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

source stringlengths 1 2.05k ⌀	target stringlengths 1 11.7k
. The sources of error hat may contribute to uncertainties tu £j4; aud consequent age are indeed utimerous aud include uncertainties in: adopted distance. spectral type and extinction. photometric variability. unresolved companions. aud accretion luminosity.	The sources of error that may contribute to uncertainties in $L_{bol}$ and consequent age are indeed numerous and include uncertainties in: adopted distance, spectral type and extinction, photometric variability, unresolved companions, and accretion luminosity.
Ou the other haud. if oue were to concede a real age spread in the ONC (plausible as star formation is still talkiug place at the present day). stellar ages 1might Carry a statistical aud relative significauce. though still uncertain on an individual basis: stars tliat we place lower on the Hayashi tracks may be older. on average. than those that we place higher.	On the other hand, if one were to concede a real age spread in the ONC (plausible as star formation is still taking place at the present day), stellar ages might carry a statistical and relative significance, though still uncertain on an individual basis: stars that we place lower on the Hayashi tracks may be older, on average, than those that we place higher.
This latter hypotliesis is supported by the time evolution of ONC stellar radii recently observed » Rhodeοἱal.(2001).	This latter hypothesis is supported by the time evolution of ONC stellar radii recently observed by \citet{rho01}.
. Although our correlation of Ly with age also seems to argue in favor of a ‘eal age spread. it could also be an artifact of au interrelation between inferred) Ly aud L4. aud he matter remains open.	Although our correlation of $L_X$ with age also seems to argue in favor of a real age spread, it could also be an artifact of an interrelation between inferred $L_X$ and $L_{bol}$, and the matter remains open.
Indeed the treud in Figure 8. corresponds to a constant ratio of X-ray to iometric lumiuosity at different ages aud cau be interpreted equally well in two different. ways: 1) Stars of equal mass have equal bolometric αμα X-ray luminosities. but there are effects that act in exactly the same way on both Ly auc Lo: 2) These stars are saturated and stay saturated hrough their contraction on the Hayashi tracks. thus keeping La/Lig constant.	Indeed the trend in Figure \ref{fig:LXvsA0a} corresponds to a constant ratio of X-ray to bolometric luminosity at different ages and can be interpreted equally well in two different ways: 1) Stars of equal mass have equal bolometric and X-ray luminosities, but there are effects that act in exactly the same way on both $L_X$ and $L_{bol}$ ; 2) These stars are saturated and stay saturated through their contraction on the Hayashi tracks, thus keeping $L_X/L_{bol}$ constant.
Figure 9 presents a plot of ουν) vs. [ουσε) [οι stars iu the 2—3AL. mass bin ancl 'eveals an age clepeucdence aud a sudden drop of Ly at LogAye)6.5 (a similar drop is also seen in Lxy/Ly4).	Figure \ref{fig:LXvsA0b} presents a plot of $Log(L_X)$ vs. $Log(Age)$ for stars in the $2-3~M_{\odot}$ mass bin and reveals an age dependence and a sudden drop of $L_X$ at $Log(Age) \sim 6.5$ (a similar drop is also seen in $L_X/L_{bol}$ ).
This is just the age at which a 2.5AL. star dissipates its convective envelope accordingo the SDF modclels. leudiug further evidence for the convection —activity counection found iu	This is just the age at which a $2.5~M_{\odot}$ star dissipates its convective envelope accordingto the SDF models, lending further evidence for the convection –activity connection found in
where r is the timescale over which the mass is lost.	where $\tau$ is the timescale over which the mass is lost.
The magnetic energv density is equal to the kinetic energy density at the Alfvénn radius which we find to be We consider this further in Section 4. where we apply our model to the recurrent nova U Sco.	The magnetic energy density is equal to the kinetic energy density at the Alfvénn radius which we find to be We consider this further in Section \ref{usco} where we apply our model to the recurrent nova U Sco.
We parametrise the dipole moment with the magnetic field strength at the stellar surface. Da. so that where fs is the stellar radius of the companion which woe find with equation (13)).	We parametrise the dipole moment with the magnetic field strength at the stellar surface, $B_{\rm star}$, so that where $R_2$ is the stellar radius of the companion which we find with equation \ref{rl}) ).
The specific angular momentum of the ejectec mass that is forced to corotate with the secondary is (63\|AUG)0 assuming that the secondary is tically locked.	The specific angular momentum of the ejected mass that is forced to corotate with the secondary is $(a_2^2+KR_{\rm A}^2)\,\Omega$ assuming that the secondary is tidally locked.
The distance from the centre of mass to the secondary star is The constant A depends on the distribution of the material within the Alfvénn radius.	The distance from the centre of mass to the secondary star is The constant $K$ depends on the distribution of the material within the Alfvénn radius.
ΕΕ it were distributed uniformly within a spherical shell. then A.=22/3.	If it were distributed uniformly within a spherical shell, then $K=2/3$.
Because the density within the shell varies and the shell itsell will not be perfectly spherical. we take fy=I.	Because the density within the shell varies and the shell itself will not be perfectly spherical, we take $K=1$.
Thus. the angular momentum. loss from the binary to the ejected material is eiven by We estimate the fraction of the ejected mass that gains angular momentum in this wav to be Now with equations (2)) : We substitute this into equation (6)) with AM,=Am, and AAI=0 to find the change in separation We note that this reduces to equation (7)) when there is no magnetic field.	Thus, the angular momentum loss from the binary to the ejected material is given by We estimate the fraction of the ejected mass that gains angular momentum in this way to be Now with equations \ref{da}) ) and \ref{jobs}) ) we find We substitute this into equation \ref{change}) ) with $\Delta M_1=-\Delta m_1$ and $\Delta M_2=0$ to find the change in separation We note that this reduces to equation \ref{over}) ) when there is no magnetic field.
In Fig.	In Fig.
2 we plot APP as a function of the mass ratio for cillerent values of the πμAlfvénn radius.	\ref{dadm2} we plot $\frac{\Delta P/P}{\Delta m_1/M}$ as a function of the mass ratio for different values of the Alfvénn radius.
The larger magnetic fields are even capable of producing an overall decrease to the orbital period. during the nova for small mass ratios.	The larger magnetic fields are even capable of producing an overall decrease to the orbital period during the nova for small mass ratios.
In Fig.	In Fig.
I. we also plot MUSmena [orfor anan. Mfvénn.Alfvénn radius of O.75e@ (long dashed line) for comparison with the other models we have considered.	\ref{dadm} we also plot $\frac{\Delta P/P}{\Delta m_1/M}$ for an Alfvénn radius of $0.75\,a$ (long dashed line) for comparison with the other models we have considered.
For this strong magnetic field. the period change is smaller than with mass transfer to the secondary or [rictional angular momentum losses.	For this strong magnetic field, the period change is smaller than with mass transfer to the secondary or frictional angular momentum losses.
Lf a secondary star has a large magnetic field then it will significantIv. alter the orbital period. change during a nova outburst.	If a secondary star has a large magnetic field then it will significantly alter the orbital period change during a nova outburst.
‘There is a critical mass ratio where the additional terni in equation (29)) causes no change to the separation (or orbital period) Ig<qon then the cjeetecd material that is forced to corotate with the secondary gains angular momentum. angular momentum is lost [rom the orbit and so the separation decreases.	There is a critical mass ratio where the additional term in equation \ref{cor}) ) causes no change to the separation (or orbital period) If $q<q_{\rm crit}$ then the ejected material that is forced to corotate with the secondary gains angular momentum, angular momentum is lost from the orbit and so the separation decreases.
Similarly the orbital period change is smaller than previously predicted.	Similarly the orbital period change is smaller than previously predicted.
Llowever. if dder. then the material Forced. to corotate loses angular moment and the orbital period increases.	However, if $q>q_{\rm crit}$, then the material forced to corotate loses angular momentum and the orbital period increases.
We plot this in the solidunm. line in Fie. 3..	We plot this in the solid line in Fig. \ref{qc}.
We note that most observed systems have a mass ratio qd1 (Ritter&Ixolb2003). because there is a critical mass ratio. that depends on AM». above which the mass transfer becomes unstable (e.g.Smith&VanelePutte 2006).	We note that most observed systems have a mass ratio $q\lesssim 1$ \citep{ritter03} because there is a critical mass ratio, that depends on $M_2$, above which the mass transfer becomes unstable \citep[e.g.][]{smith06}.
. Hence in most systems the orbital period will decrease when the effects of à magnetic field are considered.	Hence in most systems the orbital period will decrease when the effects of a magnetic field are considered.
We also plot the dashed line for the mass ratio below which the period change is negative.	We also plot the dashed line for the mass ratio below which the period change is negative.
We see that for the larger magnetic fields it is possible that the orbital period may actually decrease for all mass ratios.	We see that for the larger magnetic fields it is possible that the orbital period may actually decrease for all mass ratios.
This effect. could be significant even for larger mass ratios.	This effect could be significant even for larger mass ratios.
Frictional angular momentum losses only dominate for g=0.01 (Sharactal. so this new mechanism potentially has a greater effect on the orbital period change.	Frictional angular momentum losses only dominate for $q\lesssim 0.01$ \citep{shara86} so this new mechanism potentially has a greater effect on the orbital period change.
There are now ten recurrent novae known in our galaxy	There are now ten recurrent novae known in our galaxy \citep[the tenth one was discovered last year;][]{pagnotta09}
populate the ERO class.	populate the ERO class.
The best examples are provided by LBDS53W001 (ROAN= 5.8 2=1.55: Dunlop et al.	The best examples are provided by LBDS 53W091 $R-K=5.8$ ; $z=1.55$; Dunlop et al.
1996: Spinrad et al.	1996; Spinrad et al.
1997). ERO CL 09039143713D (29.Wy=à: 2o1.6: Soifer et al.	1997), ERO CL 0939+4713B $R-K=7$; $z\sim 1.6$; Soifer et al.
1999). and. by most of the EROs ats~1.3 around the QSO 1213-0017 (Liu et al.	1999), and by most of the EROs at $z\sim 1.3$ around the QSO 1213-0017 (Liu et al.
2000).	2000).
Also near-LR. spectroscopy confirmed that both dusty and old galaxies contribute to the ERO population (Cimatti et al.	Also near-IR spectroscopy confirmed that both dusty and old galaxies contribute to the ERO population (Cimatti et al.
1999).	1999).
Understanding the nature and deriving the abundance of EROs is important to shed light on the controversial Issue of the deficit of high-z elliptical galaxies: according to some results. the number of galaxies with the red colours expected for high-z passively evolved. spheroidals is lower compared to the predictions of passive luminosity evolution. (e.g. Ixaullmann. Charlot White 1996: Zepf 1997: Franceschini et al.	Understanding the nature and deriving the abundance of EROs is important to shed light on the controversial issue of the deficit of $z$ elliptical galaxies: according to some results, the number of galaxies with the red colours expected for $z$ passively evolved spheroidals is lower compared to the predictions of passive luminosity evolution (e.g. Kauffmann, Charlot White 1996; Zepf 1997; Franceschini et al.
1998: Bareer et al.	1998; Barger et al.
1999).	1999).
Llowever. other works cid not confirm the existence of such a deficit up to z~2 (e.g. 'l'otani Yoshii 1997: Benitez ct al.	However, other works did not confirm the existence of such a deficit up to $z\sim 2$ (e.g. Totani Yoshii 1997; Benitez et al.
1999: Broadhurst. Bowens 1999: Schade et al.	1999; Broadhurst Bowens 1999; Schade et al.
1999).	1999).
As a first step to understand the nature of EROs. we started an imaging survey program aimed at selecting complete samples of such galaxies both in "empty fields and around. high-z radio-loud. AGN.	As a first step to understand the nature of EROs, we started an imaging survey program aimed at selecting complete samples of such galaxies both in “empty” fields and around $z$ radio-loud AGN.
In this paper. we present the results of our survey around racdio-Ioud CN at ο>1.5.	In this paper, we present the results of our survey around radio-loud AGN at $z>1.5$.
The main motivations of such a survey are to investigate whether EROs are more numerous around high-z ACN (as suspected. in previous works: e.g. AleCarthy et al.	The main motivations of such a survey are to investigate whether EROs are more numerous around $z$ AGN (as suspected in previous works; e.g. McCarthy et al.
1992: Dev et al.	1992; Dey et al.
1995). to select. old. passively evolving galaxies abozI1. hy weir optical/near-I1t. colours. to study the environment of high-z racio-loud XN (e.g. Hall. Green Cohen 1998: Hall Creen 1998 and references therein). ane to search for galaxy cluster candidates at 5.	1995), to select old passively evolving galaxies at $z>1$ by their optical/near-IR colours, to study the environment of $z$ radio-loud AGN (e.g. Hall, Green Cohen 1998; Hall Green 1998 and references therein), and to search for galaxy cluster candidates at $z>1.5$ .
We recal that the most distant. spectroscopically confirmed. cluster known to date bas z=1.27 (Stanford et al.	We recall that the most distant spectroscopically confirmed cluster known to date has $z=1.27$ (Stanford et al.
LOOT: Rosati et al.	1997; Rosati et al.
1999).	1999).
Other works selected cluster candidates aroun quasars at z1.3 and suggested a significant heterogeneity of the cluster galaxies. including both passively evolving ol ellipticals as well as vounger and custy systems (e.g. Hal Green 1998: Liu et al.	Other works selected cluster candidates around quasars at $z>1.3$ and suggested a significant heterogeneity of the cluster galaxies, including both passively evolving old ellipticals as well as younger and dusty systems (e.g. Hall Green 1998; Liu et al.
2000).	2000).
Phroughout this paper we assume Ly=50 + 1 Oy2d) and QO,=0 unless otherwise statect.	Throughout this paper we assume $H_0=50$ $^{-1}$ $^{-1}$, $\Omega_0=1$ and $\Omega_{\Lambda}=0$ unless otherwise stated.
The survey presented here is based on 2 and. A7 -baux imagine.	The survey presented here is based on $R$ - and $K^{\prime}$ -band imaging.
We observed totally 12. fields: 6 around. racio galaxies (Gs) taken from the ALRC sample selected at 408 Az (AleCarthy et al.	We observed totally 14 fields: 6 around radio galaxies (RGs) taken from the MRC sample selected at 408 MHz (McCarthy et al.
1997. IxXapahii ct al.	1997, Kapahi et al.
1998). S arounc raclio-loud quasars (ILOs) taken from the PIS) sample selected at 2.7 Gllz (Wright Ostrupeek 1990).	1998), 8 around radio-loud quasars (RLQs) taken from the PKS sample selected at 2.7 GHz (Wright Ostrupcek 1990).
We note that AIRC101-220 is a broad line radio galaxy (Ixapahi e al.	We note that MRC1017-220 is a broad line radio galaxy (Kapahi et al.
1998).	1998).
All the targeted AGN have 1.5«z2.0. with the exception of the quasars PIS 1351-018 (2= 3.71) ane PINS 1556-245 (2= 2.82).	All the targeted AGN have $1.5<z<2.0$, with the exception of the quasars PKS 1351-018 $z=3.71$ ) and PKS 1556-245 $z=2.82$ ).
The radio powers at rest-frame 5 CllzrangeΓΕ for RGs to 430.1077-- for RLOs.	The radio powers at rest-frame 5 GHz range from $2-5 \times 10^{27}$ $^{-1}$ for RGs to $4-30 \times 10^{27}$ $^{-1}$ for RLQs.
Table 1 lists the relevant information about the sample.	Table 1 lists the relevant information about the sample.
“Phe fields were selected according to the redshifts of the AGN. to their Galactic latitude (>207) and to their good observability curing the telescope runs.	The fields were selected according to the redshifts of the AGN, to their Galactic latitude $>20^{\circ}$ ) and to their good observability during the telescope runs.
We have used the predictions of the Bruzual Charlot (1999) evolutionary svnthesis models to define a colour selection. criterion capable to select a complete sample of EROs at high-z.	We have used the predictions of the Bruzual Charlot (1999) evolutionary synthesis models to define a colour selection criterion capable to select a complete sample of EROs at $z$.
In particular. we used instantaneous burst models (also called simple stellar population models. SSP) with cillerent redshifts of formation (2,=2.3.4.5.6) o describe old. spheroidal galaxies.	In particular, we used instantaneous burst models (also called simple stellar population models, SSP) with different redshifts of formation $z_{f}=2,3,4,5,6$ ) to describe old spheroidal galaxies.
In addition. we also considered. models with exponential time-scale for the star ormation rate. SLRxecp(F/r7). with 7=0.1.0.3 Ce or QO=1.0.1 respectively.	In addition, we also considered models with exponential time-scale for the star formation rate, $SFR \propto exp(-t/\tau)$, with $\tau=0.1,0.3$ Gyr for $\Omega= 1,0.1$ respectively.
Such models. correspond. to cases with low residual star formation at 2< (ic. <1 M.vr \| for a galaxy with mass Maur=10 M. ). and hey are capable to reproduce the colours and the spectra of local ellipticals. as well as the faint galaxy optical colour and vecshift distributions (Pozzetti. Bruzual Zamorani 1996).	Such models correspond to cases with low residual star formation at $z<2$ (i.e. $<1$ $_{\odot}$ $^{-1}$ for a galaxy with mass $_{gal}=10^{11}$ $_{\odot}$ ), and they are capable to reproduce the colours and the spectra of local ellipticals, as well as the faint galaxy optical colour and redshift distributions (Pozzetti, Bruzual Zamorani 1996).
A Salpeter IME (0.1.<a<125 M.) and. solar metallicity have been adopted.	A Salpeter IMF $0.1<m<125$ $_{\odot}$ ) and solar metallicity have been adopted.
Models with Scalo IME or with supersolar metallicities reach even redder colours at high-:.	Models with Scalo IMF or with supersolar metallicities reach even redder colours at $z$.
Optical to near-Ht AA. colours derived from the acloptecl evolutionary models are shown in Fig.	Optical to near-IR $R-K$ colours derived from the adopted evolutionary models are shown in Fig.
1 for two cosmologies (44,=50 + \|. Q—0.1).	1 for two cosmologies $H_0=50$ $^{-1}$ $^{-1}$, $\Omega=0.1,1$ ).
The colour selection threshold adopted in our survey is Rolv>6.	The colour selection threshold adopted in our survey is $R-K>6$.
This allows us to select old galaxies at z>1.2 [ormed at zr;23 in both cosmologics.	This allows us to select old galaxies at $z>1.2$ formed at $z_{f}>3$ in both cosmologies.
Ht is relevant to note that for τ 20.3 Gyr and zy<3. colours 2.AN>6 are never reached.	It is relevant to note that for $\tau>$ 0.3 Gyr and $z_{f}<3$, colours $R-K>6$ are never reached.
In other words. colours 2A6 select old. galaxies whieh had a short episode of star formation in carly cosmological epochs.	In other words, colours $R-K>6$ select old galaxies which had a short episode of star formation in early cosmological epochs.
For instance. assuming no cust extinction. a galaxy at z 21.5 with RoA>6 should have SpIOX. an age C2 Gaver (04= 1.0) or 3 Gyr (04=0.1) and SER<<1 M.vyr. +.	For instance, assuming no dust extinction, a galaxy at $z\approx$ 1.5 with $R-K>6$ should have $z_{f}>3$, an age $>$ 2 Gyr $\Omega_0=1.0$ ) or $>$ 3 Gyr $\Omega_0=0.1$ ) and $SFR<<1$ $_{\odot}$ $^{-1}$ .
eear-infrared imagine was done on 1997 April 1-3 with rw ESO/AIPL 2.2m telescope. with the HUXC2D. camera (Moorwood ct al.	Near-infrared imaging was done on 1997 April 1-3 with the ESO/MPI 2.2m telescope with the IRAC2B camera (Moorwood et al.
1992) equipped with a 256 11οςαΤο array (0.506" /pixel).	1992) equipped with a $\times$ 256 HgCdTe array $^{\prime \prime}$ /pixel).
We used the A" filter in order to reduce 1e thermal noise.	We used the $K^{\prime}$ filter in order to reduce the thermal noise.
The sky conditions were photometric and 10 seeing. was around 1.07. o	The sky conditions were photometric and the seeing was around $^{\prime \prime}$.
sPhe observations. were done aking a number of background-limited. images (typically 10-14) with the telescope moved LO” between cach image.	The observations were done taking a number of background-limited images (typically 10-14) with the telescope moved $10^{\prime \prime}$ between each image.
In each telescope position. theexposures were typically 120 seconds long (c.g. 12 images cach with an exposuretime of 10 seconds).	In each telescope position, theexposures were typically 120 seconds long (e.g. 12 images each with an exposuretime of 10 seconds).
“Phe data reduction was performed. using the ΗΛ reduction package and using the method outlined by Villani di Serego Alighieri (1099).	The data reduction was performed using the IRAF reduction package and using the method outlined by Villani di Serego Alighieri (1999).
Photometric calibration was obtained observing standard. stars from the Carter Meadows. (1995). sample.	Photometric calibration was obtained observing standard stars from the Carter Meadows (1995) sample.
Phe typical night-to-night scatter in the zero-points was around 0.03 magnitudes.	The typical night-to-night scatter in the zero-points was around 0.03 magnitudes.
The conversion from A" to A magnitudes inthe SAAO-Carter	The conversion from $K^{\prime}$ to $K$ magnitudes inthe SAAO-Carter
showing weak or even abscut thermal features) display very high peak frequencics: in IIBL. which appear to be the less powerful objects (with typical Iuuinosities 107 ere cm? 1) the first peals falls in the UV-soft Naav baud aud the high energv component can peak near the TeV reeion.	showing weak or even absent thermal features) display very high peak frequencies: in HBL, which appear to be the less powerful objects (with typical luminosities $10^{42}$ erg $^{-2}$ $^{-1}$ ) the first peak falls in the UV-soft X-ray band and the high energy component can peak near the TeV region.
This behaviour has been iuterpreted (Cbisellini et al.	This behaviour has been interpreted (Ghisellini et al.
1998. recently revisited in (κοπια et al.	1998, recently revisited in Ghisellini et al.
2002) as the result of a balance. between a acceleration i1 cooling suffered. by the electrons producing the radiation at the peaks.	2002) as the result of a balance between a acceleration an cooling suffered by the electrons producing the radiation at the peaks.
Iu the assumption that the acceleration rate is coustant in all the sources. the sequence ean be successfully interpreted as due to au increasing cooling rate (dominated iu most of the sources by IC losses) from the low-power BL Lacs to the powerful FSRQs.	In the assumption that the acceleration rate is constant in all the sources, the sequence can be successfully interpreted as due to an increasing cooling rate (dominated in most of the sources by IC losses) from the low-power BL Lacs to the powerful FSRQs.
SEDs can be satisfactorily reproduced by simple cussion models (e.g. Clisellini et al.	SEDs can be satisfactorily reproduced by simple emission models (e.g. Ghisellini et al.
1998. Tavecchio et al.	1998, Tavecchio et al.
1998). vielding the value of the main plwvsical quautitics of the jet. such as electron density and cnereyv. magnetic fold. size of the region.	1998), yielding the value of the main physical quantities of the jet, such as electron density and energy, magnetic field, size of the region.
In particular miecasuring the N-ray spectra aud adapting a broad band model to their SEDs vields reliable estimates of the total ΙΟ: of relativistic particles involved. which is dominated by those at the lowest energies.	In particular measuring the X-ray spectra and adapting a broad band model to their SEDs yields reliable estimates of the total number of relativistic particles involved, which is dominated by those at the lowest energies.

. The sources of error hat may contribute to uncertainties tu £j4; aud consequent age are indeed utimerous aud include uncertainties in: adopted distance. spectral type and extinction. photometric variability. unresolved companions. aud accretion luminosity.

The sources of error that may contribute to uncertainties in $L_{bol}$ and consequent age are indeed numerous and include uncertainties in: adopted distance, spectral type and extinction, photometric variability, unresolved companions, and accretion luminosity.

Ou the other haud. if oue were to concede a real age spread in the ONC (plausible as star formation is still talkiug place at the present day). stellar ages 1might Carry a statistical aud relative significauce. though still uncertain on an individual basis: stars tliat we place lower on the Hayashi tracks may be older. on average. than those that we place higher.

On the other hand, if one were to concede a real age spread in the ONC (plausible as star formation is still taking place at the present day), stellar ages might carry a statistical and relative significance, though still uncertain on an individual basis: stars that we place lower on the Hayashi tracks may be older, on average, than those that we place higher.

This latter hypotliesis is supported by the time evolution of ONC stellar radii recently observed » Rhodeοἱal.(2001).

This latter hypothesis is supported by the time evolution of ONC stellar radii recently observed by \citet{rho01}.

. Although our correlation of Ly with age also seems to argue in favor of a ‘eal age spread. it could also be an artifact of au interrelation between inferred) Ly aud L4. aud he matter remains open.

Although our correlation of $L_X$ with age also seems to argue in favor of a real age spread, it could also be an artifact of an interrelation between inferred $L_X$ and $L_{bol}$, and the matter remains open.

Indeed the treud in Figure 8. corresponds to a constant ratio of X-ray to iometric lumiuosity at different ages aud cau be interpreted equally well in two different. ways: 1) Stars of equal mass have equal bolometric αμα X-ray luminosities. but there are effects that act in exactly the same way on both Ly auc Lo: 2) These stars are saturated and stay saturated hrough their contraction on the Hayashi tracks. thus keeping La/Lig constant.

Indeed the trend in Figure \ref{fig:LXvsA0a} corresponds to a constant ratio of X-ray to bolometric luminosity at different ages and can be interpreted equally well in two different ways: 1) Stars of equal mass have equal bolometric and X-ray luminosities, but there are effects that act in exactly the same way on both $L_X$ and $L_{bol}$ ; 2) These stars are saturated and stay saturated through their contraction on the Hayashi tracks, thus keeping $L_X/L_{bol}$ constant.

Figure 9 presents a plot of ουν) vs. [ουσε) [οι stars iu the 2—3AL. mass bin ancl 'eveals an age clepeucdence aud a sudden drop of Ly at LogAye)6.5 (a similar drop is also seen in Lxy/Ly4).

Figure \ref{fig:LXvsA0b} presents a plot of $Log(L_X)$ vs. $Log(Age)$ for stars in the $2-3~M_{\odot}$ mass bin and reveals an age dependence and a sudden drop of $L_X$ at $Log(Age) \sim 6.5$ (a similar drop is also seen in $L_X/L_{bol}$ ).

This is just the age at which a 2.5AL. star dissipates its convective envelope accordingo the SDF modclels. leudiug further evidence for the convection —activity counection found iu

This is just the age at which a $2.5~M_{\odot}$ star dissipates its convective envelope accordingto the SDF models, lending further evidence for the convection –activity connection found in

where r is the timescale over which the mass is lost.

where $\tau$ is the timescale over which the mass is lost.

The magnetic energv density is equal to the kinetic energy density at the Alfvénn radius which we find to be We consider this further in Section 4. where we apply our model to the recurrent nova U Sco.

The magnetic energy density is equal to the kinetic energy density at the Alfvénn radius which we find to be We consider this further in Section \ref{usco} where we apply our model to the recurrent nova U Sco.

We parametrise the dipole moment with the magnetic field strength at the stellar surface. Da. so that where fs is the stellar radius of the companion which woe find with equation (13)).

We parametrise the dipole moment with the magnetic field strength at the stellar surface, $B_{\rm star}$, so that where $R_2$ is the stellar radius of the companion which we find with equation \ref{rl}) ).

The specific angular momentum of the ejectec mass that is forced to corotate with the secondary is (63|AUG)0 assuming that the secondary is tically locked.

The specific angular momentum of the ejected mass that is forced to corotate with the secondary is $(a_2^2+KR_{\rm A}^2)\,\Omega$ assuming that the secondary is tidally locked.

The distance from the centre of mass to the secondary star is The constant A depends on the distribution of the material within the Alfvénn radius.

The distance from the centre of mass to the secondary star is The constant $K$ depends on the distribution of the material within the Alfvénn radius.

ΕΕ it were distributed uniformly within a spherical shell. then A.=22/3.

If it were distributed uniformly within a spherical shell, then $K=2/3$.

Because the density within the shell varies and the shell itsell will not be perfectly spherical. we take fy=I.

Because the density within the shell varies and the shell itself will not be perfectly spherical, we take $K=1$.

Thus. the angular momentum. loss from the binary to the ejected material is eiven by We estimate the fraction of the ejected mass that gains angular momentum in this wav to be Now with equations (2)) : We substitute this into equation (6)) with AM,=Am, and AAI=0 to find the change in separation We note that this reduces to equation (7)) when there is no magnetic field.

Thus, the angular momentum loss from the binary to the ejected material is given by We estimate the fraction of the ejected mass that gains angular momentum in this way to be Now with equations \ref{da}) ) and \ref{jobs}) ) we find We substitute this into equation \ref{change}) ) with $\Delta M_1=-\Delta m_1$ and $\Delta M_2=0$ to find the change in separation We note that this reduces to equation \ref{over}) ) when there is no magnetic field.

In Fig.

2 we plot APP as a function of the mass ratio for cillerent values of the πμAlfvénn radius.

\ref{dadm2} we plot $\frac{\Delta P/P}{\Delta m_1/M}$ as a function of the mass ratio for different values of the Alfvénn radius.

The larger magnetic fields are even capable of producing an overall decrease to the orbital period. during the nova for small mass ratios.

The larger magnetic fields are even capable of producing an overall decrease to the orbital period during the nova for small mass ratios.

In Fig.

I. we also plot MUSmena [orfor anan. Mfvénn.Alfvénn radius of O.75e@ (long dashed line) for comparison with the other models we have considered.

\ref{dadm} we also plot $\frac{\Delta P/P}{\Delta m_1/M}$ for an Alfvénn radius of $0.75\,a$ (long dashed line) for comparison with the other models we have considered.

For this strong magnetic field. the period change is smaller than with mass transfer to the secondary or [rictional angular momentum losses.

For this strong magnetic field, the period change is smaller than with mass transfer to the secondary or frictional angular momentum losses.

Lf a secondary star has a large magnetic field then it will significantIv. alter the orbital period. change during a nova outburst.

If a secondary star has a large magnetic field then it will significantly alter the orbital period change during a nova outburst.

‘There is a critical mass ratio where the additional terni in equation (29)) causes no change to the separation (or orbital period) Ig<qon then the cjeetecd material that is forced to corotate with the secondary gains angular momentum. angular momentum is lost [rom the orbit and so the separation decreases.

There is a critical mass ratio where the additional term in equation \ref{cor}) ) causes no change to the separation (or orbital period) If $q<q_{\rm crit}$ then the ejected material that is forced to corotate with the secondary gains angular momentum, angular momentum is lost from the orbit and so the separation decreases.

Similarly the orbital period change is smaller than previously predicted.

Llowever. if dder. then the material Forced. to corotate loses angular moment and the orbital period increases.

However, if $q>q_{\rm crit}$, then the material forced to corotate loses angular momentum and the orbital period increases.

We plot this in the solidunm. line in Fie. 3..

We plot this in the solid line in Fig. \ref{qc}.

We note that most observed systems have a mass ratio qd1 (Ritter&Ixolb2003). because there is a critical mass ratio. that depends on AM». above which the mass transfer becomes unstable (e.g.Smith&VanelePutte 2006).

We note that most observed systems have a mass ratio $q\lesssim 1$ \citep{ritter03} because there is a critical mass ratio, that depends on $M_2$, above which the mass transfer becomes unstable \citep[e.g.][]{smith06}.

. Hence in most systems the orbital period will decrease when the effects of à magnetic field are considered.

Hence in most systems the orbital period will decrease when the effects of a magnetic field are considered.

We also plot the dashed line for the mass ratio below which the period change is negative.

We see that for the larger magnetic fields it is possible that the orbital period may actually decrease for all mass ratios.

This effect. could be significant even for larger mass ratios.

This effect could be significant even for larger mass ratios.

Frictional angular momentum losses only dominate for g=0.01 (Sharactal. so this new mechanism potentially has a greater effect on the orbital period change.

Frictional angular momentum losses only dominate for $q\lesssim 0.01$ \citep{shara86} so this new mechanism potentially has a greater effect on the orbital period change.

There are now ten recurrent novae known in our galaxy

There are now ten recurrent novae known in our galaxy \citep[the tenth one was discovered last year;][]{pagnotta09}

populate the ERO class.

The best examples are provided by LBDS53W001 (ROAN= 5.8 2=1.55: Dunlop et al.

The best examples are provided by LBDS 53W091 $R-K=5.8$ ; $z=1.55$; Dunlop et al.

1996: Spinrad et al.

1996; Spinrad et al.

1997). ERO CL 09039143713D (29.Wy=à: 2o1.6: Soifer et al.

1997), ERO CL 0939+4713B $R-K=7$; $z\sim 1.6$; Soifer et al.

1999). and. by most of the EROs ats~1.3 around the QSO 1213-0017 (Liu et al.

1999), and by most of the EROs at $z\sim 1.3$ around the QSO 1213-0017 (Liu et al.

2000).

Also near-LR. spectroscopy confirmed that both dusty and old galaxies contribute to the ERO population (Cimatti et al.

Also near-IR spectroscopy confirmed that both dusty and old galaxies contribute to the ERO population (Cimatti et al.

1999).

Understanding the nature and deriving the abundance of EROs is important to shed light on the controversial Issue of the deficit of high-z elliptical galaxies: according to some results. the number of galaxies with the red colours expected for high-z passively evolved. spheroidals is lower compared to the predictions of passive luminosity evolution. (e.g. Ixaullmann. Charlot White 1996: Zepf 1997: Franceschini et al.

Understanding the nature and deriving the abundance of EROs is important to shed light on the controversial issue of the deficit of $z$ elliptical galaxies: according to some results, the number of galaxies with the red colours expected for $z$ passively evolved spheroidals is lower compared to the predictions of passive luminosity evolution (e.g. Kauffmann, Charlot White 1996; Zepf 1997; Franceschini et al.

1998: Bareer et al.

1998; Barger et al.

1999).

Llowever. other works cid not confirm the existence of such a deficit up to z~2 (e.g. 'l'otani Yoshii 1997: Benitez ct al.

However, other works did not confirm the existence of such a deficit up to $z\sim 2$ (e.g. Totani Yoshii 1997; Benitez et al.

1999: Broadhurst. Bowens 1999: Schade et al.

1999; Broadhurst Bowens 1999; Schade et al.

1999).

As a first step to understand the nature of EROs. we started an imaging survey program aimed at selecting complete samples of such galaxies both in "empty fields and around. high-z radio-loud. AGN.

As a first step to understand the nature of EROs, we started an imaging survey program aimed at selecting complete samples of such galaxies both in “empty” fields and around $z$ radio-loud AGN.

In this paper. we present the results of our survey around racdio-Ioud CN at ο>1.5.

In this paper, we present the results of our survey around radio-loud AGN at $z>1.5$.

The main motivations of such a survey are to investigate whether EROs are more numerous around high-z ACN (as suspected. in previous works: e.g. AleCarthy et al.

The main motivations of such a survey are to investigate whether EROs are more numerous around $z$ AGN (as suspected in previous works; e.g. McCarthy et al.

1992: Dev et al.

1992; Dey et al.

1995). to select. old. passively evolving galaxies abozI1. hy weir optical/near-I1t. colours. to study the environment of high-z racio-loud XN (e.g. Hall. Green Cohen 1998: Hall Creen 1998 and references therein). ane to search for galaxy cluster candidates at 5.

1995), to select old passively evolving galaxies at $z>1$ by their optical/near-IR colours, to study the environment of $z$ radio-loud AGN (e.g. Hall, Green Cohen 1998; Hall Green 1998 and references therein), and to search for galaxy cluster candidates at $z>1.5$ .

We recal that the most distant. spectroscopically confirmed. cluster known to date bas z=1.27 (Stanford et al.

We recall that the most distant spectroscopically confirmed cluster known to date has $z=1.27$ (Stanford et al.

LOOT: Rosati et al.

1997; Rosati et al.

1999).

Other works selected cluster candidates aroun quasars at z1.3 and suggested a significant heterogeneity of the cluster galaxies. including both passively evolving ol ellipticals as well as vounger and custy systems (e.g. Hal Green 1998: Liu et al.

Other works selected cluster candidates around quasars at $z>1.3$ and suggested a significant heterogeneity of the cluster galaxies, including both passively evolving old ellipticals as well as younger and dusty systems (e.g. Hall Green 1998; Liu et al.

2000).

Phroughout this paper we assume Ly=50 + 1 Oy2d) and QO,=0 unless otherwise statect.

Throughout this paper we assume $H_0=50$ $^{-1}$ $^{-1}$, $\Omega_0=1$ and $\Omega_{\Lambda}=0$ unless otherwise stated.

The survey presented here is based on 2 and. A7 -baux imagine.

The survey presented here is based on $R$ - and $K^{\prime}$ -band imaging.

We observed totally 12. fields: 6 around. racio galaxies (Gs) taken from the ALRC sample selected at 408 Az (AleCarthy et al.

We observed totally 14 fields: 6 around radio galaxies (RGs) taken from the MRC sample selected at 408 MHz (McCarthy et al.

1997. IxXapahii ct al.

1997, Kapahi et al.

1998). S arounc raclio-loud quasars (ILOs) taken from the PIS) sample selected at 2.7 Gllz (Wright Ostrupeek 1990).

1998), 8 around radio-loud quasars (RLQs) taken from the PKS sample selected at 2.7 GHz (Wright Ostrupcek 1990).

We note that AIRC101-220 is a broad line radio galaxy (Ixapahi e al.

We note that MRC1017-220 is a broad line radio galaxy (Kapahi et al.

1998).

All the targeted AGN have 1.5«z2.0. with the exception of the quasars PIS 1351-018 (2= 3.71) ane PINS 1556-245 (2= 2.82).

All the targeted AGN have $1.5<z<2.0$, with the exception of the quasars PKS 1351-018 $z=3.71$ ) and PKS 1556-245 $z=2.82$ ).

The radio powers at rest-frame 5 CllzrangeΓΕ for RGs to 430.1077-- for RLOs.

The radio powers at rest-frame 5 GHz range from $2-5 \times 10^{27}$ $^{-1}$ for RGs to $4-30 \times 10^{27}$ $^{-1}$ for RLQs.

Table 1 lists the relevant information about the sample.

“Phe fields were selected according to the redshifts of the AGN. to their Galactic latitude (>207) and to their good observability curing the telescope runs.

The fields were selected according to the redshifts of the AGN, to their Galactic latitude $>20^{\circ}$ ) and to their good observability during the telescope runs.

We have used the predictions of the Bruzual Charlot (1999) evolutionary svnthesis models to define a colour selection. criterion capable to select a complete sample of EROs at high-z.

We have used the predictions of the Bruzual Charlot (1999) evolutionary synthesis models to define a colour selection criterion capable to select a complete sample of EROs at $z$.

In particular. we used instantaneous burst models (also called simple stellar population models. SSP) with cillerent redshifts of formation (2,=2.3.4.5.6) o describe old. spheroidal galaxies.

In particular, we used instantaneous burst models (also called simple stellar population models, SSP) with different redshifts of formation $z_{f}=2,3,4,5,6$ ) to describe old spheroidal galaxies.

In addition. we also considered. models with exponential time-scale for the star ormation rate. SLRxecp(F/r7). with 7=0.1.0.3 Ce or QO=1.0.1 respectively.

In addition, we also considered models with exponential time-scale for the star formation rate, $SFR \propto exp(-t/\tau)$, with $\tau=0.1,0.3$ Gyr for $\Omega= 1,0.1$ respectively.

Such models. correspond. to cases with low residual star formation at 2< (ic. <1 M.vr | for a galaxy with mass Maur=10 M. ). and hey are capable to reproduce the colours and the spectra of local ellipticals. as well as the faint galaxy optical colour and vecshift distributions (Pozzetti. Bruzual Zamorani 1996).

Such models correspond to cases with low residual star formation at $z<2$ (i.e. $<1$ $_{\odot}$ $^{-1}$ for a galaxy with mass $_{gal}=10^{11}$ $_{\odot}$ ), and they are capable to reproduce the colours and the spectra of local ellipticals, as well as the faint galaxy optical colour and redshift distributions (Pozzetti, Bruzual Zamorani 1996).

A Salpeter IME (0.1.<a<125 M.) and. solar metallicity have been adopted.

A Salpeter IMF $0.1<m<125$ $_{\odot}$ ) and solar metallicity have been adopted.

Models with Scalo IME or with supersolar metallicities reach even redder colours at high-:.

Models with Scalo IMF or with supersolar metallicities reach even redder colours at $z$.

Optical to near-Ht AA. colours derived from the acloptecl evolutionary models are shown in Fig.

Optical to near-IR $R-K$ colours derived from the adopted evolutionary models are shown in Fig.

1 for two cosmologies (44,=50 + |. Q—0.1).

1 for two cosmologies $H_0=50$ $^{-1}$ $^{-1}$, $\Omega=0.1,1$ ).

The colour selection threshold adopted in our survey is Rolv>6.

The colour selection threshold adopted in our survey is $R-K>6$.

This allows us to select old galaxies at z>1.2 [ormed at zr;23 in both cosmologics.

This allows us to select old galaxies at $z>1.2$ formed at $z_{f}>3$ in both cosmologies.

Ht is relevant to note that for τ 20.3 Gyr and zy<3. colours 2.AN>6 are never reached.

It is relevant to note that for $\tau>$ 0.3 Gyr and $z_{f}<3$, colours $R-K>6$ are never reached.

In other words. colours 2A6 select old. galaxies whieh had a short episode of star formation in carly cosmological epochs.

In other words, colours $R-K>6$ select old galaxies which had a short episode of star formation in early cosmological epochs.

For instance. assuming no cust extinction. a galaxy at z 21.5 with RoA>6 should have SpIOX. an age C2 Gaver (04= 1.0) or 3 Gyr (04=0.1) and SER<<1 M.vyr. +.

For instance, assuming no dust extinction, a galaxy at $z\approx$ 1.5 with $R-K>6$ should have $z_{f}>3$, an age $>$ 2 Gyr $\Omega_0=1.0$ ) or $>$ 3 Gyr $\Omega_0=0.1$ ) and $SFR<<1$ $_{\odot}$ $^{-1}$ .

eear-infrared imagine was done on 1997 April 1-3 with rw ESO/AIPL 2.2m telescope. with the HUXC2D. camera (Moorwood ct al.

Near-infrared imaging was done on 1997 April 1-3 with the ESO/MPI 2.2m telescope with the IRAC2B camera (Moorwood et al.

1992) equipped with a 256 11οςαΤο array (0.506" /pixel).

1992) equipped with a $\times$ 256 HgCdTe array $^{\prime \prime}$ /pixel).

We used the A" filter in order to reduce 1e thermal noise.

We used the $K^{\prime}$ filter in order to reduce the thermal noise.

The sky conditions were photometric and 10 seeing. was around 1.07. o

The sky conditions were photometric and the seeing was around $^{\prime \prime}$.

sPhe observations. were done aking a number of background-limited. images (typically 10-14) with the telescope moved LO” between cach image.

The observations were done taking a number of background-limited images (typically 10-14) with the telescope moved $10^{\prime \prime}$ between each image.

In each telescope position. theexposures were typically 120 seconds long (c.g. 12 images cach with an exposuretime of 10 seconds).

In each telescope position, theexposures were typically 120 seconds long (e.g. 12 images each with an exposuretime of 10 seconds).

“Phe data reduction was performed. using the ΗΛ reduction package and using the method outlined by Villani di Serego Alighieri (1099).

The data reduction was performed using the IRAF reduction package and using the method outlined by Villani di Serego Alighieri (1999).

Photometric calibration was obtained observing standard. stars from the Carter Meadows. (1995). sample.

Photometric calibration was obtained observing standard stars from the Carter Meadows (1995) sample.

Phe typical night-to-night scatter in the zero-points was around 0.03 magnitudes.

The typical night-to-night scatter in the zero-points was around 0.03 magnitudes.

The conversion from A" to A magnitudes inthe SAAO-Carter

The conversion from $K^{\prime}$ to $K$ magnitudes inthe SAAO-Carter

showing weak or even abscut thermal features) display very high peak frequencics: in IIBL. which appear to be the less powerful objects (with typical Iuuinosities 107 ere cm? 1) the first peals falls in the UV-soft Naav baud aud the high energv component can peak near the TeV reeion.

showing weak or even absent thermal features) display very high peak frequencies: in HBL, which appear to be the less powerful objects (with typical luminosities $10^{42}$ erg $^{-2}$ $^{-1}$ ) the first peak falls in the UV-soft X-ray band and the high energy component can peak near the TeV region.

This behaviour has been iuterpreted (Cbisellini et al.

This behaviour has been interpreted (Ghisellini et al.

1998. recently revisited in (κοπια et al.

1998, recently revisited in Ghisellini et al.

2002) as the result of a balance. between a acceleration i1 cooling suffered. by the electrons producing the radiation at the peaks.

2002) as the result of a balance between a acceleration an cooling suffered by the electrons producing the radiation at the peaks.

Iu the assumption that the acceleration rate is coustant in all the sources. the sequence ean be successfully interpreted as due to au increasing cooling rate (dominated iu most of the sources by IC losses) from the low-power BL Lacs to the powerful FSRQs.

In the assumption that the acceleration rate is constant in all the sources, the sequence can be successfully interpreted as due to an increasing cooling rate (dominated in most of the sources by IC losses) from the low-power BL Lacs to the powerful FSRQs.

SEDs can be satisfactorily reproduced by simple cussion models (e.g. Clisellini et al.

SEDs can be satisfactorily reproduced by simple emission models (e.g. Ghisellini et al.

1998. Tavecchio et al.

1998, Tavecchio et al.

1998). vielding the value of the main plwvsical quautitics of the jet. such as electron density and cnereyv. magnetic fold. size of the region.

1998), yielding the value of the main physical quantities of the jet, such as electron density and energy, magnetic field, size of the region.

In particular miecasuring the N-ray spectra aud adapting a broad band model to their SEDs vields reliable estimates of the total ΙΟ: of relativistic particles involved. which is dominated by those at the lowest energies.

In particular measuring the X-ray spectra and adapting a broad band model to their SEDs yields reliable estimates of the total number of relativistic particles involved, which is dominated by those at the lowest energies.