ji}=10^{2(z_{i}-m_{ji})/5}$ is the flux in ADU of the $j$ th light curve on the $i$ th image, $z_{i}$ is the zero-point of the $i$ th image, $m_{ji}$ is the magnitude of the $j$ th light curve on the $i$ th image, $s_{i}$ is the effective sky flux of the $i$ th image, $g_{eff}$ is the effective gain of the chip, and $c$ is the constant error term for the chip. |
When perlorming PSF fitting the above equation is applicable except that the effective gain is less thau the actual gain. with the exact factor depeuciug ou the shape of the PSF aud the size of the region oue uses to fit the PSF (see eq. [ | When performing PSF fitting the above equation is applicable except that the effective gain is less than the actual gain, with the exact factor depending on the shape of the PSF and the size of the region one uses to fit the PSF (see \citealt{kjeldsen92} eq. [ |
37] for the case of a Gaussian PSF). | 37] for the case of a Gaussian PSF). |
We find values for the effective gain that are typically less than the actual gain of the CCD by a (actor of ~1.7. | We find values for the effective gain that are typically less than the actual gain of the CCD by a factor of $\sim 1.7$. |
We find that ou the second night the constant error terim ranges from 0.56 mag to 2.1 uunag with an average value (over stars) of 1.2 munag. while for the third night the coustaut error term ranges from 0.12 παπάς to 1.6 μπας with an average value of just below 1 minuag. | We find that on the second night the constant error term ranges from 0.56 mmag to 2.4 mmag with an average value (over stars) of 1.2 mmag, while for the third night the constant error term ranges from 0.42 mmag to 1.6 mmag with an average value of just below 1 mmag. |
We also uote that stars faint enough lor the errors to be domiuated by photon statistics have a lower RMS in the secoud welt compared to the third by a factor that is consistent with tle longer exposure time for the second night. | We also note that stars faint enough for the errors to be dominated by photon statistics have a lower RMS in the second night compared to the third by a factor that is consistent with the longer exposure time for the second night. |
When the data for the second aud third nights are combined. the RMS is not increased Lor stars that are below saturation iu both nights (2< 15.5). | When the data for the second and third nights are combined, the RMS is not increased for stars that are below saturation in both nights $R <
15.5$ ). |
This inuplies that there is no substautial systematic offset between the nights and suggests that one may be able to consistently achieve this level of precision for a louger time series campalgu. | This implies that there is no substantial systematic offset between the nights and suggests that one may be able to consistently achieve this level of precision for a longer time series campaign. |
There are a number of possible sources for the observed constant error terim in our photometry. | There are a number of possible sources for the observed constant error term in our photometry. |
The relative error in the photometry (in magnitudes) due to Poisson noise iu the [lat-field is (Ixjeldseu&Frandsen1992):: where Npp=Osgpepg. Qepy ds the ellective area of the PSF in pixels. and err is the total inunber of electrons in the flat-field in one pixel. | The relative error in the photometry (in magnitudes) due to Poisson noise in the flat-field is \citep{kjeldsen92}: where $N_{ff}=\Omega_{eff}e_{ff}$, $\Omega_{eff}$ is the effective area of the PSF in pixels, and $e_{ff}$ is the total number of electrons in the flat-field in one pixel. |
For the third night the combined flat-lield has eg;=3.3:10 electrons/2x2 pixel. and the FWHM ranged from 6 to 12 2x2 pixels. | For the third night the combined flat-field has $e_{ff} = 3.3 \times 10^{5}$ electrons/2x2 pixel, and the FWHM ranged from 6 to 12 2x2 pixels. |
Thus the expected constant error term due to flat-Delkding lies below 0.2 manag. well below the measured coustaut error terms | Thus the expected constant error term due to flat-fielding lies below 0.2 mmag, well below the measured constant error terms |
ANPs. | AXPs. |
However. the cueut data. namely the association of three ANPs with SNRs. taken at face value. secuinely argue for the opposite conclusion. | However, the current data, namely the association of three AXPs with SNRs, taken at face value, seemingly argue for the opposite conclusion. |
We do recognize that his inference is based on a Πα sample: six ANPs. three of have associated SNRs aud at iiost one SNR association or SCARS (namely. the object of this paper). | We do recognize that this inference is based on a small sample: six AXPs, three of have associated SNRs and at most one SNR association for SGRs (namely, the object of this paper). |
Another possibility for differing ecometry is to invoke aree scale twists of a dipole field with the twist anele ine the underlying plivsical parauceter (Thompson.Lyu-ikov.&I&ullzumi 2002). | Another possibility for differing geometry is to invoke large scale twists of a dipole field with the twist angle being the underlying physical parameter \citep*{tlk02}. |
. In this model. the BB flux arises voth from the heating of the surface due to the decav of strong maguctar fields (Thonipsou&Duncan1996:Πο]&Rulkuni1998) as well as heating of the surface x the return current. | In this model, the BB flux arises both from the heating of the surface due to the decay of strong magnetar fields \citep{td96,hk98} as well as heating of the surface by the return current. |
Resonant cvclotrou scattering of hese photous by the magnetosphere is responsible for the PL compoucut. | Resonant cyclotron scattering of these photons by the magnetosphere is responsible for the PL component. |
The twist angle could be the uucderlwing physical pavancter that differentiates ANPs from SGRs. | The twist angle could be the underlying physical parameter that differentiates AXPs from SGRs. |
We refer to Thompsonetal.(2002) for further discussion of this hypothesis. | We refer to \citet{tlk02} for further discussion of this hypothesis. |
As noted in rofsecintroduetioun and also above. there are considerable difficulties in linkine ANPs to SCRs via temporal evolution. | As noted in \\ref{sec:introduction} and also above, there are considerable difficulties in linking AXPs to SGRs via temporal evolution. |
Specifically, the period and period. derivatives of ANDPs and SCGRs overlap and in are strouely clustered. | Specifically, the period and period derivatives of AXPs and SGRs overlap and in are strongly clustered. |
Thus. the simplest interpretation of the overlap of propertics is that ANPs and SCRs are similar objects but in differing “states”. | Thus, the simplest interpretation of the overlap of properties is that AXPs and SGRs are similar objects but in differing “states”. |
As an example. we note that j)behaved like ai classical SCOR from. its discovery du 1979 until 1983. but has been sileut since then and this may account why the current spectral properties of aure similar to those of ΑΝΤΙΣ. | As an example, we note that behaved like a classical SGR from its discovery in 1979 until 1983, but has been silent since then and this may account why the current spectral properties of are similar to those of AXPs. |
We do not know the duty evele of the two states (ANP and SCR). | We do not know the duty cycle of the two states (AXP and SGR). |
I£ maguetars speud a significant fraction of time in the ANP state then the radiated cucrey (assuniue sav 50 eV low enerev cutoff for the PLcompoucut) cau be as lieh as 1.2«10 ore ὃνlobar~LO ore. | If magnetars spend a significant fraction of time in the AXP state then the radiated energy (assuming say 50 eV low energy cutoff for the PLcomponent) can be as high as $1.2\times 10^{37}\,$ erg $^{-1}\times 10^4\,{\rm yr} \sim 3\times 10^{48}\,$ erg. |
The interred B-field value (to supply this euergv) is iu excess of 102? G. As noted in rofsecintroductioun there is erowing evidence for pulsars with strong B-fields. 107<Bx10H C (IIBPSRs). | The inferred $B$ -field value (to supply this energy) is in excess of $10^{15}\,$ G. As noted in \\ref{sec:introduction} there is growing evidence for pulsars with strong $B$ -fields, $10^{13} \ale B \ale 10^{14}\,$ G (HBPSRs). |
Zhang&Tarding(2000) have suggested that neutron stars with Beποτ will not exhibit radio pulsations. | \citet{zh00} have suggested that neutron stars with $B\age
10^{14}\,$ G will not exhibit radio pulsations. |
If so. there niav exist an intermediate eroup of neutron stars with 1015<Bzx MAC which are neither radio pulsars nor πιοος of the ANP|SCR family. | If so, there may exist an intermediate group of neutron stars with $10^{14} \ale B \ale
10^{15}$ G which are neither radio pulsars nor members of the AXP+SGR family. |
Perhaps the ucarby X-rav pulsar RBS 1223 (Tambarvanetal.2002) may be a member of this intermediate eroup. | Perhaps the nearby X-ray pulsar RBS 1223 \citep{hhss2001} may be a member of this intermediate group. |
We have made extensive use of the SIMDBAD database and we are erateful to the astronomers at the Centre de Donnéces Astronomiques de Strasbourg for maintaining this database. | We have made extensive use of the SIMBAD database and we are grateful to the astronomers at the Centre de Donnéees Astronomiques de Strasbourg for maintaining this database. |
We thauk Marten van ERerkwijk. τής Thompson. Audrew Melatos and Binge Zhang for helpful discussions. | We thank Marten van Kerkwijk, Chris Thompson, Andrew Melatos and Bing Zhang for helpful discussions. |
DLE thanks the Faunie aud Johu Wertz Foundation for a fellowship. | DLK thanks the Fannie and John Hertz Foundation for a fellowship. |
Support for this work was provided bv the National Aeronautics aud Space Achuinistration (NASA) through award umuuber COL-2056X issued bv the ταν Observatory Center which is operated by the Sinithsonian Astrophysical Observatory for aud on behalf of NASA under contract NASS-39073. | Support for this work was provided by the National Aeronautics and Space Administration (NASA) through award number GO1-2056X issued by the X-ray Observatory Center which is operated by the Smithsonian Astrophysical Observatory for and on behalf of NASA under contract NAS8-39073. |
at its position iu recent. deep Slo ΛΠΙ observations (INoudviatey et al. | at its position in recent, deep 840 MHz observations (Kondriatev et al. |
2009) would only rule out a raclio-loud ACN, | 2009) would only rule out a radio-loud AGN. |
However. an AGN not detecte in would be either very much absorbed or at a much higher redshift. which would be hardly compatible with the constraiut on its (B2) and with its detection iu the D hance. respectively, | However, an AGN not detected in would be either very much absorbed or at a much higher redshift, which would be hardly compatible with the constraint on its $(B-R)$ and with its detection in the B band, respectively. |
On the other haud. à (85.R)zi0.6 would be compatible with a iain sequence star of spectral tvpe earlier than F but this should be well outside the Galaxy to reproduce the observed. B-baud Hux. | On the other hand, a $(B-R)\la0.6$ would be compatible with a main sequence star of spectral type earlier than F but this should be well outside the Galaxy to reproduce the observed B-band flux. |
Thus. both the counterpart colour aud its brightuess are only compatible with a neutron star. | Thus, both the counterpart colour and its brightness are only compatible with a neutron star. |
We compared our AZ-baud magnitude upper luit with the vest-fit model to the sspectrmu of 177 land with the available multi-band photometry (Rea et al. | We compared our $R$ -band magnitude upper limit with the best-fit model to the spectrum of 1774 and with the available multi-band photometry (Rea et al. |
2007: Zane et al. | 2007; Zane et al. |
2008: Sclavope et al | 2008; Schwope et al. |
2009). | 2009). |
We computed the interstellar extinction correction from the hydrogen column density Vy, derived roni the X-rav spectral fit. | We computed the interstellar extinction correction from the hydrogen column density $N_H$ derived from the X-ray spectral fit. |
To this aim. we fitted the spectrin anew using updated calibrations and respo-SC files, | To this aim, we fitted the spectrum anew using updated calibrations and response files. |
We re-extracted the EPIC-pu spectrui from the original λίαν 2001 observation using version 10.0.0 of the SScicuce Analysis System (SAS). a circular region of 20” radius. auc selecting only sinegle-pixel events. | We re-extracted the EPIC-pn spectrum from the original May 2004 observation using version 10.0.0 of the Science Analysis System (SAS), a circular region of $20\arcsec$ radius, and selecting only single-pixel events. |
Since colmbined fits with pu aud MOS spectra are dominated bv the higher statistical quality of the pu spectrum. we concentrated on the pu spectrum only. | Since combined fits with pn and MOS spectra are dominated by the higher statistical quality of the pn spectrum, we concentrated on the pn spectrum only. |
For conrparisou with the results obtained by Sclwope ct al. ( | For comparison with the results obtained by Schwope et al. ( |
2009) and Cropper et al. ( | 2009) and Cropper et al. ( |
2007) we fittec the spectrin with either a lack body (BB) model. a BB plus au (additive) Caussiau inc. and a BB plus a (iuultiplicative) absorption edge. | 2007) we fitted the spectrum with either a black body (BB) model, a BB plus an (additive) Gaussian line, and a BB plus a (multiplicative) absorption edge. |
Iu cach case. interstellar absorption was included by using he (version 12.6.0X) uodoel with abundances Yoni Wilts et al. ( | In each case, interstellar absorption was included by using the (version 12.6.0k) model with abundances from Wilms et al. ( |
2000), | 2000). |
The best fi parameters are sunimarsed in Table 1. | The best fit parameters are summarised in Table 1. |
As in previous analyses, inclusion of an absorption feature nuproves the fit over a simple DD 1nodel. although the current energy resolution docs rot allow to discriminate a Gaussian Lue iu absorption roni an edge. | As in previous analyses, inclusion of an absorption feature improves the fit over a simple BB model, although the current energy resolution does not allow to discriminate a Gaussian line in absorption from an edge. |
The slight differences between the model »uanmeters reported here aud those given by Sclavope et al. ( | The slight differences between the model parameters reported here and those given by Schwope et al. ( |
2009) and Cropper et al. ( | 2009) and Cropper et al. ( |
2007) are larecly due ο the different revisious of the soft spectral response calibration of the EPIC-pu camera aud possibly dciffereut elemental abuucances used for the absorption model. | 2007) are largely due to the different revisions of the soft spectral response calibration of the EPIC-pn camera and possibly different elemental abundances used for the absorption model. |
The ιο spectral fit with an absorbed BB plus an (additive) Cassia ine vields a best-fit temperature of AD=LoL2+2.1. eV. cousistent with that οΤαλος by Zane et al. ( | The new spectral fit with an absorbed BB plus an (additive) Gaussian line yields a best-fit temperature of $kT =
104.2 \pm 2.1$ eV, consistent with that obtained by Zane et al. ( |
2005) x fitting the same data but with older calibrations anc response files. | 2005) by fitting the same data but with older calibrations and response files. |
However. we note that our new spectral fit vields a best-fit Ny value of (2.640.2)&1029 2 which is somewhat lower han that obtained by Zane ct al. ( | However, we note that our new spectral fit yields a best-fit $N_H$ value of $(2.6 \pm 0.2) \times 10^{20}$ $^{-2}$ which is somewhat lower than that obtained by Zane et al. ( |
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