Disk–Jet Coupling in the 2017/2018 Outburst of the Galactic Black Hole Candidate X-Ray Binary MAXI J1535–571

UVMULTIFIT was also used to search for source variability within each epoch, where we fit for a point source in the uv-plane for different time intervals within each observation. Many of our radio observations were short in duration (generally only 15–30 minutes long), meaning that for many observational epochs the short time variability could not be well explored. Therefore, for most ATCA observations during the outburst we only observed small changes in the source flux density (by ≲3%) over the radio observation.

However, we detected source variability on three epochs, on MJDs 58,013, 58,017, and 58,019, where the source varied significantly when compared to other sources in the field. This variability is discussed in Section 3 and shown in Figures 2 and 3.

Zoom In Zoom Out Reset image size

Zoom In Zoom Out Reset image size

Following our initial radio detection on MJD 58,001, MAXI J1535−571 brightened at both radio and X-ray wavelengths over the next ∼week as the X-ray spectrum slowly softened (Figure 4), although the source remained in the HIMS. After this phase of rapid radio brightening, the radio emission then remained relatively steady (${S}_{5.5\mathrm{GHz}}\approx 170\mbox{--}180$ mJy) from ∼MJD 58,008.5 until MJD 58,012.5 (during which we observed daily; Figure 1). Over this time, although the radio spectrum evolved marginally, the spectral index generally remained flat to inverted, indicating the continued presence of the compact jet.

Zoom In Zoom Out Reset image size

Our final ATCA observation within the HIMS (starting ≈MJD 58,013.5) showed that the radio emission as a whole had faded (${S}_{5.5\mathrm{GHz}}\approx 135$ mJy) and the radio spectrum had steepened while the X-ray luminosity increased (Figure 1). This epoch is a little more complicated than just steady fading and shows strong evidence that the jet spectral break had moved into, and through, the radio band during this epoch. The evolution of the jet spectral break will be discussed in greater detail in T. D. Russell et al. (in preparation), however, we summarize the behavior during this epoch here. This evidence is highlighted by the radio spectrum initially remaining ∼flat (-to inverted) between the 5.5 and 9 GHz ATCA bands, while the 17–19 GHz emission was fainter than the lower-frequency observing bands (and continuing to fade²² ) and exhibited a steep radio spectrum (Figure 2, top). Toward the end of the observation, the 5.5 and 9 GHz radio spectrum also steepened due to rapid fading at 9 GHz. This evolution implies that the jet spectral break resided within the radio band during this observation. Such behavior indicates the beginning of the quenching of the radio jet, which is supported by a sharp drop in the IR and optical brightness between MJDs 58,012 and 58,016 (Baglio et al. 2018).

Our next radio observation (≈MJD 58,017.4) occurred within the SIMS. Dividing that radio observation into 2-minute time intervals showed that MAXI J1535−571 was brighter (${S}_{5.5\mathrm{GHz}}\,\approx 190$ mJy) at the beginning of that observation, with the radio emission fading steadily at all frequencies (to ${S}_{5.5\mathrm{GHz}}\approx 118$ mJy; Figures 1 and 2) during the ≈3.5 hr radio observation. At all times the radio spectrum was steep (α ≈ −0.5). Additionally, the Australian Square Kilometre Array Pathfinder (ASKAP) also detected a bright radio peak (≈580 mJy) in the 1.4 GHz light curve ∼MJD 58,017.17 (Chauhan et al. 2019), which are included in Figure 1. These results indicate that a bright radio flare occurred between MJDs 58,013.5 and 58,017.4.

A second bright, steep-spectrum radio flare was observed to be rising during our radio observation on MJD 58,019 (Figure 3). This second flare then faded over the next few weeks. However, instead of continuing to fade, MAXI J1535−571 temporarily returned back to a hard state between MJDs 58,044 and 58,054. During this period the radio jet re-brightened and the radio spectrum was observed to be flat (Figure 1).

After this brief return to the hard state, MAXI J1535−571 entered the soft state and faded slowly in the X-ray band over the next ∼160 days. During this steady soft-state decay, we continued to detect radio emission (with generally decreasing flux density and a steep radio spectrum) until MJD 58,161, more than 100 days after the hard-to-soft state transition. Throughout this evolution, we observed some minor radio flaring and changes to the radio spectrum (although α was always <0), indicating that the emission originated from a transient jet.

During the decay period at the end of the major outburst, we did not detect radio emission from the source. Of particular interest, during and after the soft-to-hard state transition (MJD 58,237) at the end of the outburst, we detected no radio emission from the source down to 3σ upper limits of ≈120 μJy beam⁻¹.

In their hard states, BH XRBs exhibit an empirical correlation between their radio (L_R) and X-ray (L_X) luminosities, which is observed over several orders of magnitude in luminosity (e.g., Corbel et al. 2000, 2003, 2013a; Gallo et al. 2003, 2012, 2018). This nonlinear relationship is generally described by two power-law tracks: an upper "radio-loud" track with a slope of ${L}_{{\rm{R}}}\,\propto {L}_{{\rm{X}}}^{\sim 0.6}$, and a lower "radio-quiet" track with a slope of ${L}_{{\rm{R}}}\propto {L}_{{\rm{X}}}^{\sim 1}$ (e.g., Coriat et al. 2011; Gallo et al. 2012), which show different radio spectral indices (Espinasse & Fender 2018). We note that recent results have questioned the statistical significance of there being two separate tracks (Gallo et al. 2014, 2018).

We investigated the radio and X-ray relationship of MAXI J1535−571 by placing a contemporaneous hard state 5 GHz radio and 1–10 keV X-ray luminosities on the radio/X-ray plane (Figure 8). Unfortunately, due to the low cadence of our radio observations during the rising hard state, and the non-detection of radio emission following the transition back to the hard state at the end of the major outburst, we are only able to place two detections on the radio/X-ray correlation (with the upper limits showing the non-detection of the jet at the end of the outburst). From these two points, we determine a slope of ${L}_{{\rm{R}}}\propto {L}_{{\rm{X}}}^{1.37\pm 0.05}$. However, while it appears that MAXI J1535−571 traced out the radio-quiet track, these two points only span 1 order of magnitude in L_R and a factor of ∼6 in L_X, where previous studies (e.g., Corbel et al. 2013a) have shown significant deviation from a source's standard behavior for luminosity ranges <2 orders of magnitude.

Zoom In Zoom Out Reset image size

Following a period of relatively steady radio emission from the compact jet (∼MJD 58,008 until MJD 58,012; Figure 1), we observed the initial stages of jet quenching. ATCA radio observations on MJD 58,013.5 showed that the radio emission had faded by a ≳25% (and was continuing to fade) and the jet spectral break had evolved into the radio band (Section 3.3.1, but will be discussed in detail by Russell et al. 2019, in preparation), driving the observed jet quenching (e.g., Fender et al. 1999; Coriat et al. 2011; van der Horst et al. 2013; Russell et al. 2013b, 2014).

The compact jet also re-formed during MAXI J1535−571's brief return to the hard state, quenching once again as the source transitioned back to the soft state (through the intermediate states). This second quenching was also characterized by fading radio emission and a steepening radio spectrum.

The radio emission from MAXI J1535−571 faded below our detection limits on ∼MJD 58,166. Comparing our deepest soft-state radio upper-limit with the steady, flat-spectrum radio emission measured during the HIMS provides a lower-limit of >3.5 orders of magnitude on the jet quenching factor. This is the deepest constraint on the soft state jet quenching to date, suggesting that the compact jet was not present during this soft state (see also Fender et al. 1999; Coriat et al. 2011; Russell et al. 2011). However, our observations do not rule out jets with low-radiative efficiency in the soft states (Sikora et al. 2005; Drappeau et al. 2017), although see Koljonen et al. (2018) for evidence against a dark jet during the soft state in Cygnus X-3.

At the end of its major outburst, MAXI J1535−571 transitioned from the soft to the hard state. This reverse transition was interesting for two reasons: (1) it occurred at a much lower X-ray luminosity than is generally expected, and (2) the radio jet was not detected.

In a typical BH XRB outburst, the hard-to-HIMS transition during the outburst rise occurs at an X-ray luminosity of ≥3% ${L}_{\mathrm{Edd}}$, while the lower-luminosity transition back to the hard state at the end of an outburst occurs between 0.3% and 10% of L_Edd (Dunn et al. 2010; Kalemci et al. 2013; Vahdat Motlagh et al. 2019), with an average value of ∼2% (Maccarone 2003). However, MAXI J1535−571 only transitioned back to the hard state at L_X ∼ 0.003% ${L}_{\mathrm{Edd}}$ (see also Chauhan et al. 2019). Such low-luminosity transitions are atypical for BH XRBs, and the only other source to show similar behavior is 4U 1630−47, which transitioned from the soft-to-hard state at L_X ≈0.008 % ${L}_{\mathrm{Edd}}$ (Tomsick et al. 2014).

Low transition luminosities have been attributed to either the decay being disrupted by a new mass inflow reigniting and extending the soft state to a lower than usual X-ray luminosity (Vahdat Motlagh et al. 2019), or due to a low disk magnetic field and viscosity (e.g., Petrucci et al. 2008; Begelman & Armitage 2014).

For the case of a new mass inflow extending the soft state and pushing the reverse transition to much lower than typical luminosities, the jet may have been undetected due to the low transition luminosity. Assuming typical hard state ${L}_{{\rm{R}}}$/${L}_{{\rm{X}}}$ scalings (Figure 8), the X-ray luminosity over the reverse transition at the end of MAXI J1535−571's outburst implies an expected radio flux density of ∼50–400 μJy. Therefore, while our 3σ radio upper limits are relatively radio faint in comparison with the majority of other hard state BHs, the radio non-detection is not remarkable.

In the case where low disk magnetic fields were responsible for the low-luminosity state transition, we may also expect weaker radio jets (e.g., Shibata & Uchida 1986; Kylafis et al. 2012; Begelman & Armitage 2014; Kylafis & Belloni 2015), resulting in the radio non-detection. Our data do not allow us to conclusively determine the cause of the low-luminosity transition or non-detection of the radio jets.

The proper motion of 47.2 ± 1.5 mas day⁻¹ indicates an apparent jet velocity of ≈1.1 c for a source distance of 4.1 kpc. Apparent superluminal motions of jet ejecta have been observed in a handful of BH XRBs to date (e.g., Mirabel & Rodríguez 1994; Hjellming & Rupen 1995; Tingay et al. 1995; Rupen et al. 1998; Fender et al. 1999; Mirabel & Rodríguez 1999; Rodríguez & Mirabel 1999; Corbel et al. 2002; Fender et al. 2002, see Fender 2006, for review), and indicates that S2 was the approaching component.

Our tracking of S2 can be used to constrain the inclination of the jet (at the time of the ejection event; Miller-Jones et al. 2019) and the speed of the ejection. We did not detect a counter-jet component, so we cannot uniquely solve for the jet speed, $\beta =\tfrac{v}{c}$, or inclination, θ. Instead, we can only solve for $\beta \cos \theta$, given that (e.g., Rees 1966; Mirabel & Rodríguez 1994):

$$\begin{eqnarray}&&\mathop{\mu \mathrm{app}}\limits_{\mathrm{rec}}=\displaystyle \frac{\beta \sin \theta }{1\mp \beta \cos \theta }\displaystyle \frac{c}{D},\end{eqnarray} \tag{ 3 }$$

where μ_app and μ_rec are the approaching and receding proper motions, and D is the distance to the source. As shown in Figure 9, from its apparent superluminal motion, S2 is almost certainly the approaching component, with only a small set of solutions existing for it to be the receding component (only at the lowest distance limit; Chauhan et al. 2019), $\beta \cos \theta \geqslant 0.49$, such that β ≥ 0.69 and θ ≤ 45°.

Zoom In Zoom Out Reset image size

X-ray spectral fits of the iron line during the HIMS (Miller et al. 2018; Xu et al. 2018) favored a disk inclination of ∼55°–68°, which is discrepant from our determined values for S2. However, there was some evidence for disk warping (Miller et al. 2018), suggesting that the inner disk orientation was changing over time. Such changes in the inner disk and jet orientation can be rapid (e.g., Liska et al. 2018; Miller-Jones et al. 2019), and would account for such a difference between the disk inclination and the jet when measured at different times. Additionally, the jet and outer disk may be misaligned.

S2 remained unresolved for all radio detections. Therefore, while we do not observe the expansion of S2, our observations can constrain the opening angle and transverse expansion of the jet from its width at a given distance from the core (e.g., Fender et al. 1999). Combining the 29 ATCA resolution with the ∼17'' separation we measured when S2 was detected at its maximum separation, the jet opening angle is constrained to ≤10°, similar to typical constraints on BH XRB jet opening angles (see e.g., Miller-Jones et al. 2006). In terms of the transverse expansion, combining the ATCA beam size of our final detection with the ejection date, we limit the expansion velocity to ≤0.18 c, consistent with the expansion estimates found for V404 Cygni (Tetarenko et al. 2017b).

Using the proper motion of S2, we also estimate the size scale of the radio emission based on the delay between the ejection time and the time of the first radio flare. While we do not detect the peak of the first radio flare, we constrain it to have occurred between MJDs 58,013.6 and 58,017.4 (Section 3.3). Therefore, the ∼GHz radio emitting region lies at a distance of <430 mas from the radio core. For a source distance of 4.1 kpc, this corresponds to a size of <1760 au.

From the rise time and brightness reached by the second radio flare,²³ we also place constraints on the minimum energy and magnetic field required to produce such a flare. Following Fender (2006) and assuming equipartition between electrons and magnetic field, and one proton per electron, we estimate the minimum energy E_min ∼ 10⁴²–10⁴³ erg, corresponding to a minimum mean power P_min ∼ 10³⁷ erg s⁻¹ with an equipartition magnetic field, B_eq ∼10–500 mGauss. These values are comparable to estimates from observed ejection events from a number of other sources (e.g., Fender et al. 1999; Brocksopp et al. 2007; Curran et al. 2014). Our energy estimates differ from the ejecta observed from V404 Cygni (Tetarenko et al. 2017b; Miller-Jones et al. 2019), however, in that outburst, multiple, smaller ejecta were observed, explaining the lower energies and higher magnetic fields.

While there is a connection between the morphology of the jets and structure of the accretion flow (e.g., Fender et al. 2004), the causal sequence of events leading to the changes in the jets is not well understood, and the coupling is likely to be complex. The onset of the transient jet is generally linked to the transition from the HIMS to SIMS. The X-ray properties over this transition show a marked change (e.g., Belloni et al. 2005). For example, there is a sharp decrease in the fractional rms variability of the X-ray emission and the sudden appearance of a type-B QPO. It is these changes in the X-ray emission that have often been linked to the moment of an ejection event, in particular the presence of the type-B QPO (e.g., Fender 2006; Klein-Wolt & van der Klis 2008; Fender et al. 2009; Miller-Jones et al. 2012a). However, sparse observational sampling and the delay between an ejection and the observed radio flaring (due to optical depth effects and the time required for shocks to occur) generally prevents such a connection being clearly identified (see Fender et al. 2009, for a review).

During our radio coverage of the outburst of MAXI J1535−571, we observed the onset of the compact jet quenching in the radio band (which started ∼MJD 58,013.6). Our next radio observation (starting on MJD 58,017.4) detected the end of a radio flare. Therefore, this initial radio flare occurred between MJDs 58,013.6 and 58,017.4 and was likely associated with the ejection of S2, meaning that we would typically expect the ejection to have occurred at an earlier time (due to the time delay between an ejection event and the subsequent radio flaring; Fender et al. 2009).

Comparing our ejection window (between MJDs 58,008.3 and 58,013.4) to the X-ray properties produces interesting results. MAXI J1535−571 transitioned from the HIMS to SIMS sometime between MJDs 58,014.18 and 58,015.37 (Huang et al. 2018; Tao et al. 2018), close in time to the estimated ejection of S2. Huang et al. (2018) report a detection of a possible type-B QPO during HXMT observations taken on MJD 58015.97 (although this QPO was observed at a relatively high QPO frequency of ∼10 Hz). While this QPO could be consistent with being close to the peak of the first radio flare (but certainly not before the beginning of the flare), it is after our t_ej window. In observations after this potential type-B QPO detection, only type-C QPOs were observed until MJD 58,017.5. HXMT and AstroSat observations taken during our ejection event window show high fractional X-ray variability (≳10%) with only the presence of type-C QPOs, with a changing QPO frequency (the QPO frequency first decreased, reaching a minimum ∼MJD 58,010, before increasing during the ejection window, although some scatter was observed; Huang et al. 2018; Bhargava et al. 2019; Sreehari et al. 2019).

Conveniently, the Neutron Star Interior Composition Explorer (NICER) X-ray telescope on board the International Space Station densely monitored MAXI J1535−571 around the time of the HIMS-to-SIMS transition (observing the source multiple times nearly every day). As reported by Stevens et al. (2018), the NICER X-ray observations also detected the appearance of a possible type-B QPO during observations starting on MJD 58,016.8, that remained until MJD 58,025 (but not during observations earlier and in particular, at similar times to when HXMT reported a type-B QPO). While this type-B QPO could be coincident with the second radio flare, it is not consistent with t_ej. Additionally, NICER's high observing cadence showed the X-ray fractional variability was relatively high (dropping from ≈15% to ≈11%) over our t_ej window, making it seem unlikely that a type-B QPO was present previously (typically, type-B QPOs occur at times of low fractional X-ray variability, ≈3–5% rms; Belloni 2010).

The X-ray observations allow us to further investigate the X-ray properties at the time of t_ej. There was an initially steady drop in the X-ray rms variability (from ≈15% to ≈11% rms) between MJDs 58,011 and 58,014, followed by a more rapid decrease (≈11% to ≈7% rms) between MJDs 58,013 and 58,015 (see Figure 1 in Stevens et al. 2018). Over this time, the X-ray observations also showed a steady drop in X-ray hardness, as well as an interesting and relatively sudden increase in the soft X-rays, where the count rate increased rapidly by a factor of ∼2 between MJDs 58,010 and 58,014 (see Figure 1 in Stevens et al. 2018). While this increase could be due to a change in the accretion properties, similar short–high-energy brightenings are commonly associated to ejection events in active galactic nucleus (e.g., Marscher et al. 2008; King et al. 2016; Lisakov et al. 2017), where the X-ray increase can arise from inverse Compton scattering of the synchrotron radiation from the knot after it was ejected. If this X-ray increase was indeed related to the ejection of S2, it favors an ejection time of ∼MJD 58,010. Additionally, around the same time the QPO frequency rapidly decreased, reaching a local minimum ∼MJD 58,010 (after which it increased again; Huang et al. 2018; Sreehari et al. 2019). Speculatively, if the QPO frequency is related to the radius of the X-ray emitting material (e.g., Ingram et al. 2009), we may expect the QPO frequency to decrease due the extraction of accreting material in a jet ejection (Rapisarda et al. 2014). In addition, there was also a change in the QPO behavior at around MJD 58,013, when the QPO frequency varied rapidly (Bhargava et al. 2019).

Using high-resolution radio observations of the 2009 outburst of H1743−322, Miller-Jones et al. (2012b) were able track the motion of a bipolar ejection over two epochs, allowing tight constraints on the time of ejection (to within 1 day). While their estimated ejection date unfortunately coincided with a 3 day gap in X-ray monitoring, they determined that it occurred immediately prior to the HIMS-to-SIMS transition. Over this time, H1743−322 displayed a short increase (also by a factor of ∼2) in the X-ray count rate, a rapid reduction in the X-ray rms variability, an evolution in the type-C QPOs, and the onset of compact jet quenching. Additionally, Miller-Jones et al. (2012b) only report the first appearance of a type-B QPO ∼4 days after the estimated ejection event.

The similarities shown by H1743−322 and MAXI J1535−571 are striking. While in both cases we are unable to conclusively rule out the presence of a type-B QPO at the time of the jet ejection, the results suggest that, for these two outbursts, the X-ray signature of the ejection was not the appearance of the type-B QPO. It could be that the jet ejection and type-B QPO are both a result of some other effect. Similar to findings reported by Miller-Jones et al. (2012b), our monitoring implies that the ejection event was instead related to the rapid drop in X-ray rms variability immediately prior to the HIMS-to-SIMS transition, the sudden increase in soft X-ray count rate, or the change in the type-C QPO frequency.

Fender et al. (2009) and Miller-Jones et al. (2012b) also compared outburst data from a number of other BH LMXBs, using either very long baseline interferometry data to trace ejecta back in time, or by connecting the timing of the radio flares to the X-ray behavior. Their analysis found no clear evidence of an association between the jet ejection and the appearance of type-B QPOs. In most cases, it appeared that the time of ejection was contemporaneous with a change in the type-C QPO and a decrease in the X-ray rms variability. However, as discussed by Fender et al. (2009) and Miller-Jones et al. (2012b), this did not hold true for all systems, or even outbursts from the same system. Radio and X-ray observations of the 2002 outburst of GX 339−4 showed that while the type-C QPO was changing at the time of the radio flare (Gallo et al. 2004; Belloni et al. 2005), the drop in X-ray rms variability was observed a few days after the radio flare (where the rms drop occurred at around the time of the detection of type-B QPO; Fender et al. 2009). Additionally, the 2003 outburst of H1743−322 did not appear to show an evolution of the type-C QPO during the estimated time of ejection (Miller-Jones et al. 2012b).

Therefore, while our MAXI J1535−571 results agree well with the 2009 outburst of H1743−322, as well as a number of other systems (as presented by Fender et al. 2009 and Miller-Jones et al. 2012b) comparisons with the 2002 outburst of GX 339−4 and 2003 outburst of H1743−322 muddy the picture. These two results imply that the events driving an ejection event may vary between systems and even outbursts of the same system. Alternatively, as discussed by Fender & Gallo (2014) there may be no clear X-ray signature to the moment of ejection, and the observed knots could be a result of internal shocks arising within the jets from rapid (but not instantaneous) changes in the injection, or speed of the jet-channeled accretion material (e.g., Jamil et al. 2010; Malzac 2013).